{
    "0404/astro-ph0404257.txt": {
        "abstract": "\\vskip-4.6in \\begin{flushright} UMN-TH-2302/04 \\\\ TPI-MINN-04/13 \\\\ astro-ph/0404257 \\\\ April 2004 \\end{flushright} \\vskip+3.6in Evidence from a large sample of quasar absorption-line spectra in damped Lyman-$\\alpha$ systems has suggested a possible time variation of the fine structure constant $\\alpha$. The most statistically significant portion of this sample involves the comparison of Mg and Fe wavelength shifts using the many-multiplet (MM) method. However, the sensitivity of this method to the abundance of heavy isotopes, especially Mg, is enough to imitate an apparent variation in $\\alpha$ in the redshift range $0.5 < z < 1.8$. We implement recent yields of  intermediate mass (IM)  stars into a chemical evolution model and show that the ensuing isotope distribution of Mg can account for the observed variation in $\\alpha$ provided the early IMF was particularly rich in intermediate mass stars (or the  heavy Mg isotope yields from AGB stars are even higher than in present-day models).  As such, these observations of quasar absorption spectra can be used to probe the nucleosynthetic history of low-metallicity damped Lyman-$\\alpha$ systems in the redshift range $0.5 < z < 1.8$. This analysis,  in conjunction with other abundance measurements of low-metallicity systems, reinforces the mounting evidence that star formation at low metallicities may have been strongly influenced by a population of         IM         stars. Such        IM         stars have a significant influence on other abundances, particularly nitrogen. We constrain our models with independent measurements of N, Si, and Fe in damped Lyman-alpha systems as well as C/O in low-metallicity stars. In this way, we obtain consistent model parameters for this chemical-evolution interpretation of the MM method results. ",
        "introduction": "%===================================================================== The origin and dynamics of the fundamental constants of Nature is one of the deepest questions in physics. One of the most widely held tenets in physics is that the laws of nature are universal, constant, and favor symmetries. Nevertheless,  in various  unified theories (including string theory), gauge and Yukawa couplings often appear as dynamical variables which are only ``fixed\" when a related scalar field (such as a dilaton in string theory) picks up a vacuum expectation value. While one may naturally expect that couplings become constant at or near the unification scale, only experimental evidence can constrain the degree to which these constants vary at late times.  In this context there has been considerable excitement in recent years over the prospect that a time variation in the dimensionless fine structure constant, $\\alpha$, may have been observed (Webb et al.~1999, Murphy et al.~2001a,b, Murphy et al.~2003a).  This evidence is based upon an application of the ``many multiplet\" (MM) method to quasar absorption lines in damped Lyman-$\\alpha$ systems (DLAs). Attempts at constraining or measuring time variations of $\\alpha$ in quasar absorption-line spectra have a long history going back to work by Bahcall \\& Salpeter (1965) using O$III$ and Ne$III$ emission lines. This method was reexamined recently  by Bahcall, Steinhardt, \\& Schlegel (2004). Other recent attempts include measurements of absorption line spectra in alkali-like atoms (Potekhin  \\& Varshalovich 1994, Varshalovich \\& Potekhin 1994, Murphy et al.~2001c, Fiorenzano, Vladilo, \\& Bonifacio 2003). While many observations have led to interesting limits on the temporal variation of $\\alpha$ (for a recent review, see Uzan 2003), only the MM method has led to a quantitatively positive result. Murphy et al.~(2003a) deduce that ${\\delta \\alpha  /\\alpha}  = (-0.54 \\pm 0.12) \\times 10^{-5}$ over a redshift range of $0.5 < z < 3$, where $\\delta \\alpha$ is defined as the deviation from present value. The implications of this deduced variation in $\\alpha$ at around a 5$\\sigma$ significance are phenomenal, and several cosmological models to explain its origin have been proposed (see e.g., Beckenstein 1982, Sandvik, Barrow, \\& Magueijo 2002, Olive \\& Pospelov 2002, Wetterich 2003, Anchordoqui \\& Goldberg 2003, Copeland, Nunes, \\& Pospelov 2004, Lee, Lee, \\& Ng 2003; Byrne \\& Kolda 2004). The caveat of implementing such a precise method, however, is its possible sensitivity to unknown systematic errors.  Some of the excitement concerning the evolution of the fine structure constant has been focused on finding alternative explanations of the observed line structures or other systematic errors.  Chand et al.~(2004) and Srianand et al.~(2004) probed the sensitivity of the MM method with respect to synthetic signal alterations and found that the MM method may break down in well blended, multi-cloud systems. They also applied the MM method independently and found   ${\\delta \\alpha  / \\alpha}  = (-0.06 \\pm 0.06) \\times 10^{-5}$. Another group (Quast, Reimers, \\& Levshakov 2004) has also recently applied the MM method utilizing exceptionally high-resolution QSO absorption-line spectra. Their results are also consistent with a null hypothesis regarding the fine structure evolution. Other potential systematic errors in the MM method have been elicited by others. They  involve the cloud velocity structure and line blending (Bahcall et al.~2004) or  cloud inhomogeneity and spectrographic inconsistencies (Levshakov 2003).    A number of sources of possible systematic error in this method have been well documented (Murphy et al.~2001b and 2003b; see also Bahcall et al.~2004). Here, we will focus on one of these possible systematic errors for which there is recent evidence of a new interpretation, namely the isotopic abundances of Mg assumed in the analysis. In this paper,  we expand on our earlier studies (Ashenfelter, Mathews \\& Olive 2004) of possible systematic effects from the chemical evolution of magnesium isotopes within DLA quasar absorption-line systems. All of the results quoted above are based on the {\\em assumption} that the isotopic abundance ratios of Mg are Solar in these systems. Based upon galactic chemical evolution studies previously available (Timmes, Woosley \\& Weaver 1995), one could argue that the ratio of $^{25,26}$Mg/$^{24}$Mg is expected to decrease at low metallicity. In this case, if it is assumed that only $^{24}$Mg is present in the absorbers, the Murphy et al.~(2003a) result becomes significantly stronger ${\\delta \\alpha  / \\alpha}  = (-0.86 \\pm 0.10) \\times 10^{-5}$ (assuming also that only $^{28}$Si is present) and the Chand et al.~(2004) limit becomes a detection  ${\\delta \\alpha  / \\alpha}  = (-0.36 \\pm 0.06) \\times 10^{-5}$. Hence, it is possible that the detections of time-varying $\\alpha$ are even more significant than the quoted confidence limits. However, we show that it is also plausible that the $^{25,26}$Mg/$^{24}$Mg ratio was in fact sufficiently {\\em higher} at low metallicity to account for the apparent variation in $\\alpha$ as seen in the so-called ``low redshift\"  ($0.5 < z < 1.8$) data. Thus,  the MM method of analysis may provide important new insights into the chemical evolution of damped Lyman-$\\alpha$ quasar absorption-line systems rather than conclusive evidence for a time-varying fine-structure constant.  We begin the present discussion by briefly reviewing the current observational limits on the variations of the fine structure constant.  In section 3, we discuss the theory and observations of the Mg isotopes. The sensitivity of $(\\delta \\alpha / \\alpha)$ to the Mg isotopic abundances is explained in section 4.  In section 5, we describe a simple chemical evolution model, which we then utilize to address the question of the history of the Mg isotopes and other elements as a function of metallicity. Results of this study  and the sensitivity to the model are discussed in section 6. Our summary and conclusions are given in section 7.  %\\begin{figure} %\\mbox{\\epsfig{file=agbimf1.ps,height=7.5cm,angle=270}} %\\caption{\\label{fig:epsart} %\\end{figure} ",
        "conclusions": "We have made a study of possible relations among stellar nucleosynthesis, the galactic chemical evolution of damped Lyman-$\\alpha$ systems, and the apparent detections of a time-varying fine structure constant. In particular, we have explored the important effects of high temperature thermonuclear burning in low-metallicity AGB stars.  We have shown that ejecta from these stars could have had a dominant effect on the early galactic chemical evolution of the crucial Mg isotopes in DLA systems.  We have explored a variety of models in which the early initial mass function favors the formation of        IM         stars.  Such an enhanced contribution from early IM stars allows for sufficient modification of the Mg isotope ratios to explain the many-multiplet results without a time-varying fine structure constant.  Such a modified IMF may to some extent be a simple parametrization of uncertainties in theoretical estimates of ejected yields from AGB stars, but it is motivated by both theoretical and observational constraints.  To compare with the MM method results, we have utilized an approximate treatment that qualitatively relates computed Mg isotopic abundances to the deduced fine structure constant for DLA systems in the redshift range of $0.5 < z < 1.8$.  Although this is only an approximation to the MM approach, we have shown that it reproduces the  basic conclusions of detailed analysis (Murphy et al.~2003a). There is a real need to redo the MM method analysis in the context of evolving isotopic abundances as derived here.  Incorporating isotopic variation would help to better quantify the need (or lack thereof) for IM stars and AGB nucleosynthesis in DLA systems. We hope that the present work will stimulate efforts along this line.  In the context or our schematic analysis, we have explored a variety of chemical evolution models with an eye toward unraveling the time-varying alpha mystery while still satisfying the available constraints from observed elemental abundances in DLA systems. We have concentrated on the chemical evolution of N abundances, which are also produced in AGB stars. We also considered C as well as  O and Si largely from Type II supernovae, and Fe from SNIa and SNII. We find that the observed high nitrogen abundances in DLA systems indeed  confirm the need for enhanced ejecta from low-metallicity AGB stars. Even so, our previous model (Ashenfelter et al.~2004), which attempted to explain the MM results of (Murphy et al 2003a) tends to overproduce nitrogen and is therefore constrained by the observations.  In this paper, however, we report on a parameter search which considers both data sets.  We find  a new optimum model (Model  2 in the present work) which simultaneously fits the observed N/H, C/O, and N/Si trends vs [Fe/H] while still eliminating the need for the time-varying fine structure constant as deduced from the Murphy et al.~(2003a) data for systems in the redshift range $0.5 < z < 1.8$.  At the same time, we have also constructed a new model (Model 3) that can account for the results of the  independent MM method analysis of Chand et al.~(2004), which indicate smaller apparent variations in the fine structure constant.  Even though these authors claim results that are consistent with no variation in $\\alpha$, this conclusion is based on the assumption that the Mg isotopic ratio is equal to the Solar one.  As we have shown, in order to obtain a Solar isotope ratio at low metallicity, we must again rely on the role of IM stars and AGB nucleosynthesis.   One conclusion of the present study is that important tests can be made of the hypothesis that AGB nucleosynthesis can account for the apparent variation in the fine structure constant.  The best measurement (though probably impossible) would be to directly detect Mg isotopic abundances from spectral lines.  In our picture, the apparent variations in $\\alpha$ should correlate directly with the fraction of heavy Mg isotopes. Conversely, if heavy magnesium abundances are significantly depleted relative to Solar, then the MM method results are actually understating the variation in the fine structure constant. Furthermore, if sufficiently precise data could be obtained to distinguish the $^{25}$Mg and $^{26}$Mg abundances, then large enhancements of $^{26}$Mg observed in some systems could be indicative of Mg production that is specifically attributable to the Mg-Al cycle.  If so, large enhancements of $^{26}$Mg may also be anti-correlated with Al abundances.   We further suggest that nitrogen (and/or carbon) abundances provide an easier test of the present hypothesis with regards to an enhanced IMF for low-metallicity        IM         stars. Nitrogen  should be measured and correlated with $\\delta \\alpha/\\alpha$ in the same DLA systems to which the MM method is applied. A correlation of [N/H] with the largest variations in $\\alpha$ would argue in favor of the present hypothesis. As another test of an enhanced IMF, the highest [Fe/H] or [Si/H] (Si and Fe are very correlated) in DLA systems should exhibit a significant over abundance of [N/$\\alpha$] if our IMF is correct. One caveat in using Si abundances is its dependence on the explosion conditions for the Type II supernovae. Future work regarding these tests should examine the consequences of adopting these different explosion criteria and compare with different yield models.  Determining whether or not a metallicity condition exists for Type Ia supernovae is decisively related to an IMF enhanced with intermediate-mass stars. It may very well be that the reduced efficiency of Type Ia supernovae is offset by the enhanced numbers of intermediate mass stars, which may satisfy the previous work of Matteucci \\& Recchi ~(2001) and the metallicity conditions of Nomoto et. al ~(2003).  Comparison of abundances of N, C, and Mg with [O/H] may be able to test the viability of IM mass enhanced IMF models independent of Sn Ia rates.  One should interpret the large scatter in the inferred variation in $\\alpha$, as well as the observed variation in element abundances like Mg, in the context of the stochastic nature of the star formation process at low metallicity.  Our models simply represent a global average, and the contribution of IM stars may very well vary in individual DLAs.     Clearly, more work needs to be done in constraining the effect of chemical evolution on the interpretations of DLA observations. Nevertheless, we have established that at least some fraction of the deviations in the fine-structure constant deduced from the MM method could be due to chemical evolution effects. Among work that needs to be done, more stellar models of low-metallicity AGB stars  over a broader mass and metallicity range are needed to quantify the model predictions, particularly in the extrapolated mass range of $7 <$ M $< 12$ M$_\\odot$ between the AGB and Type II regimes. In particular, it would be very useful to have have a full complement of stellar yields (including C,  N and Mg) in the IM range derived from a self-consistent set of stellar models. Out of necessity, we supplemented the Mg yields of Karakas \\& Lattanzio (2003) with the Padova CNO yield models; however, a self-consistent model would more accurately quantify the correlation of Mg from AGB sources to N.   At the same time, more observations of abundances in DLA systems are needed. Most importantly, the DLAs that are used in the MM-method should have their associated abundances quantified in order to see if there are correlations between the apparent variations in $\\alpha$ and various abundances. Additionally, more atomic physics work needs to be done to more accurately quantify the possible isotopic shifts in the absorption lines (particularly for the high-$z$ data).  The MM method has presented a very important question as to whether the fine structure constant varies with time. It is hoped that the present study will stimulate further efforts along all of the above  lines with a goal of clarifying  this important physical question. If the fine structure constant does vary in time or space,  it provides an important window into the physics beyond the standard model. As such, the chemical evolution effects described herein should be carefully quantified to reduce the systematic uncertainties in the deduced result.  Even if our chemical-evolution interpretation of the MM results prove to be verified, the MM method will have provided valuable insight into the mysteries of early cosmic star formation and galaxy  evolution.   %{\\bf Acknowledgments}  We thank V. Flambaum for helpful conversations. The work of K.A.O. was partially supported by DOE grant DE--FG02--94ER--40823. Work at the University of Notre Dame was supported by the U.S.~Department of Energy under Nuclear Theory Grant DE-FG02-95-ER40934. One of the authors (T.A.) also  wishes to acknowledge partial support from NSF grant PHY02-16783 through the Joint Institute for Nuclear Astrophysics physics (JINA)."
    },
    "0404/hep-ph0404061_arXiv.txt": {
        "abstract": "We provide a summary  of the current knowledge, theoretical and experimental, of solar neutrino fluxes and of the masses and mixing angles that characterize solar neutrino oscillations. We also summarize the principal reasons for doing new solar neutrino experiments and what we think may be learned from the future measurements. ",
        "introduction": "\\label{sec:introduction}  We record in this paper a snapshot (taken on March 1, 2004) of where we stand with solar neutrino theoretical research.  We do not attempt to review the many papers written on this subject. For details of the extensive literature, the reader is referred to earlier, more comprehensive studies~\\cite{pontecorvo,msw,book,bp00,cabibbo02,fiorentini02,conchayossi,smirnov,roadmap,bargerreview, bilenky03,kayser03,murayama,analysispostkamland,valle03,chitre01,sackmann03,turckchieze,haxton04}.  The related subject of solar neutrino experiments will be reviewed in this volume by A. McDonald~\\cite{mcdonald}. We therefore do not discuss the experimental aspects of solar neutrino research in this article, although we do emphasize the relation between theoretical ideas and predictions and solar neutrino measurements.  We begin in Section~\\ref{sec:models} by summarizing our current theoretical knowledge of the solar neutrino fluxes.  We then summarize in Section~\\ref{sec:parameters} the numerical results regarding solar neutrino parameters and neutrino fluxes that have been inferred from solar neutrino and reactor experiments. Neutrinos are the first cosmological dark matter to be discovered. We describe in Section~\\ref{sec:darkmatter} what solar and atmospheric neutrino experiments have taught us about the cosmological mass density in neutrinos. Finally, in Section~\\ref{sec:newexperiments} we discuss the reasons for doing future solar neutrino experiments and the scientific results that may be obtained from the proposed new experiments. ",
        "conclusions": ""
    },
    "0404/astro-ph0404248_arXiv.txt": {
        "abstract": "{ We present \\textit{K}-band commissioning observations of the Mira star prototype $o$ Cet obtained at the ESO Very Large Telescope Interferometer (VLTI) with the VINCI instrument and two siderostats. The observations were carried out between 2001 October and December, in 2002 January and December, and in 2003 January. Rosseland angular radii are derived from the measured visibilities by fitting theoretical visibility functions obtained from center-to-limb intensity variations (CLVs) of Mira star models \\citep{BSW,HSW,TLSW}. Using the derived Rosseland angular radii and the SEDs reconstructed from available photometric and spectrophotometric data, we find effective temperatures ranging from $T_{\\rm eff}=3192 \\pm 200$ K at phase $\\Phi=0.13$ to $2918 \\pm 183$ K at $\\Phi=0.26$. Comparison of these Rosseland radii, effective temperatures, and the shape of the observed visibility functions with model predictions suggests that $o$ Cet is a fundamental mode pulsator. Furthermore, we investigated the variation of visibility function and diameter with phase. The Rosseland angular diameter of $o$~Cet increased from $28.9 \\pm 0.3$~mas (corresponding to a Rosseland radius of $332 \\pm 38 ~R_{\\odot}$ for a distance of $D=107\\pm12~{\\rm pc}$) at $\\Phi=0.13$ to $34.9 \\pm 0.4$~mas ($402 \\pm 46~R_{\\odot}$) at $\\Phi=0.4$. The error of the Rosseland linear radius almost entirely results from the error of the parallax, since the error of the angular diameter is only approximately 1\\%. ",
        "introduction": "Mira stars are long-period variables (LPVs) which evolve along the asymptotic giant branch (AGB), and are characterized by stellar pulsation with amplitudes as large as $\\Delta V \\sim 9$ and well-defined pulsation periods (80-1000 days). In recent years, the comparison of theoretical pulsation models with MACHO observations of LPVs in the LMC, in particular the reproduction of period ratios in multimode pulsators, has shown that Miras are fundamental-mode pulsators \\citep{woodM}. However, radius measurements of Mira variables when compared to theoretical pulsation calculations have generally yielded the large values expected for first overtone pulsators (e.g., \\citealt{Feast,VANB}). There is clearly a problem with the interpretation of radius measurements that needs examination.\\\\ High-resolution interferometric studies of Mira stars allow the determination of the size of the stellar disk, its center-to-limb intensity variation, surface inhomogeneities, and the dependence of diameter on wavelength and variability phase (see, e.g., \\citealt{PEA}; \\citealt{BON}; \\citealt{LAB77}; \\citealt{BON82}; \\citealt{KAR}; \\citealt{HAN92}; \\citealt{QUI}; \\citealt{WIL92}; \\citealt{TUT94}; \\citealt{DAN}; \\citealt{HAN95}; \\citealt{WEI96}; \\citealt{VANB}; \\citealt{BURNS}; \\citealt{PER}; \\citealt{HOF00}; \\citealt{WEI00}; \\citealt{WEIN}; \\citealt{THOM}; \\citealt{MEIS}). The results of such interferometric observations can be compared with predictions from theo\\-retical models of stellar pulsation and the atmosphere of Mira stars (e.g., \\citealt{WAN}; \\citealt{SCHO}; \\citealt{BES}; \\citealt{BSW} = BSW;  \\citealt{HSW} = HSW; \\citealt{TLSW} = TLSW; \\citealt{ISW} = ISW). Confrontation of detailed theoretical models with high-resolution observations is crucial for improving our understanding of the physical properties of Mira stars (e.g., \\citealt{HOFRCAS}; \\citealt{HOF01}; \\citealt{WEI03}; \\citealt{SCH03}). \\\\ In this paper we present ESO VLTI/VINCI visibility measurements of $o$~Cet and compare the measured visibility shape and the phase dependence of the visibility with model predictions. $o$~Cet, the prototype of oxygen-rich Mira stars, is a very suitable object for these studies, since VINCI observations exist for different phases and baselines, its distance is known (revised HIPPARCOS distance $107.06\\pm12.26~{\\rm pc}$, \\citealt{KNAPP}), and a large amount of spectroscopic and photometric data is available for different phases. ",
        "conclusions": "We presented \\textit{K}-band observations of $o$~Cet obtained with the VLTI and its beam combiner instrument VINCI. From these VINCI observations at six different epochs we derived Rosseland angular radii using 5 different theoretical model series from BSW, HSW, TLSW, and ISW. Using the derived Rosseland angular diameter and the SEDs reconstructed from various photometry and spectrophotometry data, we obtained effective temperatures ranging from $T_{\\rm eff}=3192 \\pm 200$~K at $\\Phi=0.13$ to $2918 \\pm 183$~K for $\\Phi=0.26$.\\\\ We found that there is fair agreement between the Rosseland linear radii derived from the observed visibilities and those predicted by the fundamental  mode pulsation model series P and M, while there is no agreement for other models. The nonlinear pulsation models all start from a static ``parent'' star. The parent star for the best-fitting model series, the fundamental-mode P series, has a radius of $\\sim$240 R$_{\\odot}$. It is clear from Fig. 4 that $o$~Cet, and indeed the nonlinear pulsation models from the P series, nearly all have radii of $\\sim$300-400 R$_{\\odot}$. Thus, the effect of large amplitude pulsation is to expand the surface layers of the star so that its apparent radius is considerably larger at most phases than the radius it would have if static (note that this effect is milder with respect to a continuum radius like $R_{1.04}$ ; cf. Table \\ref{models2} and the discussion in ISW). This expansion of the apparent radius does not greatly affect the interior of the model or the pulsation period. Hence, if one compares observed radii of Miras in (say) a radius-period diagram with the radii and pulsation periods of the parent stars, then one will clearly not find agreement. This is the reason that the Miras were thought for so long to be first-overtone pulsators (see, e.g., \\citealt{FEREV,TUCH}, and references therein). It is only by a detailed comparison of interferometric angular diameter measurements with models of large-amplitude, pulsating atmospheres that this problem has been solved.\\\\ On the other hand, the effective temperatures derived from the observations are very close to the effective temperatures predicted by the D-model series but higher than those predicted by the P and M~models. Given the definition of effective temperature $L = 4 \\pi\\sigma R^{2} T^{4}_{\\rm eff}$, this tells us that the P-series models are too low in luminosity, consistent with the luminosities derived in section \\ref{teff} for $o$ Cet (compared to the model values in Table \\ref{models2}). The shape of the measured visibilities for $o$~Cet at phase $\\Phi=0.13$ are best fitted with the P~model series, whereas all other model series and simple UD models show much poorer agreement with the observations. Taking all this into account, it is clear that a higher-luminosity, fundamental-mode model series is required for a more accurate modeling of $o$ Cet.\\\\ Furthermore, we found that the observed visibility functions and diameters change considerably from phase 0.13 to phase 0.40. The Rosseland angular diameter of $o$~Cet increases from $d_{\\rm{Ross}}^{\\rm{a}} = 28.9 \\pm 0.3$~mas (corresponding to a Rosseland linear radius of $R_{\\rm Ross} = 332 \\pm 38 ~R_{\\odot}$) at $\\Phi=0.13$ to $d_{\\rm{Ross}}^{\\rm{a}} = 34.9 \\pm 0.4$~mas ($R_{\\rm Ross} = 402 \\pm 46 ~R_{\\odot}$) at $\\Phi=0.4$. Thus, the diameter of $o$~Cet increases by 18\\% between these two phases, which is in good agreement with the approximately 14\\% diameter increase derived from linear interpolation of the results by \\cite{THOM} for the oxygen-rich Mira S~Lac."
    },
    "0404/astro-ph0404281_arXiv.txt": {
        "abstract": "Here we present a self-consistent stationary solution for spherically symmetric winds driven by massive star clusters under the impact of radiative cooling. We demonstrate that cooling may modify drastically the distribution of temperature if the rate of injected energy approaches a critical value. We also prove that the stationary wind solution does not exist whenever the energy radiated away at the star cluster center exceeds $\\sim 30\\%$ of the energy deposition rate. Finally we thoroughly discuss the expected appearance of super-star cluster winds in the X-ray and visible line regimes. The three solutions here found: the quasi-adiabatic, the strongly radiative wind and the inhibited stationary solution, are then compared to the winds from Arches cluster, NGC 4303 central cluster and to the supernebula in NGC 5253. ",
        "introduction": "In the stationary solution for spherically symmetric winds (Chevalier and Clegg 1985; hereafter referred to as CC85) as well as in the former approach of Holzer and Axford (1970) and in the more recent numerical calculations of Cant{\\'o} et al. (2000) and Raga et al. (2001) the flow has been assumed to be adiabatic and thus predicts a very extended X-ray envelope around the sources. The impact of cooling on the stationary wind solution, was discussed by Silich et al. (2003, hereafter referred to as Paper I) for winds driven by powerful and compact stellar clusters, and by Wang (1995) for gas outflows from galaxies. Winds driven by compact star clusters establish a temperature distribution radically different from that predicted by the adiabatic solution, bringing the X-ray emitting boundary much closer to the star cluster surface. However, in none of the above studies, the effects of radiative cooling within the star forming volume itself were taken into consideration.  Here we present a self-consistent semi-analytical model of stationary winds driven by massive stellar clusters taking full account of radiative cooling (see sections 2 and 3). We first discuss how to find proper wind central values and then use them to integrate numerically the basic equations. We also indicate the threshold value of the energy deposition rate above which a stationary solution is inhibited. In sections 4 and 5, the three regimes found when radiative cooling is considered: the quasi-adiabatic, the strongly radiative wind and the inhibited stationary wind, are then compared to well observed examples. Our conclusions are given in section 6. ",
        "conclusions": "We have developed a self-consistent stationary solution for spherically symmetric winds driven by compact star clusters taking into consideration radiative cooling.  We have shown that stationary radiative winds differ strongly from their adiabatic counterparts. In particular we have shown that in the energy-size plane, there is a regime where the stationary wind solution is inhibited. This occurs whenever the energy radiated away per unit volume and per unit time ($n^2_c \\Lambda(Z_w,T_w)$) surpasses a value of $\\sim 30\\%$ of the energy injection rate. In this catastrophic cooling regime, the sonic point cannot be accommodated at the star cluster surface and the stationary wind solution does not exist. Below such a limit the flow, despite radiative cooling, behaves within the star cluster volume in a quasi-adiabatic manner and is able to set the sonic point at the star cluster boundary and evolve into a stationary wind.  Stationary winds driven by stellar clusters with an energy input rate or a size that approaches the critical value, establish a temperature distribution radically different from that predicted by the adiabatic solution. In these stationary wind cases the fast fall of temperature brings the boundaries of the X-ray zone, and of the line cooling zone and the photoionized envelope, closer to the star cluster center. This promotes the establishment of a compact ionized gaseous envelope which should be detected as a week and broad ($\\sim 1000$ km s$^{-1}$) emission line component at the base of a much narrower line caused by the central HII region.  Note that the threshold energy input rate approaches an asymptotic value for large values of $R_{sc}$ (see Figure 3). This implies that single supermassive star clusters are not able to generate stationary outflows whatever their radii may be. The fate of the ejected material in this case remains unclear. A self-regulating star forming region may form and may keep the injected gas bound because of catastrophic radiative cooling or the gas may be blown away in a quasi-recurrent regime. The outflows driven by supermassive or super-compact star clusters should be studied with a full non-stationary hydrodynamic approach.  We have speculated that the super-nebula in NGC 5253 seems a good example of this inhibited stationary wind regime. Radiative cooling enforces a rapid drop in the sound speed value and the injected matter is to remain near the cluster. In such a case we predict that the metallicity of the super-nebula is above solar, making the cluster lie above the threshold limit for stationary winds (Figures 3 and 5).  Our calculations show that the Arches cluster wind seems to evolve in the quasi-adiabatic regime and predict the H$_{\\alpha}$ and Br$\\gamma$ broad component luminosities around $L_{H\\alpha} \\approx 5 \\times 10^{34}$ erg s$^{-1}$ and $L_{Br\\gamma} \\approx 5 \\times 10^{32}$ erg s$^{-1}$, respectively.  The temperature distribution derived for the NGC 4303 central 1.55~pc star cluster wind is radically different from the adiabatic temperature distribution even for a solar wind metallicity. The calculated X-ray luminosity is in reasonable agreement with the observed diffuse component luminosity. The radiative model also predicts a compact (between 6 pc and 30 pc) broad line emission with $L_{H\\alpha} \\approx 10^{36}$ erg s$^{-1}$ and L$_{Br\\gamma} \\approx 10^{34}$ erg s$^{-1}$, respectively.  We are pleased to thank D. Strickland who provided us with his X-ray emissivity tables. Our thanks also to C. Mu\\~noz-Tu\\~n\\'on for multiple suggestions and to C. Law for a useful discussion about the Arches cluster observational parameters during the X-ray - radio connection Santa Fe workshop. We thank prof. J. Palou\\v s for his comments and suggestions as a referee. We also appreciate the financial support given by M\\'exico (CONACYT) research grant 36132-E."
    },
    "0404/astro-ph0404554_arXiv.txt": {
        "abstract": "A starquake mechanism for pulsar glitches is developed in the solid quark star model. It is found that the general glitch natures (i.e., the glitch amplitudes and the time intervals) could be reproduced if solid quark matter, with high baryon density but low temperature, has properties of shear modulus $\\mu=10^{30\\sim 34}$ erg/cm$^3$ and critical stress $\\sigma_{\\rm c}=10^{18\\sim 24}$ erg/cm$^3$. The post-glitch behavior may represent a kind of damped oscillations.  \\vspace{5mm} \\noindent {\\it PACS codes:} 97.60.G, 97.60.J, 11.80.F ",
        "introduction": "Pulsars are unique objects, with which all types of elemental interaction could be tested extremely. However, the most elementary question relevant is still open: {\\em What is the nature of pulsars?} It is conventionally thought that pulsars are simply a kind of boring big ``nuclei'' --- neutron stars, but more and more attention is paid to the quark star model for pulsars \\citep{xzq02,xu03a} since {\\em no} convincing work, neither in theory from first principles nor in observation, has confirmed Baade-Zwicky's original idea that supernovae produce neutron stars. The bare quark surface is suggested to be a new probe for identifying quark stars with strangeness, and possible observational evidence for bare strange stars appears: the drifting subpulses of radio pulsars, ultra-high luminosity of soft $\\gamma$-ray repeaters, non-atomic thermal spectra of isolated ``neutron'' stars \\citep{xu03b}. However, can the bare strange star model reproduce most of the general features of pulsars (especially glitches)?  The observation of free precession in PSR B1828-11 \\cite{sls00} and PSR B1642-03 \\citep{slu01} challenges astrophysicists today to re-consider the internal structure of radio pulsars \\citep{horv04}. The current model for glitches involves neutron superfluid vertex pinning and the consequent fluid dynamics. However, the pinning should be much weaker than predicted in the glitch models, otherwise the vortex pinning will damp out the precession on timescales being much smaller than observed. In addition the picture, that a neutron star core containing coexisting neutron vertices and proton flux tubes, is also inconsistent with observations of freely precessing pulsars \\citep{link03}. It is then supposed that the hydrodynamic forces presented in a precessing star are probably sufficient to unpin all of the vortices of the inner crust \\citep{lc02} since a definitive conclusion on the nature of vertex pinning has not been reached yet due to various uncertainties in the microscopic physics. But recently, Levin \\& D'Angelo \\citep{ld04} studied the magnetohydrodynamic (MHD) coupling between the crust and the core of a rotating neutron star, and found that the precession of PSR B1828-11 should decay over a human lifetime. This well-defined MHD dissipation should certainly be important in order to test the stellar models.  An alternative way to understand both glitch and free-precession could be through the suggestion that radio pulsars are solid quark stars \\citep{xu03,xu04}. A solid quark star is just a rigid-like body, no damping occurs, and the solid pulsar model may survive future observational tests if the free precession keeps the same over several tens of years. A neutron star could not be in a solid state, whereas a cold quark star could be. Such a solid state of quark matter could be very probably Skyrme-like\\footnote{% Skyrme \\citep{skyrme62} considered baryons as solitons. The $n-$quark clusters might also be described as solitons in a similar way. } \\citep{ob99,lee03}, % the study of which may help us to understand dense quark matter with low temperature.  Fluid strange-star (even with possible crusts) models were noted to be inconsistent with the observations of pulsar glitches more than one decade ago \\citep{alpar87}. Modifications with the inclusion of possible stable particles to form a differentiated structure of so-called strange pulsars was also suggested \\citep{benv90}, but is not popular because of a disbelief in the employed physics \\citep{hps93,horv04}. However, can a fully solidified quark star proposed \\citep{xu03} really reproduce the glitch behaviors observed? One negative issue is that giant glitches are generally not able to occur at an observed rate in a solid neutron star \\citep{bp71}. Nonetheless, more strain energy could be stored in a solid quark star due to an almost uniform distribution of density (the density near the surface of a bare strange star is $\\sim 4\\times 10^{14}$ g/cm$^3$) and high shear modulus introduced phenomenologically for solid quark matter with strangeness. More energy is then released, and this may enable a solid pulsar to glitch frequently with large amplitudes. Furthermore, the post-glitch behaviors may represent damped vibrations. In this paper, we try to model glitches in a starquake scenario of solid quark stars. ",
        "conclusions": "A starquake model for pulsar glitches is developed in the regime of solid quark stars, and it is found that the general glitch behaviors (i.e., the glitch amplitude $\\Delta\\Omega/\\Omega$ and the time interval $t_{\\rm q}$) could be reproduced if solid quark matter has properties of shear modulus $\\mu=10^{30\\sim 34}$ erg/cm$^3$ and critical stress $\\sigma_{\\rm c}=10^{18\\sim 24}$ erg/cm$^3$. It is suggested that the post-glitch process could be described as damped oscillations, especially in the critical and the overdamped cases. Anyway, this is only a primary and simplified study of quakes in solid quark stars, more elaborate work, with possible modifications, on both quake and postquake processes is necessary in order to understand the nature of solid quark matter through glitching pulsars.  We are dealing with solid quark stars in this paper. The quark Cooper pairing of the BCS type is suggested in quark matter of low-temperature but high baryon density, which may result in a color superconducting state \\cite{arw98}, with a large pairing gap on the order of 100 MeV. This kind of condensation in momentum space takes place in case of same Fermi momenta; whereas ``LOFF''-like state may occur if the Fermi momenta of two (or more) species of quarks are different \\cite{abr02}. For 3 flavors of massless quarks, all nine quarks pair in a pattern which locks color and flavor symmetries, as called color-flavor locking (CFL) state \\cite{arw99}. However, for such quark matter, there exists a {\\em competition} between color superconductivity and solidification, just like the case of laboratory low-temperature physics. One needs weak-interaction and low-mass in order to obtain a quantum fluid before solidification. This is why only helium, of all the elements, shows superfluid phenomenon though other noble elements have similar weak strength of interaction due to filled crusts of electrons. The strong color interaction (and the Coulomb interaction in the system with strangeness) may be responsible for a possible solidification of dense quark matter with low temperatures. Further experiments (in low-energy heavy ion colliders) may answer whether quark matter is in a state of solid or color-superconductivity. Can a solid neutron star be possible? The answer might be {\\em no}, because at least the part of neutron matter with approximate nuclear saturation density should be in a fluid state. In this sense, only solid quark matter is possible, and a quark star is identified if one convinces that a pulsar is in a solid state.  We have assumed that the entire strain energy $E_{\\rm strain}$ is relieved in a quake, which results in the reference oblateness of Eq.(\\ref{epsiloni}). However, it is possible that not entire, but most of, the energy $E_{\\rm strain}$ is released in a real situation, and the actual reference points are near but larger than $\\varepsilon_i$. Of course, there is a tendency of $\\varepsilon\\rightarrow \\varepsilon_i$ after the $i-$th quake, but the effective shear modulus, $\\mu_{\\rm eff}$, of matter broken could be much smaller than that of perfect elastic solid, $\\mu$. The recovery timescale could be $\\tau \\sim 15$ days if $\\mu_{\\rm eff}$ is order of $10^{15}$ erg/cm$^3$.  A large Vela glitch on 2000 January 16.319 was noted \\citep{dmc00}, and {\\em Chandra} observations were carried out $\\sim 3.5$ and $\\sim 35$ days after the glitch \\citep{hgh01}, but no temperature change expected in conventional models with released thermal energy of $\\sim 10^{42}$ ergs is detected. This could be understood in this starquake model since (1) the thermal conductivity of quark matter is much larger than that of hadron matter and (2) the thermal energy released, $E_{\\rm therm}$, to be much smaller than $10^{42}$ ergs is possible.  {\\em Acknowledgments}: This work is supported by National Nature Sciences Foundation of China (10273001), and by the Special Funds for Major State Basic Research Projects of China (G2000077602). The valuable suggestions from an anonymous referee are sincerely acknowledged."
    },
    "0404/astro-ph0404412_arXiv.txt": {
        "abstract": "We present the results of an interferometric study of 38 millimeter-wave lines of $^{12}$CH$_3$OH in the vicinity of the massive star forming region W3(OH/H$_2$O). These lines cover a wide range of excitation energies and line strengths, allowing for a detailed study of excitation mechanisms and opacities. In this paper we concentrate on the region around the water maser source W3(H$_2$O) and a region extending about 30 arcsec to the south and west of the hydroxyl maser source W3(OH).  The methanol emitting region around W3(H$_2$O) has an extent of 2.0 x 1.2 arcsec (4400 x 2600 AU). The density is of order 10$^7$ \\cc, sufficient to thermalize most of the methanol lines. The kinetic temperature is approximately 140 K and the methanol fractional abundance greater than 10$^{-6}$, indicative of a high degree of grain mantle evaporation. The W3(H$_2$O) source contains sub-structure, with peaks corresponding to the TW source and Wyrowski's B/C, separated by 2500 AU in projection. The kinematics are consistent with these being distinct protostellar cores in a wide binary orbit and a dynamical mass for the region of a few tens of M$_{\\sun}$.  The extended methanol emission to the southwest of W3(OH) is seen strongly only from the lowest excitation lines and from lines known elsewhere to be class I methanol masers, namely the 84.5 GHz 5$_{-1}$--4$_0$ E and 95.2 GHz 8$_0$--7$_1$ A$^+$ lines. This suggests that this region, like class I maser sources, is dominated by collisional excitation. Within this region there are two compact clumps, which we denote as swA and swB, each about 15 arcsec (0.16 pc projected distance) away from W3(OH). Excitation analysis of these clumps indicates the presence of lines with inverted populations but only weak amplification. The sources swA and swB appear to have kinetic temperatures of order 50--100 K and densities of order 10$^5$--10$^6$ \\cc. The methanol fractional abundance for the warmer clump is of order 10$^{-7}$, suggestive of partial grain mantle evaporation. The clumping occurs on mass scales of order 1 M$_{\\sun}$. ",
        "introduction": "Methanol is an abundant interstellar molecule, especially in regions of star formation where its high abundance is thought to be the result of thermal evaporation of dust grain mantles \\citep{CHH93}. A slightly asymmetric rotor with hindered internal rotation and significant a- and b-axis dipole moments, methanol has a complex spectrum which is sensitive to interstellar conditions.  Observable lines cover a wide range of energies and line strengths, allowing for detailed analysis of excitation conditions and opacities.  Two classes of interstellar methanol masers are known. Class I methanol masers, produced by collisional pumping, are seen in regions of massive star formation, but are generally well separated from the centers of activity as traced by embedded infrared sources and ultracompact \\HII\\ (\\UCHII) regions. Class II methanol masers are thought to be produced by radiative excitation at infrared wavelengths and are spatially well correlated with young stellar objects, \\UCHII\\ regions, and OH masers. Even when not masing, methanol may be expected to exhibit distinctly non-LTE excitation, except under the highest density conditions where collisions could thermalize the level populations.  The southeastern portion of W3 contains the hydroxyl maser source W3(OH), which is associated with an \\UCHII\\ region around a young O7 star, and the nearby water maser source W3(H$_2$O), which contains a young stellar object known as the TW object \\citep{TW84}. Numerous studies have been made of molecular material associated with these two compact objects and throughout the region surrounding them \\citep{W91,W93,W94,HvD97,W99,N00b}. W3 is thought to be about 2.2 kpc from the Sun \\citep{H78}. At such a distance, resolution of order 1 arcsec is necessary to study the nature of protostellar cores on scales of order 1000 AU.  \\cite{S01}\\ presented results on new class II methanol masers in W3(OH) and discussed excitation conditions needed to explain observed fluxes in those and several other methanol maser-candidate lines. \\cite{S03}\\ extended that analysis to include absorption and emission components in a larger set of methanol lines in W3(OH) and discussed the implications in terms of excitation and source structure. Here we wish to complete our analysis of methanol in the vicinity of W3(OH) and W3(H$_2$O) by presenting our complete set of millimeter-wave observations and by discussing the nature of the methanol associated with the W3(H$_2$O) source and methanol in other nearby regions. ",
        "conclusions": "To the southwest of W3(OH) there is an extensive region of gas, dense enough to be seen in a number of methanol lines. Emission is most readily seen in lines of low upper state energy (E$_\\mathrm{u}$ $\\lesssim $ 15 K) and in lines known elsewhere for their class I methanol maser emission. In this case, strong class I masers do not seem to be present, but the physical conditions are similar to those seen in class I maser regions.  Excitation is collisionally driven, with densities of order 10$^5$ - 10$^6$ \\cc. Inverted populations and weak amplification appear to be present. We predict that the 36 GHz and 44 GHz lines will also be inverted, although it is doubtful that their optical depths will be sufficient to produce strong masers. The region is clumpy on scales of about 1 to a few tens of solar masses. Kinetic temperatures range up to about 100 K. The warmest regions have methanol fractional abundances of order 10$^{-7}$, suggesting that some grain mantle evaporation has taken place, but that the evaporation is incomplete.  In contrast, towards W3(H$_2$O) the methanol emission is nearly thermalized and many lines are optically thick. None of the lines studied have inverted populations. Excitation is due to a combination of collisional and radiative processes. The size of the methanol emitting region is about 2600 x 4400 AU. This is resolved into at least two components, corresponding to the TW object and a blend of Wyrowski's B and C components. It is likely that both are massive young stellar objects. Their separation and line of sight velocity difference are consistent with a wide binary orbit with a period of order 50,000 years.  Each source may be accompanied by circumstellar material containing of order one solar mass or less of gas. There may be as much as a few tens of M$_{\\sun}$ of circumbinary material, although its distribution is unclear. Methanol fractional abundance in the circumstellar material is of order 10$^{-6}$, indicative of a warm environment and extensive grain mantle evaporation."
    },
    "0404/nlin0404058_arXiv.txt": {
        "abstract": "We use the Smaller Alignment Index (SALI) to distinguish rapidly and with certainty between ordered and chaotic motion in Hamiltonian flows. This distinction is based on the different behavior of the SALI for the two cases: the index fluctuates around non--zero values for ordered orbits, while it tends rapidly to zero for chaotic orbits. We present a detailed study of SALI's behavior for chaotic orbits and show that in this case the SALI exponentially converges  to zero, following a time rate depending on the difference of the two largest Lyapunov exponents $\\sigma_1$, $\\sigma_2$ i.e. $\\mbox {SALI} \\propto e^{-(\\sigma_1-\\sigma_2)t}$. Exploiting the advantages of the SALI method, we demonstrate how one can rapidly identify even tiny regions of order or chaos in the phase space of Hamiltonian systems of 2 and 3 degrees of freedom. ",
        "introduction": "\\label{intro} Knowing whether the orbits of a dynamical system are ordered or chaotic is fundamental for the understanding of the behavior of the system. In the dissipative case, this distinction is easily made as both types of motion are attracting. In conservative systems, however, distinguishing between order and chaos is often a delicate issue (e.g.\\ when the chaotic or ordered regions are small) especially in systems with many degrees of freedom where one cannot easily  visualize the dynamics. For this reason it is of great importance to have quantities that determine if an orbit is ordered or chaotic, independent of the dimension of its phase space.     The well--known and commonly used method for this purpose is the evaluation of the maximal Lyapunov Characteristic Exponent (LCE) $\\sigma_1$. If $\\sigma_1 > 0$ the orbit is chaotic. Benettin et al.\\ \\cite{BGGS80a} studied theoretically the problem of the computation of all LCEs and proposed in \\cite{BGGS80b} an algorithm for their numerical computation. In particular, $\\sigma_1$ is computed as the limit for $t \\rightarrow \\infty$ of the quantity \\begin{equation} L_t=\\frac{1}{t}\\, \\ln  \\frac{|\\vec{w}(t)|}{|\\vec{w}(0)|}\\, ,\\, \\mbox{i.e.}\\,\\, \\sigma_1 = \\lim_{t\\rightarrow \\infty} L_t \\, , \\label{eq:lyap} \\end{equation} where $\\vec{w}(0)$, $\\vec{w}(t)$ are  deviation vectors from a given orbit, at times $t=0$ and $t>0$ respectively. The time evolution of $\\vec{w}$ is given by solving the so--called {\\it variational equations} (see Sec. \\ref{Behavior}). Generally, for almost all choices of initial deviations $\\vec{w}(0)$, the limit for $t \\rightarrow \\infty$ of Eq.~(\\ref{eq:lyap}) gives the same $\\sigma_1$.   In practice, of course, since the exponential growth of $\\vec{w}(t)$ occurs for short time intervals, one stops the evolution of $\\vec{w}(t)$ after some time $T_1$, records the computed $L_{T_1}$, normalize vector $\\vec{w}(t)$ and repeats the calculation for another time interval $T_2$, etc. obtaining finally $\\sigma_1$ as an average over many $T_i$, $i=1,2,\\ldots,N$ as \\begin{displaymath} \\sigma_1 = \\frac{1}{N} \\sum_{i=1}^{N} L_{T_i} \\end{displaymath} The basic problem of the computation of $\\sigma_1$  is that, after every $T_i$, the calculation starts from the beginning and may yield an altogether different $L_{T_i}$ than the $T_{(i-1)}$ interval. Thus, since $\\sigma_1$ is influenced by the whole evolution of $\\vec{w}(0)$, the time needed for $L_t$ (or the $L_{T_i}$) to converge is not known a priori and may become extremely long. This makes it often difficult to tell whether $\\sigma_1$ finally tends to a positive value (chaos) or converges to zero (order).  In recent years, several methods have been introduced which try to avoid this problem by studying  the evolution of deviation vectors, some of which are briefly discussed in Sec.\\ \\ref{Compare}. In the present paper, we focus our attention on the method of the Smaller Alignment Index (SALI)~\\cite{SALI}, performing a systematic study of its behavior in the case of autonomous Hamiltonian systems with 2 (2D) and 3 (3D) degrees of freedom. This method has been applied successfully to several 2--dimensional (2d) and multidimensional maps \\cite{SALI}, where SALI was found to converge rapidly to zero for chaotic orbits, while it exhibits small fluctuations around non--zero values for ordered orbits. It is exactly this ``opposite\" behavior of the SALI which makes it an ideal indicator of chaoticity: Unlike the maximal LCE, it does not start at every step a new calculation of the deviation vectors, but takes into account information about their convergence on the unstable manifold from all the previous steps. The method has already been used successfully as a chaos detection tool in some specific Hamiltonian systems \\cite{SABV03c,PBS04,S03,SESS04,VKS02,KVC04}, although some authors \\cite{VKS02,KVC04} use different names for the SALI.   The paper is organized as follows: In Sec.\\ \\ref{Ap} we recall the definition of the SALI and present  results distinguishing between ordered and chaotic motion in 2 and 3--degrees of freedom (2D and 3D) Hamiltonians, comparing also the efficiency of the SALI with the computation of $\\sigma_1$. In Sec.\\ \\ref{Behavior} we explain the behavior of the SALI  for ordered and chaotic orbits, showing that in the latter case SALI converges exponentially to zero following a rate which depends on the difference of the two largest Lyapunov exponents $\\sigma_1$ and $\\sigma_2$. In Sec.\\ \\ref{Discussion} we demonstrate the ability of the method to reveal the detailed structure of the dynamics in the phase space. In Sec.\\ \\ref{Compare} we compare the SALI method with some other known methods of chaos detection and in Sec.\\ \\ref{Summary} we summarize our results. ",
        "conclusions": "\\label{Summary}  In this paper we have applied the SALI method to distinguish between order and chaos in  2D and  3D autonomous Hamiltonian systems, and have also analyzed the behavior of the index for chaotic  orbits. Our results  can be summarized as follows.  \\begin{itemize}   \\item The SALI proves to be an ideal indicator of chaoticity independent of the dimensions of the system. It tends to zero for chaotic orbits, while it exhibits small fluctuations around non--zero values for ordered ones and so it clearly distinguishes between these two cases. Its  advantages are its simplicity, efficiency and reliability as it can rapidly and accurately determine the chaotic vs ordered nature of a given orbit. In regions of `stickiness', of course, along the borders of ordered motion it displays transient oscillations. However, once the orbit enters a large chaotic domain the SALI converges exponentially to zero, often at  shorter times than it takes the maximal Lyapunov exponent to converge to its limiting value.  \\item We emphasize that the main advantage of the SALI in chaotic regions is that it uses two deviation vectors  and exploits at every step, their convergence to the unstable manifold from all previous steps. This allows us to show that the SALI tends to zero for chaotic orbits at a rate which is related to the difference of the two largest Lyapunov characteristic exponents $\\sigma_1$, $\\sigma_2$ as $\\mbox{SALI} \\propto e^{-(\\sigma_1-\\sigma_2)t}$. By comparison, the computation of the maximal LCE, even though it requires only one deviation vector and one exponent, $\\sigma_1$, often takes longer to converge, since it needs to average over many time intervals, where the calculation of this exponent is independent from all previous intervals. The SALI was also proved to have similar or even better performance than other methods of chaos detection which were briefly discussed in Sec.\\ \\ref{Compare}.  \\item The $\\mbox{SALI} \\in [0,\\sqrt{2}]$ and its value characterize an orbit of being chaotic or ordered. Exploiting this feature of the index  we have plotted detailed  phase space portraits both for 2D and 3D Hamiltonian systems, where the chaotic and ordered regions are clearly distinguished. We were thus able to trace in a fast and systematic way very small islands of ordered motion, whose detection by traditional methods would be very difficult and time consuming. This approach is therefore expected to provide useful tools for the location of stable periodic orbits, or the computation of the phase space volume occupied by ordered or chaotic motion in multidimensional systems, where the PSS is not easily visualized, and very few other similar techniques of practical value are available. \\end{itemize}   \\ack We acknowledge fruitful discussions on the contents of this work  with Professors Giulio Casati and Tomas Prosen. We would also  like to thank the anonymous referee for very useful comments which helped us improve the clarity of the paper. This research was partially supported by the `Heraclitus' research program of the Greek Ministry of Development. Ch.\\ Skokos was supported by the `Karatheodory' post--doctoral fellowship No 2794 of the University of Patras and Ch.\\ Antonopoulos was supported by the `Karatheodory' graduate student fellowship No 2464 of the University of Patras."
    },
    "0404/astro-ph0404468_arXiv.txt": {
        "abstract": "Recent attempts to fit Type Ia supernova data by modeling the dark energy density as a truncated Taylor series have suggested the possibility of {\\em metamorphosis}, i.e., a rapidly evolving equation of state parameter, $w_{DE}(z)$. However, we show that fits using that parametrization have significant problems. Evolution of $w_{DE}(z)$ is both favoured and in some sense forced, and the equation of state parameter blows up or diverges in large regions of the parameter space. To further elucidate these problems we have simulated sets of supernova data in a $\\Lambda$--universe to show that the suggested ``evidence'' for metamorphosis is also common for $w_{DE}=-1$. ",
        "introduction": "Revealing the true nature of dark energy (DE) has become one of the most important tasks in cosmology. Considering the plethora of DE models proposed in the literature, a model independent reconstruction of DE would be an appealing alternative to testing all models separately. In two recent papers~\\cite{alam:dec03,alam:mar04} attempts were made by Alam \\etal to reconstruct the dark energy equation of state parameter $w_{DE}(z)$ in a model independent manner, using the latest supernova data~\\cite{hzt:tonry,hzt:barris,hzt:riess,scp:knop}. In these two papers a truncated Taylor series was used to model the dark energy density $\\rho_{DE}(z)$.  The results indicate an evolution of $w_{DE}(z)$, a behaviour they call metamorphosis. From the reported analysis, it would seem that this is a significant effect prompting for ``exotic'' models for the DE.  Since other parameterizations~\\cite{hzt:tonry,hzt:riess,scp:knop} suggest that the Type Ia supernova data collected so far are consistent with the simplest DE model of all, that of a cosmological constant $\\Lambda$ ($w_{DE}=-1$), it is important to investigate how such different conclusions can be reached starting from the same sets of data.  In this paper we argue that the method of model independent reconstruction proposed by Alam \\etal suffers from a number of serious shortcomings. For alternative methods of DE reconstruction see e.g. references~\\cite{wang1,wang2,wang3,huterer,daly}.  A fundamental requirement on a model independent reconstruction of $w_{DE}(z)$ must be that all DE alternatives are treated on an equal footing. Although the method of Alam \\etal at first sight seems to be capable of an exact reproduction of the equation of state parameter of the cosmological constant, it actually favours evolving $w_{DE}(z)$. The confidence contours, describing the level to which the reconstructed $w_{DE}(z)$ is known, exhibit two related problems. First, contours enclosing regions of high confidence level (CL) and high redshift tend to diverge. Second, by construction, the low level CL regions shrink for high redshifts. ",
        "conclusions": "The ansatz proposed by Alam \\etal may be useful for modelling the dark energy density, but its usefulness for revealing the nature of the DE seems limited. The parametrization of the dark energy equation of state parameter based on a truncated Taylor series involves a number of severe problems. Evolution is both favoured and forced by this parametrization. The cosmological constant is thus mistreated by the ansatz proposed in~\\cite{alam:dec03, alam:mar04}. Not even an extension of the ansatz seems to be able to overcome these difficulties.  The equation of state parameter expressed as in equation~(\\ref{eq:de_eos}) also diverges in large regions of the parameter space. More data, which focus the solutions to the stable regions, may solve this problem. However, the region close to the point describing a universe with a cosmological constant will be in a disfavoured part of the parameter space.  The dark energy equation of state parameter reconstructed from the data sets presented in reference~\\cite{hzt:tonry} and~\\cite{hzt:barris} is inconsistent with the cosmological constant at the $68\\%$ confidence level. Simulations show that the rapidly changing behaviour of $w_{DE}(z)$ could also be expected with this ansatz for a $\\Lambda$ universe. The scatter of the best fit parameters  can be reduced if additional data points at high redshift are added to the simulated data sets. Our best fit to real data with $16$ additional high redshift supernovae was consistent with the cosmological constant at the $68\\%$ confidence level.  We concluded that the suggested ``model independent'' method of reconstructing the dark energy equation of state parameter in \\cite{alam:dec03}, is in fact model dependent, and that the seemingly striking results are likely to be due to this deficiency.  \\ack A.G. is a Royal Swedish Academy Research Fellow supported by a grant from the Knut and Alice Wallenberg Foundation.  \\vspace{1cm}"
    },
    "0404/astro-ph0404232_arXiv.txt": {
        "abstract": "Automatic classification of variability is now possible with tools like neural networks. Here, we present two neural networks for the identification of microlensing events -- the first discriminates against variable stars and the second against supernovae. The inputs to the networks include parameters describing the shape and the size of the lightcurve, together with colour of the event. The network computes the posterior probability of microlensing, together with an estimate of the likely error. An algorithm is devised for direct calculation of the microlensing rate from the output of the neural networks.  We present a new analysis of the microlensing candidates towards the Large Magellanic Cloud (LMC). The neural networks confirm the microlensing nature of only 7 of the possible 17 events identified by the MACHO experiment. This suggests that earlier estimates of the microlensing optical depth towards the LMC may have been overestimated. A smaller number of events is consistent with the assumption that all the microlensing events are caused by the known stellar populations in the outer Galaxy/LMC. ",
        "introduction": "Microlensing is rare and out-numbered by stellar variability by at least a factor of ten thousand. Despite this, the selection of microlensing candidates in variability surveys seems straightforward at an optimistic first glance.  Unlike almost all forms of stellar variability, microlensing is achromatic, time-symmetric and does not repeat. The theoretical form of the microlensing lightcurve is well-known (e.g., Paczy\\'nski 1986) and so events can seemingly be selected by their goodness-of-fit in two passbands.  In practice, the selection of candidates is fraught with difficulties. The lightcurves are usually sparsely sampled and noisy -- for example, the median seeing at the site of one of the most prominent microlensing experiment (MACHO) is $\\sim 2.0''$.  More awkwardly still, the clear-cut set of characteristics of microlensing only holds good in the simplest case of an isolated point-mass lensing a point-source. In fact, microlensing lightcurves may show colour variations because of blending (e.g., Di Stefano \\& Esin 1995). They may show substantial deviations from time-symmetry because of parallax or xallarap effects (Dominik 1998; Mao et al. 2002) or because the lens is a binary star (e.g., Mao \\& Paczy\\'nski 1991; An et al. 2004).  As a consequence, the results of the microlensing experiments towards the Magellanic Clouds by the MACHO and EROS collaborations remain controversial (e.g., Evans 2002). From 5.7 years of data, the MACHO collaboration identified between 13 and 17 candidates towards the Large Magellanic Cloud (LMC) and reckoned that the optical depth is $1.2^{+0.4}_{-0.3} \\times 10^{-7}$ (Alcock et al. 2000). The first set of 13 events comprises the most convincing candidates, whilst the second set of 17 candidates includes an additional 4 events less firmly established. This is in astonishing contrast to the results reported by the EROS collaboration, who found just 3 events towards the LMC (Lasserre et al. 2000). The two experiments are not directly comparable as EROS monitor a wider solid angle of less crowded fields than do MACHO. Even though EROS do not analyze their data in terms of optical depth, it is clear that the results point to a lower value than that claimed by MACHO. Tellingly, a similar discord prevails in the results towards the Galactic Centre; MACHO (Alcock et al. 1997) recorded that the microlensing optical depth to the red clump stars as $3.9^{+1.8}_{-1.2} \\times 10^{-6}$, while EROS (Afonso et al. 2003b) found a value of $0.94 \\pm 0.26 \\times 10^{-6}$ at almost the same location.  These discrepancies strongly suggest that the systematic effects in the experiments are not yet fully understood, with candidate selection fingered as the most likely culprit.  All this motivated Belokurov, Evans \\& Le Du (2003) to introduce neural networks as an automatic way of classifying lightcurve shapes in massive variability surveys.  They constructed a working neural network for identification of microlensing events and applied it to microlensing data towards the Galactic Centre. In this paper, the ideas and methods of analysis are extended to the variability datasets taken towards the LMC. This is a harder problem, as the source stars are fainter and hence the microlensing events less clear-cut. A particular difficulty already identified by Alcock et al. (2000) is the contamination of samples of microlensing events by supernovae in distant galaxies behind the LMC.  \\begin{table*} \\begin{center} \\begin{tabular}{l|l|c}\\hline Variable Type & Specific Examples & Number\\\\ \\hline Eruptive      & Pre-Main Sequence, R Corona Borealis stars & 34 \\\\ Pulsating     & RV Tauris, Mira, Semi-Regular variables & 595 \\\\ \\null         & Cepheids & 372 \\\\ \\null         & Bumpers  & 300 \\\\ Cataclysmic   & Supernovae, novae, recurrent novae & 45\\\\ Eclipsing     & \\null & 135 \\\\ MACHO samples & \\null & 531 \\\\ Microlensing  & \\null & 1500 \\\\ \\hline \\end{tabular} \\end{center} \\caption{Composition of the training set. There are 1500 examples of microlensing and 2014 examples of other classes of lightcurves.  The sources for the data are reported in the main text. } \\label{table:tset} \\end{table*} ",
        "conclusions": "This paper has demonstrated the power of machine learning techniques, such as neural networks, for the classification of events in massive variability datasets. Using the specific example of the microlensing surveys, committees of neural networks have been devised to discriminate against common forms of stellar variability and against supernovae. The output of the neural network is the posterior probability of microlensing, given the prior distribution in the training set. The error on the probability can be straightforwardly calculated.  The networks have been used to process some of the data ($\\approx 22000$ lightcurves) taken towards the Large Magellanic Cloud by the MACHO collaboration (Alcock et al. 2000).  The latter authors provide a set of 13 events whose identification as microlensing is believed to be secure and a further 4 events whose identification is possible. The neural networks confirm the microlensing nature of only 7 of these possible 17 events.  Without processing the entire dataset ($\\sim 11.9$ million lightcurves), we cannot be sure that there are no events missed by Alcock et al. (2000) which would be classified as microlensing by the networks.  It is reasonable to argue that this is unlikely, as the $\\approx 22000$ MACHO lightcurves we have re-processed provide no new candidates. But, this remains a plausible speculation rather than an empirically derived fact.  Hence, we can only speculate that, as the number of events has been roughly halved, so the optical depth will be similarly reduced.  For comparison, Alcock et al. (1997) estimate the optical depths of the thin disk, thick disk and spheroid to be $2.2 \\times 10^{-8}$, whilst the optical depth of the stellar content of the LMC to be $3.2 \\times 10^{-8}$ on average. In other words, from the known stellar populations in the outer Galaxy and the LMC, the optical depth must be at least $5.4 \\times 10^{-8}$. This may well be enough to provide the 7 events whose microlensing nature we confirm.  There is supporting evidence for the belief that the known stellar populations are providing the bulk of the lenses both from the exotic events and from the lensing signal towards the Small Magellanic Cloud (SMC).  First, the exotic events yield additional information which can break some of the microlensing degeneracies and thus give indirect evidence on the location of the lens. There are two exotic events towards the LMC and two towards the SMC (Bennett et al. 1996; Palanque-Delabrouille 1998; Kerins \\& Evans 1999; Afonso et al. 2000; Alcock et al. 2001a; Evans 2002). In all cases, the exotic events favour an interpretation in which the lens lies in the Magellanic Clouds.  Additionally, Alcock et al. (2001b) imaged one of the events towards the LMC and identified the lens as a nearby low mass disk star.  Second, as Afonso et al. (2003a) point out, the duration of the events towards the SMC is very different from the duration towards the LMC. The EROS collaboration constrain the optical depth towards the SMC to be $< 10^{-7}$ at better than the 90 \\% confidence level, based on an admittedly small sample. Both these facts militate against the idea that a single population of objects in the Milky Way halo is causing the microlensing events. The mass function, internal kinematics and proper motions of the SMC and LMC are different, so that differences in the distributions of microlensing events are expected if the lenses lie predominantly in the Magellanic Clouds. Based on roughly spherical models of the dark halo, the optical depth towards the SMC is expected to be greater than that towards the LMC if the halo provides most of the lenses. Hence, the paucity of events towards the SMC is beginning to be highly problematic for halo interpretations of the events."
    },
    "0404/astro-ph0404004_arXiv.txt": {
        "abstract": "{Spectroscopic studies indicate that gas in the photospheres of young O stars moves at speeds up to the sound speed.  We show, using two-dimensional radiation MHD calculations and results from a local linear analysis, that the motions may be due to photon bubble instability if young O stars have magnetic fields.  }  \\resumen{ }  \\addkeyword{stars: early-type} \\addkeyword{stars: formation}  \\def\\msol{\\rm\\,M_\\odot} \\def\\divv{{\\bf\\nabla\\cdot v}}  \\begin{document} ",
        "introduction": "Young high-mass stars may trigger or terminate nearby star and planet formation through ionizing radiation and winds launched from their surfaces.  The properties of the radiation and winds depend on the structure of the stellar surface layers.  Stars of more than 15 Solar masses reach the main sequence while still accreting material \\citep{ys02}.  High-mass main sequence stars are thought to have stably stratified outer layers, yet gas motions are present in the atmospheres of these stars.  Absorption lines formed in the photospheres of O stars are broader than expected based on the temperature, pressure and stellar rotation. If the rotation axes are randomly oriented with respect to our line of sight, some stars will be viewed pole-on and show no rotational broadening.  However samples of O stars appear to contain no such pole-on cases \\citep{s56,ce77,p96,hs97}.  The minimum apparent equatorial rotation speed is a function of spectral type and luminosity.  Among main-sequence O stars in the Small Magellanic Cloud cluster NGC 346, the excess broadening of UV metal lines ranges from 25 km s$^{-1}$ at O2, similar to the speed of sound at the photosphere, to 5 km s$^{-1}$ at O9.5 \\citep{bl03}.  Processes that might be responsible for the additional line broadening include \\begin{enumerate} \\item {\\em Vertical velocity gradients in the stellar wind acceleration region} \\citep{k92}: Minimum line widths are similar in stars with and without strong outflows.  Also, stellar atmosphere plus wind calculations including a variation in velocity through the acceleration region produce photospheric lines similar to those in hydrostatic calculations \\citep{bl03}. \\item {\\em Global non-radial pulsations:} Low levels of spectral line shape variations in main-sequence O stars \\citep{fg96} suggest that any global non-radial modes are either of low amplitude or high order. An unresolved issue is how such modes might be excited. \\item {\\em Strange modes:} Instability occurs only for masses greater than $80\\msol$ in zero-age main sequence stars with metallicity $Z=0.02$ \\citep{gk93}, while excess photospheric line widths are observed in main-sequence stars with masses down to $20\\msol$. \\item {\\em Shaviv modes:} The instabilities are found only at luminosities greater than half the Eddington value \\citep{s01}, while excess line broadening occurs in stars with one-tenth the Eddington luminosity. \\end{enumerate}  Here we explore the possibility that a fifth mechanism, photon bubble instability \\citep{a92,g98}, may lead to small-scale motions in the surface layers of O stars.  A general local WKB analysis of linear disturbances in optically-thick radiating atmospheres indicates instability under a wide range of conditions \\citep{bs03}.  Photon bubbles are driven by the radiative flux (Figure~\\ref{fig:howitworks}), and grow if (1) Rosseland mean optical depth per wavelength is greater than unity, (2) a magnetic field is present, and (3) the radiative flux exceeds a critical value that in radiation-supported atmospheres is approximately the gas sound speed times the sum of gas and photon energy densities.  The surface layers of O stars generally satisfy the first and third criteria, but until recently were thought to be unmagnetized.  However $\\theta^1$~Ori~C, the illuminating star of the Orion Nebula and one of the nearest young O stars, shows polarimetric variations in photospheric lines with rotational phase, consistent with a 1.1~kG dipole field inclined $42^\\circ$ from the rotation axis \\citep{db02}.  Also, narrow X-ray emission lines indicate some gas near the star moves more slowly than the wind.  The line ratios show that most of the X-ray flux is thermal emission from plasma hotter than 15 million~K, too hot to arise in shocks in the wind, and as hot as coronae in some magnetically active lower-mass stars \\citep{sc03}.  Overall, there is good evidence for a magnetic field in $\\theta^1$~Ori~C.  \\begin{figure}[t!] \\vspace*{-6mm} \\hspace*{-3mm} \\includegraphics[width=0.95\\columnwidth]{f1.eps} \\vspace*{-1mm} \\caption{How photon bubbles work.  {\\em Top:} A patch of atmosphere (gray square) is initially in hydrostatic balance.  When density is reduced slightly in a small elongated region, photons diffuse more easily up the long axis.  The perturbed radiation flux exerts an extra force with a component along the magnetic field.  Gas is driven along the field to the right, out of the region of low density, and the perturbation grows.  {\\em Bottom:} As fresh gas enters the region of low density from the left, the pattern propagates to the left at about the gas sound speed.  } \\label{fig:howitworks} \\end{figure}  Photon bubbles, unlike strange modes, can grow in scattering atmospheres.  They differ from Shaviv instabilities in that they require magnetic fields and may grow when the luminosity is substantially less than the Eddington value. ",
        "conclusions": "Photon bubble instability is present in the surface layers of magnetized main-sequence O stars according to the results of a local linear WKB analysis.  We use two-dimensional radiation MHD calculations with horizontal magnetic fields and closed vertical boundaries to show that the instability can lead to small-scale movements of gas in the surface layers, with velocities about equal to the observed linewidth excesses.  In a calculation with parameters appropriate for an early O star, photon bubbles result in horizontal density variations near the photosphere.  Radiation diffuses more rapidly through the inhomogeneities than through the hydrostatic atmosphere with the same column depth, and the radiative flux is 10\\% greater than that calculated assuming hydrostatic equilibrium.  An enhanced flux may affect measurements of stellar parameters.  Issues remaining to be resolved \\linebreak \\adjustfinalcols include the dependence on the local magnetic field orientation and the shapes of the spectral lines expected from emission integrated over the stellar disk.  Photon bubbles reaching amplitudes similar to those in the hottest case we considered may eject blobs of gas through the photosphere, and could lead to density variations in the winds of O stars.  The motions might transfer energy into an initially weak magnetic field. Additional radiation MHD calculations indicate the saturation amplitude generally increases with decreasing surface gravity, so photon bubbles may reach larger amplitudes in O giants and supergiants.  Solutions of the dispersion relation show that photon bubbles may be present also in accretion disks around young massive stars."
    },
    "0404/astro-ph0404010_arXiv.txt": {
        "abstract": "We present a library of 6410 synthetic spectra with resolution $\\lambda/\\Delta\\lambda = 250\\,000$ based on the revised Kurucz 1993 model atmospheres. The library covers the wavelength range $3000 - 10\\,000$ {\\AA} with 54 values of effective temperature in the range $5250 - 50\\,000$ K, 11 values of log surface gravity between 0.0 and 5.0 and 19 metallicities in the range $-5.0$ to 1.0. We find that, with a few caveats, the library compares well with both the original 20 {\\AA} Kurucz spectra and also with observed spectra. The library is intended for use in population synthesis and physical parameterisation of stellar spectra. We assess the suitability of the library for these tasks. ",
        "introduction": "Progress in population synthesis and automatic classification of stellar spectra has been limited by the spectral resolution of the available synthetic stellar spectra. The existing synthetic libraries are not at high enough resolution to be useful for classifying stars from recent surveys such as the SDSS \\citep{stoughton02} ($\\lambda/\\Delta\\lambda\\sim1800$), or future surveys such as RAVE \\citep{steinmetz02} and GAIA \\citep{lindegren96}. Most classification techniques smooth the observed spectra to the resolution of the synthetic spectra. This means much of the detailed information in the observed spectra is lost, which may reduce the quality of the classifications. Population synthesis packages such as \\pegase\\ \\citep{fioc97} or \\textsc{gissel} \\citep{bruzual93} use a grid of stellar spectra to generate galaxy spectra. The resulting galaxy spectra are limited by the resolution of the input stellar spectra. To study the high resolution features, it is necessary to have a grid of observed or synthetic stellar spectra which match the resolution of the observed galaxy spectra. For example, galaxy spectra synthesised from 20 {\\AA} spectra cannot be used to measure standard line indices like the Lick indices.  Perhaps the most widely used library of synthetic spectra are the flux distributions from the Kurucz \\atlas\\ model atmospheres \\citep{kurucz93CD13}. It is important to note that while these are usually referred to as spectra, they are flux distributions predicted directly from the model atmospheres, rather than spectra generated by a spectral synthesis program. The Kurucz atmospheres have several disadvantages which are discussed in various sources such as \\cite{kurucz92}. However, one of their advantages is the wide range of parameters they cover, which is important for generating a grid of stellar spectra for population synthesis. The need for higher resolution spectra has been recognised for some time but, because of the immense computational expense involved, the synthesis of the spectra has been limited to partial wavelength ranges and specific regions of parameter space. Several groups have generated libraries of spectra from the Kurucz model atmospheres. For example, \\citet*{chavez97} provide a set of 711 Kurucz spectra at $\\lambda/\\Delta\\lambda=250\\,000$ in the wavelength region $4850 - 5400$ {\\AA} and \\citet{castelli01} have generated a set of 698 Kurucz spectra at $\\lambda/\\Delta\\lambda=20\\,000$, in the wavelength region $7650 - 8750$ {\\AA} for use with GAIA spectra. \\citet{gonzalez99} have generated synthetic spectra in very small spectra regions necessary for a particular application. They created a grid of synthetic profiles of stellar H Balmer and HeI lines at $\\Delta\\lambda = 0.3$ {\\AA} for the purposes of evolutionary synthesis.  There are also several libraries of observed spectra now available at much higher resolution. For example, the \\elodie\\ database \\citep{prugniel01}, consists of 709 stars observed in the wavelength range $4100-6800$ {\\AA} with a resolution of $\\lambda/\\Delta\\lambda \\sim 42\\,000$. STELIB \\citep{leborgne03} provides spectra for 249 stars observed in the wavelength range $3200-9500$ {\\AA} with a resolution of $\\lambda/\\Delta\\lambda \\sim2000$. The observed libraries are crucial for evaluating the accuracy of synthetic spectra and can also be used directly for population synthesis and classification. STELIB has been used by \\citet{kauffmann03} with a new version of \\textsc{gissel}, and \\elodie\\ has been used to assign physical parameters to stars observed by the SDSS. However, a limitation of the observed spectral libraries is that they do not cover the full range in parameter space needed for galaxy population synthesis. Complete coverage of the parameter space is even more important for stellar spectral classification. The most successful approaches to classification have used methods from machine learning \\citep{bailerjones01a}. In these methods, the distribution of spectra in the training set has a direct impact on the accuracy of the classification assigned to new spectra.  There are several efforts to generate higher resolution Kurucz spectra that are currently in progress. \\citet{bertone02} have generated a grid of 832 spectra at $\\lambda/\\Delta\\lambda=500\\,000$ over the wavelength range $3500-7000$ {\\AA}. They intend to extend the wavelength range down to 850 {\\AA} at a resolution of $\\lambda/\\Delta\\lambda=50\\,000$. \\citet*{zwitter03} are in the process of generating a grid of Kurucz spectra at $\\lambda/\\Delta\\lambda=20\\,000$ over the wavelength range $2500-10\\,500$ {\\AA} for use in radial velocity correction work.  We have generated a larger library of 6410 spectra from the Kurucz model atmospheres. Previously these spectra were only available either at much lower resolution (20 {\\AA}) or over small wavelength ranges. Our spectra were generated from \\atlas\\ model atmospheres, using John Lester's Unix version of the \\synthe\\ spectral synthesis package (Lester 2002, private communication). We have modified this package to improve the efficiency of the code, making it possible to generate the complete range of Kurucz spectra in a reasonable time.  This paper compares our higher resolution spectra with the original 20 {\\AA} Kurucz spectra and the STELIB library of observed spectra. In Section \\ref{generate} we describe the main characteristics of the new library of spectra. In Section \\ref{comp20} we compare the spectra with the 20 {\\AA} Kurucz spectra from \\citet{kurucz93CD13}. Finally, in Section \\ref{compstellib} we compare the spectra with observed spectra from the STELIB library. We will make this library of spectra available for general use on request. ",
        "conclusions": "We have generated a grid of theoretical spectra from the Kurucz model atmospheres. Since the intended use of these spectra is in population synthesis and stellar classification, we have made several comparisons to check the validity of using the spectra for these purposes. The broadband properties of the spectra compare well with observed spectra, as do the line index measurements. The comparisons do not guarantee the accuracy of the Kurucz spectra as models of observed stellar spectra. Rather, they demonstrate that our new, high resolution Kurucz spectra are as good for use in population synthesis as the commonly used library of 20 {\\AA} spectra. The advantage of these spectra over the previously available Kurucz spectra is that they allow the modelling of spectral line features such as the Lick indices.  When using these spectra at high resolutions, it should be recognised that the spectra were generated using line lists that include `predicted' lines. This is necessary to reproduce the broadband colours of the spectra accurately. However, it does mean that many of the individual lines present at high resolutions do not have measured properties. Also, the line lists that we used (such as LOWLINES) are known to have problems with the values for specific lines. As this mostly involves weak lines, it should not present any difficulties for population synthesis at the resolution of the SDSS. More accurate properties for these lines could be obtained using alternative values, for example from the Vienna Atomic Line Database (VALD) \\citep{kupka00}.  Also, the 20 {\\AA} Kurucz flux distributions are often used with the corrections of \\cite{lejeune97,lejeune98} applied. The corrections are an attempt to calibrate the spectra by comparing the synthetic model colours with empirical stellar colours. We investigated applying the same corrections to our high resolution spectra, but found that artifacts were introduced when using this technique directly. The Lejeune corrections mainly affect the lower temperature spectra, which we have not included in this library. For most of our spectra, the corrections are negligibly small.  We intend to extend the library to the lower temperature models in which TiO lines become important and also generate spectra for the NOVER models as these are more accurate. We are in the process of building a large library for population synthesis -- a higher resolution version of that done by \\cite{lejeune97} -- which will also incorporate the NextGen models \\citep*{hauschildt99a}, \\citep{hauschildt99b}."
    },
    "0404/astro-ph0404226_arXiv.txt": {
        "abstract": "We present results of an age and metallicity gradient analysis inferred from both optical and near-infrared surface photometry.  The analysis is based on a sample of 36 nearby early-type galaxies, obtained from the Early Data Release of the Sloan Digital Sky Survey and the Two Micron All Sky Survey. Surface brightness profiles were derived in each band, and used to study the color gradients of the galaxies. Using simple stellar population models with both optical and near infrared colors, we may interpret the color gradients in term of age and metallicity gradients of galaxies. Using $g_Z \\equiv d \\log Z_{\\rm met} / d \\log R $ and $g_A = d \\log {\\rm Age} / d \\log R $ to represent the metallicity and age gradients, we found a median value of $g_Z=-0.25\\pm 0.03$ for the metallicity gradient, with a dispersion $\\sigma_{g_Z}=0.19\\pm0.02$. The corresponding values for the age gradients were $g_A=0.02\\pm 0.04$ and $\\sigma_{g_A}=0.25\\pm0.03$. These results are in good agreement with recent observational results, as well as with recent simulations that suggest both monolithic collapse and major mergers have played important roles in the formation of early-type galaxies.  Our results demonstrate the potential of using multi-waveband colors obtained from current and future optical and infrared surveys in constraining the age and metallicity gradients of early-type galaxies. ",
        "introduction": "\\label{sec intro}  The color gradients of early type galaxies have been known for quite a long time (de Vaucouleurs 1961; Boroson, Thompson \\& Shectman 1983).  Such gradients are believed to be due to the variation of the properties of the underlying stellar population, such as age and metallicity.  Theoretically, radial variations in metallicity are expected in some formation scenarios for early-type galaxies.  For example, early simulations of the monolithic collapse of a gas cloud tended to predict metallicity gradients that are too steep to match observations (e.g. Larson 1974; Carlberg 1984). Later simulations based on mergers of gas-rich galaxies predict that interactions between merging galaxies can effectively dampen their metallicity gradients (e.g. White 1980; Kobayashi 2004). Moreover, radial age gradients are also theoretically conceivable. Therefore, detailed color gradient data are essential to discriminate different formation models of early-type galaxies.  Spectroscopic indices are the most commonly used indicators of metallicity in early-type galaxies. Earlier measurements indicate the existence of systematic metallicity gradients at the level $\\Delta \\log Z / \\Delta \\log R \\approx -0.1$ to $-0.3$ (Baum, Thomsen, \\& Morgan 1986; Carollo, Danziger, and Buson 1993; Davies, Sadler, \\& Peletier 1993; Henry \\& Worthey 1999; Mehlert et al. 2000 and 2003). However, it is not easy to obtain a large sample using this technique, because of the difficulties associated with observing long-slit (two-dimensional) spectra. In addition, the rapid decrease of surface brightness with radius limits the radial range over which accurate spectroscopic measurements can be obtained. Such measurements are usually possible only out to a radius of about one or two effective radius, making it difficult to study gradients in the outer part of early-type galaxies.  On the other hand, broad band surface photometry with CCD is relatively easy to expand the measurements to several effective radius.  Furthermore, radial profiles obtained from surface photometry are an average measurement within an isophotic annulus, and so are a better representation of the average properties at a given radius.  With modern telescopes, accurate multi-waveband photometry can be obtained with little difficulty for a large number of galaxies, making it possible to study the gradients of the stellar population in a statistical way.  Because of these reasons, photometric measurements are also widely used for studying the metallicity and age gradients in early-type galaxies, although the interpretation of such measurements has to contend with the well-known age-metallicity degeneracy (e.g. Worthey 1994). Early analyses of nearby galaxies in optical bands all indicated that early-type galaxies have systematic color gradients (Boroson, Thompson, \\& Shectman 1983; Davis et. al. 1985; Cohen 1986; Franx, Illingworth, \\& Heckman 1989; Peletier et al. 1990; Michard 1999; Scodeggio 2001; Idiart, Michard, \\& de Freitas Pacheco 2002). This is supported by the more recent analysis of the surface photometry of E/S0 galaxies in the nearby rich cluster Abell 2199 by Tamura \\& Ohta (2003).  Because of the age-metallicity degeneracy, the observed color gradients are usually interpreted either as the metallicity gradients or the age gradients, assuming the gradients of the other quantity is known. Assuming the age of the stellar population to be the same over an entire galaxy, both Peletier, Valentijn, \\& Jameson (1990) and Idiart, Michard, \\& de Freitas Pacheco (2003) obtained a metallicity gradient of $\\sim -0.16$ (in terms of $\\Delta\\log Z/\\Delta\\log R$). Under the same assumption, Tamura \\& Ohta (2003) found an average metallicity gradient of $-0.3\\pm 0.1$ from a sample of 40 galaxies in Abell 2199 and 11 galaxies in Abell 2634. Using stellar population synthesis models, Saglia et al. (2000) examined the origin of the observed color gradients by comparing the 20 brightest early-type galaxies in CL0949+44, a cluster at a redshift of $\\sim 0.4$ taken from HST WF2 frames, with local galaxies. They concluded that their results are better explained in terms of passive evolution of metallicity gradients than pure age gradients. Hinkley \\& Im (2001) investigated the optical and near-infrared color gradients in the HST WFPC2 and NICMOS images for six field early-type galaxies with redshifts from 0.4 to 1.0. By comparing with stellar synthesis models, they found that 5 out of the 6 galaxies show negligible age gradients and are dominated by a metallicity gradient. Similarly, Mehlert et al. (2003) found nearly zero age gradients in 35 early-type galaxies in the Coma cluster. However, Silva \\& Elston (1994) investigated both the optical and near-infrared color gradients in eight early-type galaxies and found that all of them show both metallicity and age gradients. Therefore, whether the color gradients can be ascribed to pure metallicity gradients is still controversial.  As pointed out by de Jong (1996; see also Cardiel et al. 2003; MacArthur et al. 2004), the age-metallicity degeneracy may be partially broken by adding infrared photometry to the optical colors. Including infrared data also has another advantage, since contamination by dust is expected to be less important in infrared than in optical. Thus, accurate near infrared photometry are extremely useful in studying stellar population gradients in galaxies.  In this paper, we use the broad band photometry obtained by the Sloan Digital Sky survey (SDSS) in the optical, together with the 2 Micron All Sky Survey (2MASS) photometric data in the near-infrared, to study the color gradients for nearby early-type galaxies. The relative high image quality in both the SDSS and 2MASS surveys makes it possible to trace the spectral energy distribution (SED) to the outer part of individual galaxies. In addition, since the combined SDSS and 2MASS data give a uniform coverage of the SEDs over a large wavelength range, we may hope to use such data to derive stringent constraints on the variations of the stellar populations in early-type galaxies. As we will show, although 2MASS images are shallower and have lower resolution than the SDSS images, the S/N is sufficient for probing the color gradients to more than 2 effective radii. In this paper, we present results based on 36 early-type galaxies.  Although this sample is not larger than those used in early analyses, this is the first time where SEDs from both optical and near infrared are used to study the stellar population gradients. As we will see, the results we obtain are in good agreement with those obtained from line indices measurements, which demonstrates the strength of combining optical and near infrared photometry in constraining the stellar population gradients in early-type galaxies. In the future, much larger samples can be obtained, making this approach a very promising one for such study.  The outline of this paper is as follows. We describe our galaxy sample and related data reduction in section~\\ref{sec data}.  In sections~\\ref{sec color-gra} and ~\\ref{sec gradient of A-M} we examine the color gradients of these galaxies, and use stellar population synthesis models to obtain constraints on their age and metallicity gradients. We discuss our results in section~\\ref{sec discussion} and give a summary of our results in section~\\ref{sec summary}. Except where stated otherwise, we assume the Hubble constant to be $H_0=70 {\\rm kms^{-1}} {\\rm Mpc}^{-1}$. ",
        "conclusions": "\\label{sec discussion}  \\subsection{The Effect of Dust} \\label{subsec dust}  It must be kept in mind that the results presented above have neglected dust absorption in the host galaxies. Since dust obscuration reddens stellar light, the existence of dust changes the color of a galaxy, and can mimic a color gradient if the distribution of dust has a gradient. Although many elliptical galaxies contain significant amounts of interstellar matter (Roberts et al. 1991), most is in the form of hot X-ray gas, and only a small fraction (about $10^7$ M$_{\\odot}$) is in the form of dust (Kormendy \\& Djorgovski 1989; Forbes 1991; Goudfrooij 1994; Wise \\& Silva 1996). van Dokkum \\& Franx (1995) used HST images to study the dust properties of a sample of 64 early-type galaxies. They found that 48\\% of them show dust absorption, but that dust absorption is highly concentrated to the nuclear regions. The sizes of these absorption regions at the major axes are much smaller than 1 kpc for almost all their sample galaxies (only 3 have a dust distribution size of 2 kpc). In recent observations of 6 early type galaxies in infrared PAH bands by the Spitzer Space Telescope (Pahre et al.  2004) 3 of these 6 early-type galaxies exhibit dust features and the dust distribution have a size of about 1 to 3 kpc.  However, dust distribution in early type galaxies is still poorly understood; only simple models are available to investigate possible effect of the dust extinction on color gradients.  Here we follow Wise \\& Silva (1996) and consider three simple models for the dust distribution. The most likely case is that of a concentrated dust distribution, which is suggested by current observations. If dust distribution has a spatial density distribution of $\\rho_{d}(r)\\sim r^{-3}$ or is even more concentrated, then dust absorption is only significant in the very central region of the galaxy ($\\sim 1\\,kpc$; see Fig. 8 in Wise \\& Silva 1996). In this case, dust extinction will not have any effect on our results, since we have excluded the central region of a galaxy (within $0.4 R_{50}$, $\\sim$ 1.4 kpc on average) in our analysis. The second case is a flat density distribution of dust $\\rho_{d}(r)\\sim r^{0}$. In this case, dust absorption is almost the same over the whole body of a galaxy, and it will not produce any radial gradients in colors (or in the SEDs). A third case is a distribution between these two extreme cases. According to Wise \\& Silva's analysis, if the dust is distributed as $\\rho_{d}(r)\\sim r^{-1}$, it will cause a linear color gradient with $\\log R$, which can be confused with true metallicity and/or age gradients. To be conservative, we consider here the effect of dust extinction assuming such a distribution.  Following Charlot \\& Fall (2000), we assume a power-law form for the dust absorption curve: $\\tau_{\\lambda}=\\tau_{V}(\\lambda/5500\\AA)^{-0.7}$.  The absorption by dust can then be parameterized with a single parameter $\\tau_{V}$, which is the effective optical depth at a wavelength of $5500\\AA$. Another assumption is that $\\tau_{V}(R)$ decreases linearly with $\\log R$, and approaches {\\it zero} in the outer region (in practice, we set $\\tau_{V}(10R_{50})\\equiv 0$). This absorption model can also be derived from Wise \\& Silva's (1996) calculations (see the dotted line in their Fig.8). Thus only one additional free parameter, $\\tau_{V}(R_{50})$, is needed to describe the absorption of such a diffuse distribution of dust as a function of radius $R$.  \\begin{table} \\caption{Metallicity gradients, age gradients, and their dispersions obtained with different dust absorption levels.} \\label{tab dust} \\begin{tabular}{ccccc} \\tableline\\tableline $\\tau_{V}(R_{50})$   & $g_Z$ & $\\sigma_{g_Z}$ & $g_A$ & $\\sigma_{g_A}$ \\\\ \\tableline 0.0 & -0.25$\\pm$0.03 & 0.19$\\pm$0.03 & 0.02$\\pm$0.04 & 0.25$\\pm$0.03 \\\\ 0.1 & -0.19$\\pm$0.04 & 0.22$\\pm$0.03 & 0.02$\\pm$0.04 & 0.23$\\pm$0.03  \\\\ 0.2 & -0.13$\\pm$0.04 & 0.22$\\pm$0.03 & 0.00$\\pm$0.04 & 0.23$\\pm$0.03  \\\\ 0.4 & -0.02$\\pm$0.04 & 0.21$\\pm$0.03 & 0.00$\\pm$0.03 & 0.18$\\pm$0.03  \\\\ \\tableline \\end{tabular} \\end{table}  The SSP fitting procedure was then carried out by convolving the absorption curves at different radii with the given values of $\\tau_{V}(R_{50})$, and the effects on age and metallicity gradients recorded. We have done several tests with different values of $\\tau_{V}(R_{50})$, and the results are summarized in Table~\\ref{tab dust}. Compared with the dust-free case $\\tau_{V}(R_{50})=0$, it is clear that the existence of a diffuse dust distribution can indeed cause a radial color gradient and thereby reduce the degree of metallicity gradients. However, it does not change the average value of the age gradients. If the total amount of diffusing dust is large enough so that $\\tau_{V}(R_{50})\\sim 0.4$, the different levels of absorption at different radii can explain all the observed color gradients in our sample. >From Wise \\& Silva's calculation, about $10^{6}M_{\\odot}$ of diffuse dust is needed to produce the required $\\tau_{V}(R_{50})\\sim 0.4$.  The dust model we have used is likely too simple to describe the dust distribution in galaxies realistically. More precise description of the dust distribution in early-type galaxies may be possible with the future observations, such that from the Spitzer Space Telescope. In summary, the presence of diffuse dust distribution can mimic the color gradients we have observed, and the metallicity gradients we obtained ($g_Z\\sim -0.25$) is only robust if the amount of dust in the diffuse distribution is negligible.  \\subsection{Theoretical Implications} \\label{subsec implications}  Simulations of the formation of elliptical galaxies through monolithic collapse (e.g. Carlberg 1984) have shown that the metallicity gradients predicted by this model are  quite steep, with $g_Z$ at least as steep as -0.5. This is clearly in conflict with our average result of $g_Z\\sim -0.25$. On the other hand, recent simulations of Kobayashi (2004) showed that metallicity gradients can be significantly flattened by strong mergers. They showed that galaxies formed through major mergers have $\\Delta \\log Z / \\Delta \\log R \\sim -0.22$ and none of them have a gradient steeper then $-0.35$, while galaxies formed through minor mergers or monolithic collapse are quite similar, with metallicity gradients of approximately $-0.3$. Our average metallicity gradient is between $-0.22$ and $-0.30$, but individual gradients span a wide range from -0.6 to above zero. Since Kobayashi used a similar radius range as ours to measure gradients, we can compare our observational results directly with her simulations. To proceed, we adopt a metallicity gradient of $-0.26$ to separate galaxies into two subsamples. According to Kobayashi's results, the subsample with metallicity gradients shallower than $-0.26$ would be galaxies dominated by a major merger, while those with steeper metallicity gradients are the ones whose formation would be dominated by minor mergers or monolithic collapse. In Figure~\\ref{fig13} we plot the absolute magnitude ($M_B$) distribution for these two subsamples of galaxies. Although the average $M_B$ is similar for the two subsamples (about $M_B \\sim -20.5$), the two distributions are different. Most of the galaxies in the subsample with $g_Z<-0.26$ (i.e. the one with steeper gradient) are fainter. Since nearly half of our galaxies are S0 galaxies which have luminosities lower than elliptical galaxies, it is possible that the subsample with lower gradients is in fact dominated by S0 galaxies. Unfortunately, the sample is too small to allow us to probe the detailed dependence on morphological type. The subsample with lower gradients also includes the three brightest galaxies in the whole sample (UGC05515, UGC00797 and NGC~5865). It is interesting that UGC00797 is a cD galaxy and the other two also look like cD galaxies from their environment and radial profiles. If merger plays an important role in determining the metallicity gradient, as discussed above, this result would mean that cD-like galaxies and faint ellipticals form through many minor mergers and/or monolithic collapse. Our current sample is still too small to make a strong statement, but this question can be addressed with large samples to be constructed from the SDSS and the 2MASS.  Another interesting result is that we obtain a  quite large scatter in $g_Z$, $\\sim 0.2$. This cannot be explained by the uncertainties in measurement or fitting alone, which are approximately an order of magnitude lower. Thus, such scatter must be intrinsic. In this study we have not separated galaxies according to morphology (e.g. lenticulars versus ellipticals, normal ellipticals versus cDs) and environment, and the large scatter may be partly due to morphology and environment dependence. Our result is also consistent with the simulation of Kobayashi (2004), who obtained similar scatter among all the galaxies in her simulation.  The average value of the age gradient we found, $g_A \\sim 0$, suggests that stars over a wide range of radius in early-type galaxies (0.4 to 5.0 $R_{50}$ as probed in this paper), have approximately the same age. This is consistent with the monolithic collapse formation scenario, where the bulk of stars in a galaxy formed in a single burst, or with similar star formation history over the whole galaxy. However, the lack of a significant age gradient does not contradict with the hierarchical formation scenario, if early-type galaxies formed their stars early in their progenitors.    In this paper we have analyzed the optical and near-infrared surface photometry   of a sample of 36 nearby early-type galaxies based on broad-band images obtained from SDSS and 2MASS. The optical and near-infrared color gradients of each galaxy have been derived and modelled in terms of metallicity and age gradients in their underlying stellar populations. Our main results are summarized as follows:  \\begin{enumerate}  \\item Almost all galaxies show shallow color gradients in the five SDSS bands and the three 2MASS bands. The measured average color gradients in $U-R$, $B-R$, and $J-K_{\\rm S}$ are $-0.21$, $-0.08$, and $-0.09$ respectively, in agreement with previous results.  \\item For the first time, SEDs from optical to near-infrared wavelengths have been used to analyze the stellar population distribution of early-type galaxies. Both tests and results have proved the strength of combining SDSS and 2MASS data in the study of stellar populations.  \\item Fitting both the age and metallicity simultaneously with stellar population synthesis models, and using $g_Z = {\\rm d}\\log Z / {\\rm d}\\log R $ and $g_A = {\\rm d} \\log {\\rm Age} /{\\rm d}\\log R$ to represent the metallicity and age gradients, a median value of $g_Z=-0.25\\pm 0.03$ was found for the metallicity gradient, with a dispersion of $\\sigma_{g_Z}=0.19\\pm0.02$. Corresponding values for the age gradient are $g_A=0.02\\pm 0.04$ and $\\sigma_{g_A}=0.25\\pm0.03$. These results are in good agreement with other observational results and with the expectations of current theories of galaxy formation.  \\item A diffuse distribution of dust [$\\rho_{d}(r)\\sim r^{-1}$] could produce a linear color gradient with $\\log R$, but it will not change the result that the average age gradient is $\\sim 0$.  If any diffusely distributed dust exists, it would decrease metallicity gradients, which would further support the hypothesis that major mergers play an important role in the formation of the early-type galaxies.  \\end{enumerate}  Given the amount of new data now available from the 2MASS and  the current SDSS (DR2), and data soon to be available from the SDSS and the Spitzer Space Telescope in the near future, it will be possible to compile a sample of several hundred large early-type galaxies to carry out the same analysis as presented in this paper. Our results presented here suggest that with such a sample, it will be possible to study the age-metallicity gradients of early-type galaxies in unprecedented detail."
    },
    "0404/hep-ph0404239_arXiv.txt": {
        "abstract": "The current status of neutrino cosmology is reviewed, from the question of neutrino decoupling and the presence of sterile neutrinos to the effects of neutrinos on the cosmic microwave background and large scale structure. Particular emphasis is put on cosmological neutrino mass measurements. ",
        "introduction": "%  Next to photons neutrinos are the most abundant particles in the universe. This means they have a profound impact on many different aspects of cosmology, from the question of leptogenesis in the very early universe, over big bang nucleosynthesis, to late time structure formation. In the present review I focus mainly on late-time aspects of neutrino cosmology, and particularly on issues relevant to cosmological bounds on the neutrino mass.  The absolute value of neutrino masses are very difficult to measure experimentally. On the other hand, mass differences between neutrino mass eigenstates, $(m_1,m_2,m_3)$, can be measured in neutrino oscillation experiments.  The combination of all currently available data suggests two important mass differences in the neutrino mass hierarchy. The solar mass difference of $\\delta m_{12}^2 \\simeq 7 \\times 10^{-5}$ eV$^2$ and the atmospheric mass difference $\\delta m_{23}^2 \\simeq 2.6 \\times 10^{-3}$ eV$^2$ \\cite{Maltoni:2003da,Aliani:2003ns,deHolanda:2003nj}.  In the simplest case where neutrino masses are hierarchical these results suggest that $m_1 \\sim 0$, $m_2 \\sim \\delta m_{\\rm solar}$, and $m_3 \\sim \\delta m_{\\rm atmospheric}$. If the hierarchy is inverted \\cite{Kostelecky:1993dm,Fuller:1995tz,Caldwell:1995vi,Bilenky:1996cb,King:2000ce,He:2002rv} one instead finds $m_3 \\sim 0$, $m_2 \\sim \\delta m_{\\rm atmospheric}$, and $m_1 \\sim \\delta m_{\\rm atmospheric}$. However, it is also possible that neutrino masses are degenerate \\cite{Ioannisian:1994nx,Bamert:vc,Mohapatra:1994bg,Minakata:1996vs,% Vissani:1997pa,Minakata:1997ja,Ellis:1999my,Casas:1999tp,Casas:1999ac,% Ma:1999xq,Adhikari:2000as}, $m_1 \\sim m_2 \\sim m_3 \\gg \\delta m_{\\rm atmospheric}$, in which case oscillation experiments are not useful for determining the absolute mass scale.  Experiments which rely on kinematical effects of the neutrino mass offer the strongest probe of this overall mass scale. Tritium decay measurements have been able to put an upper limit on the electron neutrino mass of 2.3 eV (95\\% conf.) \\cite{kraus}. However, cosmology at present yields an much stronger limit which is also based on the kinematics of neutrino mass.  Very interestingly there is also a claim of direct detection of neutrinoless double beta decay in the Heidelberg-Moscow experiment \\cite{Klapdor-Kleingrothaus:2001ke,Klapdor-Kleingrothaus:2004wj}, corresponding to an effective neutrino mass in the $0.1-0.9$ eV range. If this result is confirmed then it shows that neutrino masses are almost degenerate and well within reach of cosmological detection in the near future.  Another important question which can be answered by cosmological observations is how large the total neutrino energy density is. Apart from the standard model prediction of three light neutrinos, such energy density can be either in the form of additional, sterile neutrino degrees of freedom, or a non-zero neutrino chemical potential.  The paper is divided into sections in the following way: In section 2 I review the present cosmological data which can be used for analysis of neutrino physics. In section 3 I discuss neutrino physics around the epoch of neutrino decoupling at a temperature of roughly 1 MeV, including the relation between neutrinos and Big Bang nucleosynthesis. Section 4 discusses neutrinos as dark matter particles, including mass constraints on light neutrinos, and sterile neutrino dark matter. Section 5 contains a relatively short review of neutrino physics in the very early universe from the perspective of leptogenesis. Finally, section 6 contains a discussion. ",
        "conclusions": "%  In the present paper I have discussed how cosmological observations can be used for probing fundamental properties of neutrinos which are not easily accessible in lab experiments. Particularly the measurement of absolute neutrino masses from CMB and large scale structure data has received significant attention over the past few years. In Table \\ref{table:summary} I summarize neutrino mass bounds from cosmological observations and other astrophysical and experimental bounds.       \\begin{table} \\caption{Summary table of cosmological neutrino mass limits. For completeness bounds from other sources, astrophysical and experimental, are also listed.} \\begin{center} \\begin{tabular}{@{}lll} \\hline  Method & Bound on $\\sum m_\\nu$ & Data used \\\\ \\hline  $\\Omega_\\nu h^2 \\lesssim 0.15$  &   14 eV     &  $\\Omega_\\nu < \\Omega_m$ \\\\   CMB and LSS     &   0.7-1 eV    &   WMAP, 2dF, SDSS \\\\   \\hline  SN1987A         &   $m_{\\nu_e} < 5-20$ eV  $(\\bar{\\nu}_e)$ & SN1987A cooling curve \\cite{Kernan:1994kt,Loredo:2001rx}\\\\  $\\beta$-decay   &   $m_{\\nu_e} < 2.2$ eV $(\\nu_e)$  & Mainz experiment \\cite{kraus} \\\\   $0\\nu 2\\beta$-decay    & $m_{\\nu,{\\rm eff}} < 0.35$ eV & Heidelberg-Moscow \\cite{Klapdor-Kleingrothaus:2000sn} \\\\  & 0.1 eV $< m_{\\nu,{\\rm eff}} <$ 0.9 eV & Heidelberg-Moscow \\cite{Klapdor-Kleingrothaus:2001ke,Klapdor-Kleingrothaus:2004wj} \\\\   \\hline \\end{tabular} \\end{center} \\label{table:summary} \\end{table}     Another cornerstone of neutrino cosmology is the measurement of the total energy density in non-electromagnetically interacting particles. For many years Big Bang nucleosynthesis was the only probe of relativistic energy density, but with the advent of precision CMB and LSS data it has been possible to complement the BBN measurement. At present the cosmic neutrino background is seen in both BBN, CMB and LSS data at high significance.  Finally, cosmology can also be used to probe the possibility of neutrino warm dark matter, which could be produced by active-sterile neutrino oscillations.  In the coming years the steady stream of new observational data will continue, and the cosmological bounds on neutrino will improve accordingly. For instance, it has been estimated that with data from the upcoming Planck satellite it could be possible to measure neutrino masses as low as 0.1 eV.  Certainly neutrino cosmology will continue to be a prospering field of research for the foreseeable future."
    },
    "0404/astro-ph0404406_arXiv.txt": {
        "abstract": " ",
        "introduction": "The OH and H$_2$O maser clusters in the W3(OH)/W3(H$_2$O) region are separated by about 5$^{\\prime\\prime}$ (0.05 pc at a distance of 2.2 kpc) and are associated with two young stellar objects at different stages of evolution. The OH and class II methanol masers are seen toward W3(OH) which contains an expanding ultra-compact HII region (UCHII) surrounding a young O7 star (Kawamura$\\&$Masson, 1998). To its east is W3(H$_2$O), a region containing H$_2$O masers and a young stellar object known as the TW object. Both objects show outflow activity (see, e.g., descriptions of the W3(OH) champagne outflow in Keto etal.,1995 and the W3(H$_2$O) jet in Wilner etal., 1999). Situated in close proximity and embedded in the same molecular core (see, e.g., Wilson etal. 1993) these objects have rather similar velocities.  Here we report observations of molecular lines with BIMA which reveal the distibution of shock tracing molecules, in particular methanol and water, molecules which also produce maser emission. Additional observations at Onsala are used to reveal overall characteristics of the outflows. New detections of maser lines are used to construct model of the class II methanol maser formation region. ",
        "conclusions": ""
    },
    "0404/astro-ph0404589_arXiv.txt": {
        "abstract": "We investigate the subhalo populations of dark matter haloes in the concordance $\\Lambda$CDM cosmology. We use a large cosmological simulation and a variety of high resolution resimulations of individual cluster and galaxy haloes to study the systematics of subhalo populations over ranges of 1000 in halo mass and 1000 in the ratio of subhalo to parent halo mass. The subhalo populations of different haloes are not scaled copies of each other, but vary systematically with halo properties.  On average, the amount of substructure increases with halo mass. At fixed mass, it decreases with halo concentration and with halo formation redshift. These trends are comparable in size to the scatter in subhalo abundance between similar haloes. Averaged over all haloes of given mass, the abundance of low mass subhaloes per unit parent halo mass is independendent of parent mass.  It is very similar to the abundance per unit mass of low mass haloes in the universe as a whole, once differing boundary definitions for subhaloes and haloes are accounted for. The radial distribution of subhaloes within their parent haloes is substantially less centrally concentrated than that of the dark matter. It varies at most weakly with the mass (or concentration) of the parent halo and not at all with subhalo mass. It does depend on the criteria used to define the subhalo population considered. About $90$ per cent of present-day subhaloes were accreted after $z=1$ and about $70$ per cent after $z=0.5$. Only about $8$ per cent of the total mass of all haloes accreted at $z=1$ survives as bound subhaloes at $z=0$. For haloes accreted at $z=2$, the survival mass fraction is just 2 per cent. Subhaloes seen near the centre of their parent typically were accreted earlier and retain less of their original mass than those seen near the edge. These strong systematics mean that comparison with galaxies in real clusters is only possible if the formation of the luminous component is modelled appropriately. ",
        "introduction": "According to the standard CDM scenario, structure in our Universe formed hierarchically. Small-scale fluctuations were the first to collapse as virialised objects. These then merged to form larger systems. The inner regions of early virialised objects are very compact and often survive accretion onto a larger system to become self-bound subhaloes of their host. Since galaxies form by the condensation of gas at the centres of early haloes, most cluster galaxies may well be associated with subhaloes in their host cluster. Only in recent years have numerical techniques and computer capabilities advanced to the point where it is possible to study in detail the properties of such subhaloes (Moore et al. 1998, 1999; Tormen, Diaferio \\& Syer 1998; Klypin et al. 1999a,b; Ghigna et al. 1998, 2000; Springel et al. 2001; Stoehr et al. 2002, 2003). These studies indicate that the `overmerging' problem in early simulations, i.e. the failure to resolve subhaloes corresponding to galaxies in cosmological simulations of cluster haloes, was in part a result of insufficient mass and force resolution.  Using high resolution resimulations of individual cluster or galaxy haloes, it is possible to study the properties of subhaloes in detail. Recent papers by De Lucia et al. (2004), Diemand et al. (2004) and Gill, Knebe \\& Gibson (2004) discuss many aspects of this topic and present results compatible with but complementary to those presented below. Most studies to date have been limited because their analysis has been performed on a small number of individual haloes. Since halo-to-halo variations are large, this may prevent the derivation of statistically significant results. In addition, all studies are still affected at some level by numerical resolution. The available tests show that the subhaloes seen in a particular object are reproduced moderately well in mass, but not in position or velocity, when the same object is resimulated multiple times with varying resolution (Ghigna et al. 2000; Springel et al. 2001; Stoehr et al 2002, 2003). This is a result of the well known divergence of neighboring trajectories in nonlinear dynamical systems.  In this paper, we carry out a systematic study of the properties of subhaloes in the halo population of a single, large-scale cosmological simulation, and we complement this by analysing a multi-resolution set of resimulations of a single `Milky Way' halo, together with a set of high-resolution resimulations of eight different rich clusters. These resimulations allow us to investigate how numerical resolution and halo-to-halo variation affect the conclusions from our cosmological simulation.  We do not, however, carry out a full study of the numerical requirements for fully converged numerical results for the properties of subhaloes.  Previous studies of subhaloes within haloes of different scale have emphasised similarities -- to a large extent the internal structure of a `Milky Way' halo looks like a scaled version of that of a rich cluster halo (Moore et al. 1999; Helmi \\& White 2001; Stoehr et al. 2003; De Lucia et al. 2004; Desai et al. 2004). We show below that this scaling is not exact, and that a better model assumes the mass distribution of low-mass subhaloes to be the same as in the Universe as a whole, once the differing definitions of an object's boundary are accounted for. We show that galaxy haloes have fewer high-mass subhaloes than rich clusters because of their earlier formation times. Indeed, even among haloes of given mass, the number of massive subhaloes correlates well with formation time, as reflected in the halo's central concentration.  The emphasis of earlier high resolution work on solving the `overmerging problem' has given rise to the impression that the subhaloes are typically objects which formed at very early times. We demonstrate below that this is not the case. Even at low subhalo masses, most subhaloes were accreted onto the main halo at low redshift, in most cases well below $z=1$. This is important when considering the formation paths of present-day cluster galaxies.  Our paper is organized as follows. We introduce our various simulation sets in Section 2. In Section 3, we compare the halo mass abundance function measured from our cosmological simulation with theoretical predictions and with earlier numerical data. In Section 4, we investigate the subhalo population as a function of halo mass and of redshift. The spatial distribution of subhaloes within haloes is also discussed in Section 4. In Section 5 we investigate the infall and mass-loss histories of present-day subhaloes, as well as the fate of objects that are accreted onto bigger clusters at early times. We discuss our results and set out our conclusions in Section 6. ",
        "conclusions": "We have used a single, large-scale cosmological simulation together with two sets of resimulations of the formation of individual cluster and galaxy haloes to carry out a systematic study of the properties of dark halo substructure in the concordance $\\Lambda$CDM universe. In agreement with the earlier work of Jenkins et al. (2001), Reed et al. (2003) and Yahagi et al. (2004) we find the abundance of haloes (defined using a friends-of-friends group finder with linking length $b=0.2$) to be well described by the Sheth \\& Tormen (1999) mass function down to masses of a few times $10^{10}{\\rm M_\\odot}$ and out to a redshift of 5. Our main results for the subhalo populations within these haloes can be summarized as follows:  \\begin{description}  \\item[(1)] The subhalo populations of different haloes are not simply scaled copies of each other, but vary systematically with global halo properties. On average, massive haloes contain more subhaloes above any given fraction of parent mass than do lower mass haloes, and these subhaloes contain a larger fraction of the parent's mass. At given halo mass, subhaloes are more abundant in haloes which are less concentrated, or formed more recently.  \\item[(2)] There is considerable scatter in the abundance of subhaloes between haloes of similar mass, concentration or formation time. This presumably reflects differences in the details of halo assembly.  \\item[(3)] For subhalo masses well below that of the parent halo the mean subhalo abundance {\\it per unit parent mass} is independent of the actual mass of the parent. It is very similar to the abundance of haloes per unit mass in the universe as a whole, once a correction is made for the differing bounding density within which the masses of haloes and subhaloes are defined.  \\item[(4)] Normalised in this way to total parent halo mass, the mean abundance of subhaloes as a function of maximum circular velocity is also quite similar to the abundance per unit mass of haloes as a function of $V_{\\rm max}$. For subhaloes the abundance per unit mass is about a factor of two lower at given $V_{\\rm max}$ than for haloes. Equivalently, the $V_{\\rm max}$ values of subhaloes at given abundance per unit mass are about 25 per cent lower than those for haloes.  \\item[(5)] In agreement with previous studies, we find the the radial distribution of subhaloes within their parent haloes to be much less concentrated than that of the dark matter. We find no significant dependence of this radial profile on the mass of the subhaloes and only a very weak dependence on the mass (or concentration) of the parent halo. To a good approximation the radial distribution of subhaloes appears `universal' and we give a fitting formula for it in equation~\\ref{pro}.  \\item[(6)] The subhalo number density profile does depend on how the population is defined. Subhalo populations defined above a minimum circular velocity limit are significantly more concentrated than those defined above a minimum mass limit.  \\item[(7)] Most subhaloes in present-day haloes fell into their parent systems very recently. Only about 10 per cent of them were accreted earlier than $z=1$ and 70 per cent were accreted at $z<0.5$. These fractions depend very little on the mass of the subhaloes or on that of their parents  \\item[(8)] The rate at which tidal effects reduce the mass of subhaloes is  not strongly dependent on the mass of the accreted object or on that of the halo it falls into. About 92 per cent of the total mass of haloes accreted at $z=1$ is removed to become part of the `smooth' halo component by $z=0$. For haloes which fall in at $z=2$ this fraction is about 98 per cent. Note that the highest mass accreted objects merge into the central regions more quickly because of dynamical friction effects.   \\item[(9)] Subhaloes seen near the centre of their parent haloes typically fell in earlier and retain a smaller fraction of their original mass than subhaloes seen near the edge. Thus inner subhaloes may be expected to host brighter galaxies than outer subhaloes of similar mass (see Springel et al. 2001). \\end{description}  These properties suggest a relatively simple picture for the evolution of subhalo populations. A substantial fraction of the mass of most haloes has been added at relatively recent redshifts, and this mass is accreted in clumpy form with a halo mass distribution similar to that of the Universe as a whole. Since tidal stripping rapidly reduces the mass of subhaloes, the population at any given mass is dominated by objects which fell in recently and so had lower mass (and thus more abundant) progenitors. The orbits of recently accreted objects spend most of their time in the outer halo, so that subhaloes of given mass are substantially less centrally concentrated than the dark matter as a whole. Subhaloes which are seen near halo centre have shorter period orbits and so must have fallen in earlier. They thus retain a relatively small fraction of their initial mass.  Comparison of these subhalo properties with observation is far from simple. The recent accretion of most subhaloes means that the galaxies at their centres were almost fully formed by the time they became part of their current host. We might therefore expect their observable properties to be more closely related to the mass of their progenitor haloes and to their accretion redshifts than to the current masses of their subhaloes. Explicit tracking of galaxy formation during the assembly of cluster haloes shows that these differences can be large. For example, both Diaferio et al. (2001) and Springel et al. (2001) find radial number density profiles for magnitude limited samples of galaxies which are similar both to the underlying dark matter profiles and to the observed profiles of real clusters, but which are very different from the number density profiles for mass limited subhalo samples. Similar differences are to be expected between the velocity biases of galaxies and subhaloes. Models for the stellar content of subhaloes which are based purely on their current mass and internal structure are very unlikely to be successful. The past history of subhaloes must be included to get realistic results, as must galaxies associated with apparently disrupted subhaloes. We investigate these issues further in a companion paper (Gao et al. 2004b);"
    },
    "0404/astro-ph0404295_arXiv.txt": {
        "abstract": "Using a combination of observations involving the VLA, MERLIN and global VLBI networks we have made a detailed study of the radio continuum and the neutral hydrogen (H{\\sc i}) kinematics and distribution within the central kiloparsec of the radio galaxy 3C\\,293. These observations trace the complex jet structure and identify the position of the steeply inverted radio core at 1.3\\,GHz.  Strong H{\\sc i} absorption is detected against the majority of the inner kiloparsec of 3C\\,293. This absorption is separated into two dynamically different and spatially resolved systems. Against the eastern part of the inner radio jet narrow H{\\sc i} absorption is detected and shown to have higher optical depths in areas co-spatial with a central dust lane. Additionally, this narrow line is shown to follow a velocity gradient of $\\sim$50\\,km\\,s$^{-1}$\\,arcsec$^{-1}$, consistent with the velocity gradient observed in optical spectroscopy of ionised gas. We conclude that the narrow H{\\sc i} absorption, dust and ionised gas are physically associated and situated several kiloparsecs from the centre of the host galaxy. Against the western jet emission and core component, broad and complex H{\\sc i} absorption is detected. This broad and complex absorption structure is discussed in terms of two possible interpretations for the gas kinematics observed. We explore the possibility that these broad, double absorption spectra are the result of two gas layers at different velocities and distances along these lines of sight. A second plausible explanation for this absorbing structure is that the H{\\sc i} is situated in rotation about the core of this radio galaxy with some velocity dispersion resulting from in-fall and outflow of gas from the core region. If the latter explanation were correct, then the mass enclosed by the rotating disk would be at least 1.7$\\times$10$^9$ solar masses within a radius of 400\\,pc. ",
        "introduction": "Nuclear activity in galaxies manifests itself in a variety of forms from nearby low luminosity active galactic nuclei (AGN), such as Seyferts and LINERS, to powerful distant quasars and radio galaxies.  In these sources, the nuclear activity is responsible for radiation detected across the entire electromagnetic spectrum.  In the radio loud active galaxies, such as quasars and radio galaxies, the radio emission demonstrates the influence of the AGN, for example via the formation of powerful jets (Fanaroff \\& Riley 1974).  Additionally ample evidence is also available from other wavelength ranges ({\\it e.g} optical) for the interaction of the nuclear activity with the surrounding galactic interstellar medium (ISM) such as via the detection of outflows from nuclear regions of some Seyfert galaxies ({\\it e.g.} NGC\\,3079, Cecil et al. 2001). The commonly accepted standard model for nuclear activity asserts that the AGN is fuelled by the release of gravitational potential energy as galactic material is accreted onto a central super-massive black hole. As such it is incumbent upon investigators to study not only the effects of this nuclear activity ({\\it e.g.} jets) but also the physical and kinematic environment that surrounds the AGN, since this provides a method by which we can study how gas, dust and stars act as fuel for the activity we observe as well as how this activity impacts the surrounding ISM.  At the present time most high angular resolution studies of powerful active galaxies have concentrated upon investigating the consequences of the activity ({\\it e.g.} the synchrotron emission such as radio jets), rather than the cause of the activity ({\\it e.g.} cold neutral and molecular gas that fuels the nuclear activity). This has primarily been a direct consequence of observational constraints resulting from the relatively small collecting areas of the current generation of mm-wavelength aperture synthesis instruments and the surface brightness sensitivity of decimetre wavelength interferometers precluding observations of cold thermal gas emission at angular resolutions of $\\ltsim$1\\,arcsec.  However using current radio aperture synthesis techniques it is possible to observe cold gas, via decimetre transitions such as H{\\sc i}, OH and H$_2$CO, in absorption against the bright background radio continuum of some galaxies, on sub-arsecond angular scales. \\begin{figure} \\begin{center} \\setlength{\\unitlength}{1mm} \\begin{picture}(80,75) \\put(0,0){\\special{psfile= fig1sm.ps hoffset=-20 voffset=230 hscale=43 vscale=43 angle=270}} \\end{picture} \\caption{The {\\it u-v} coverage of the combined VLBI and MERLIN data-sets.} \\label{fig1a} \\end{center}  \\end{figure} \\begin{center} \\begin{table} \\caption[]{Summary of Observations. All three of these observations were made in spectral line mode with the bandwidth centred at 1.359\\,GHz.} \\begin{tabular}{lccc} \\hline Interferometer&Date&BW&{\\it uv} Range\\\\ &&(MHz)&(k$\\lambda$)\\\\ \\hline VLA-A$+$PT&15/12/2000&3.125&3.1-235\\\\ MERLIN&08/04/1998&8&11-989\\\\ VLBI&18/11/1999&8&162-4850\\\\ \\hline \\end{tabular} \\label{tab1} \\end{table} \\end{center}  3C\\,293 is a nearby radio galaxy associated with the peculiar elliptical galaxy {\\sc vv}5-33-12. On scales of several tens of kiloparsecs, the radio jet structure of 3C\\,293 has been well studied by Bridle, Fomalont \\& Cornwell (1981) and van Breugel et al. (1984) and resembles a moderately large two-sided FR-II radio galaxy.  However, 3C\\,293 is peculiar in that an unusually high proportion of the galaxy's radio power is emitted from a steep-spectrum extended core component.  This core region when observed at higher angular resolution, is found to be a composite of several radio components forming a kiloparsec scale east-west orientated jet (Bridle et al. 1981; Akujor et al. 1996; Beswick, Pedlar \\& Holloway 2002). At other wavelengths, 3C\\,293 and its associated galaxy display several distinctive characteristics. VV5-33-12 has a closely interacting small companion galaxy situated $\\sim$37\\arcsec\\,\\,($\\sim$30\\,kpc) toward the south-west (Heckman et al. 1985; Evans et al. 1999) and the central region of the galaxy is criss-crossed by several filamentary dust lanes aligned in an approximately north-south direction (van Breugel et al. 1981; Martel el. 1999; Allen et al. 2002). Additionally {\\it Hubble Space Telescope} ({\\it HST}) observations have detected an optical/IR jet within the central kiloparsec (Leahy, Sparks \\& Jackson 1999) partially obscured from previous observations by the nuclear dust lanes. Arcsecond resolution observations of molecular gas in 3C\\,293 by Evans et al. (1999) have revealed large concentrations of CO(1$\\rightarrow$0) detected in both emission and absorption within the central few kiloparsecs. The CO emission is primarily distributed in an asymmetric disk rotating about an unresolved continuum component that Evans et al. conclude is the AGN.  \\begin{figure} \\setlength{\\unitlength}{1mm} \\begin{picture}(80,77) \\put(0,0){\\special{psfile= fig2sm.ps hoffset=-20 voffset=230 hscale=43 vscale=43 angle=270}} \\end{picture} \\caption{VLA B configuration image at 1.35\\,GHz of the large scale structure of the radio galaxy 3C\\,293 along with inset MERLIN image of the central $\\sim$3\\,kpc. The VLA image has been contoured at $\\sqrt2$ times 0.5\\,mJy\\,beam$^{-1}$ with a peak flux of 3.44\\,Jy. The lowest contour of the inset MERLIN image is 5\\,mJy\\,beam$^{-1}$ and follows the same multiplying factors as the VLA image. The peak of the MERLIN image is 1.29\\,Jy\\,beam$^{-1}$. The MERLIN radio continuum image can be seen in more detail in Fig.\\,2.} \\label{fig1} \\end{figure}  Broad neutral hydrogen absorption was first detected against 3C\\,293 by Baan \\& Haschick (1981) using the Arecibo telescope and has been studied in great detail with ever improving sensitivities and angular resolutions using a variety of radio interferometers over the last two decades (Shostak et al. 1983; Haschick \\& Baan 1985; Beswick et al. 2002; Morganti et al. 2003). 3C\\,293 has proved to be worthy of these numerous studies because of its complex and exceptionally broad H{\\sc i} absorption structure which has been observed against the extended nuclear radio continuum source. In order to fully sample the wide range of physical scales of the nuclear radio continuum of 3C\\,293 from several arcseconds to angular resolutions of a few tens of mas we have combined global Very Long Baseline Interferometry (VLBI) observations with previously published Multi-Element Radio Linked Interferometric Network (MERLIN) data (Beswick et al. 2002) and Very Large Array (VLA: A configuration including the VLBA Pie Town antenna) observations to provide a wide range of {\\it u-v} spacings (see Tab.\\,1 and Fig.\\,1). This combined study allows this extended radio source to be imaged with high fidelity at a variety of angular resolutions.  \\begin{figure*} \\setlength{\\unitlength}{1mm} \\begin{picture}(80,105) \\put(0,0){\\special{psfile= fig3sm.ps hoffset=-155 voffset=320 hscale=95 vscale=95 angle=270}} \\end{picture} \\vskip 16pc \\caption{Sub-arcsecond continuum structure of the inner few kiloparsecs of 3C\\,293. The top contour map shows the 1.359 GHz radio continuum structure observed with MERLIN at a resolution of 0\\farcs23$\\times$0\\farcs20. Contour levels of the MERLIN image are the same as for Fig.\\,1. The lower panel shows the global VLBI, MERLIN and VLA+PT contoured image of the inner jet of 3C\\,293 with angular resolution of 30\\,mas. This map is contoured at multiples of $\\sqrt2$ times 1.3\\,mJy\\,beam$^{-1}$. The peak flux of the 30\\,mas image is 38.23\\,mJy. Labels A, B, C and D at the top of the MERLIN image follow the convention of labeling continuum components in this source used by Bridle et al. (1981) and Beswick et al. (2002).} \\label{fig2} \\end{figure*}  This paper is split into four additional sections. The first of these will describe the observations presented and the data processing that has been applied to them. This will be followed by the presentation of the observational results and a more detailed discussion of their implications. The discussion will initially concentrate upon the radio continuum structure of 3C\\,293 from arcmin to milliarcsec scales, comparing these new observations with previously published data-sets, followed by a detailed discussion of these new mas resolution H{\\sc i} absorption observations against the central kiloparsec of 3C\\,293. The final section of this paper will outline the key conclusions of this work and place them in context of other radio galaxies and their environments.  Throughout this paper we assume H$_0=$75\\,km\\,s$^{-1}$\\,Mpc$^{-1}$. At a redshift of z=0.045 this implies a distance for 3C\\,293 of 180\\,Mpc so an angular size of 1\\,mas corresponds to 0.815\\,pc.  \\begin{center} \\begin{table*} \\caption[]{Spectral indices between 1.359 and 4.546 GHz\\footnotemark[1] for components in the inner jet. Flux densities for both frequencies have been obtained from matched angular resolution (50 mas) images.  Positions of components have been derived from the 1.359 GHz image presented pictorially in Fig.\\,2, bracketed labels equivalent to the position labels of the spectra in Fig.\\,3 \\& 4 and Tab.\\,3. The flux densities quoted are for the peaks for each radio continuum component at both frequencies in addition to the integrated flux density of each component. Spectral indices have been calculated using S$_\\nu\\propto\\nu^{-\\alpha}$.} \\begin{tabular}{l||c|c|c|c|c|c|c} \\hline Component&RA (J2000) &Dec (J2000)&S$_{\\rm 1.359GHz}$&S$_{\\rm 1.359GHz}$&S$_{\\rm 4.546GHz}$&S$_{\\rm 4.546GHz}$&Spectral\\\\ &&&Peak&Total&Peak&Total&Index\\\\ &13$^{\\rm h}$\\,52$^{\\rm m}$&31\\degr\\,26\\arcmin&(mJy\\,beam$^{-1}$)&(mJy)&(mJy\\,beam$^{-1}$)&(mJy)&$\\alpha{^{1.3}_{4.5}}$\\\\ \\hline E1(2)&17$^{\\rm s}\\!\\!$.913&46\\farcs48&90.48&564.13&55.00&256.25&0.65\\\\ E2(5)&17.895&46.34&53.81&427.06&28.50&242.80&0.47\\\\ E3(6)&17.873&46.36&17.44&134.10&10.42&67.69&0.57\\\\ Core(7)&17.800&46.48&16.38&27.10&20.87&23.55&0.11\\\\ W1(8)&17.791&46.56&38.52&460.69&9.60&133.38&1.03\\\\ W2(11)&17.773&46.59&59.13&484.60&25.84&214.67&0.67\\\\ \\hline \\end{tabular} \\label{tab2} \\end{table*} \\end{center} ",
        "conclusions": "We have used observations made using the VLA including Pie Town, MERLIN and global VLBI to study the 1.3\\,GHz radio continuum structure of the central kiloparsec of the peculiar radio galaxy 3C\\,293 and to investigate the kinematics and the distribution of H{\\sc i} via absorption against this radio continuum.  We confirm the component identified originally by Akujor et al. (1996) as the most probable site of the central engine and place limits upon its size of $\\ltsim$17\\,pc. Using both MERLIN 4.5\\,GHz data (Akujor et al. 1996) and new high resolution 1.3\\,GHz observations presented here we determine a radio spectral index of $\\alpha^{1.3}_{4.5}=$0.11 for the core component. We discuss in detail the differences in the value of this spectral index compared to that derived for the same component at higher radio frequencies ($\\alpha^{15}_{22}\\approx-1$; Akujor et al. 1996). It is concluded the flat spectral index derived in this study may be affected by a significant amount of unresolved radio emission from the jet close to the AGN and that this has been incorporated in our 1.3\\,GHz flux density measurement of the component resulting in the observed flat spectrum. In addition to the core we have mapped the radio continuum emission of the jet in 3C\\,293 across a variety of angular scales. From this multi-scale approach it is apparent that the trajectory of the radio jet emission in 3C\\,293 has changed significantly over the source's lifetime.  In this region, observational data suggest the presence of an intrinsically symmetric jet with a highly relativistic spine, surrounded by a low velocity shear layer because of the jet interaction with the dense ISM. The jet orientation with respect to the line of sight is $\\sim$50$^\\circ$ with the eastern jet approaching us.  Extensive H{\\sc i} absorption has been detected against both the eastern and western jet components within the central kiloparsec of 3C\\,293, consistent with lower angular resolution studies by Haschick \\& Baan (1985) and Beswick et al. (2002). As was previously known, the structure of the H{\\sc i} absorption against the eastern jet components primarily consists of strong and narrow features with a small velocity gradient whereas the absorption against the western and core components are much broader.  The narrow H{\\sc i} absorption detected against the eastern radio jet traces a small velocity gradient of $\\sim50$\\,km\\,s$^{-1}$\\,arcsec$^{-1}$, consistent with the velocity gradient observed in ionised gas by van Breugel et al. (1984). Additionally we have re-confirmed (following Beswick et al. 2002) that this narrow H{\\sc i} absorption is co-spatial with the location of dust lanes observed by the {\\it HST}. From the association of these three components we conclude that they are all most likely situated $\\sim$8\\,kpc from the central part of the galaxy and are all probably undergoing galactic rotation.  Against the western radio jet and core complex H{\\sc i} absorption is also detected. This absorption is discussed in terms of either tracing two gas structures at undetermined distances along the line of sight to the jet or a steep velocity gradient which may be interpreted as neutral gas in rotation about the core.  If this is interpreted as rotation by a gas disk, it would imply an enclosed mass of  at least 1.7$\\times$10$^9$  solar masses within a radius of 400\\,pc of the core."
    },
    "0404/gr-qc0404076_arXiv.txt": {
        "abstract": "% The nonsymmetric gravitational theory predicts an acceleration law that modifies the Newtonian law of attraction between particles. For weak fields a fit to the flat rotation curves of galaxies is obtained in terms of the mass (mass-to-light ratio $M/L$) of galaxies. The fits assume that the galaxies are not dominated by exotic dark matter. The equations of motion for test particles reduce for weak gravitational fields to the GR equations of motion and the predictions for the solar system and the binary pulsar PSR 1913+16 agree with the observations. The gravitational lensing of clusters of galaxies can be explained without exotic dark matter. ",
        "introduction": "A gravitational theory explanation of the acceleration of the expansion of the universe~\\cite{Perlmutter,Riess,Spergel} and the observed flat rotation curves of galaxies was proposed~\\cite{Moffat}, based on the nonsymmetric gravitational theory (NGT)~\\cite{Moffat2,Moffat3,Moffat4}. Since no dark matter has been detected so far, it seems imperative to seek a possible modified gravitational theory that could explain the now large amount of data on galaxy rotation curves. The same holds true for the need to explain the acceleration of the expansion of the universe without having to invoke a cosmological constant, because of the serious problems related to this constant~\\cite{Weinberg}.  In the following, we summarize the derivation of the motion of test particles in NGT. We consider the derivation of test particle motion from the NGT conservation laws. The motion of a particle in a static, spherically symmetric gravitational field is derived, yielding the modified Newtonian law of motion for weak gravitational fields. Two parameters $\\sqrt{M_0}$ and $r_0$ occur in the generalized Newtonian acceleration law. The parameter $\\sqrt{M_0}$ is modelled for a bound system by a dependence on the mean orbital radius of a test particle, and the range parameter $r_0$ is determined for galaxies and clusters of galaxies from the acceleration $cH_0$ where $H_0$ is the measured Hubble constant. A fit to both low surface brightness and high surface brightness galaxies is achieved in terms of the total galaxy mass $M$ (or $M/L$) without exotic dark matter. A satisfactory fit is achieved to the rotational velocity data generic to the elliptical galaxy NGC 3379. Fits to the data of the two spheroidal dwarf galaxies Fornax and Draco and the globular cluster $\\omega$ Cenauri are also obtained. The predicted light bending and lensing can lead to agreement with galaxy cluster lensing observations.  The modelled values of the parameter $\\sqrt{M_0}$ for the solar system and Earth, lead to agreement with solar system observations, terrestrial gravitational experiments and the binary pulsar PSR 1913+16 observations. ",
        "conclusions": "There is a large enough sample of galaxy data which fits our predicted NGT acceleration law to warrant taking seriously the proposal that NGT can explain the flat rotational velocity curves of galaxies without exotic dark matter. We do predict that there will be galaxy matter additional to that due to visible stars and baryons, associated with the energy density $\\rho_m$ residing in the skew field $g_{[\\mu\\nu]}$. It is interesting to note that we can fit the rotational velocity data of galaxies in the distance range $0.02\\,{\\rm kpc} < r < 70\\,{\\rm kpc}$ and in the mass range $10^5\\, M_{\\odot}< M < 10^{11}\\,M_{\\odot}$. without exotic dark matter halos. We are required to investigate further the behavior of the NGT predictions for distances approaching the cores of galaxies, using a disk profile density. The lensing of clusters can also be explained by the theory without exotic dark matter in cluster halos.  We are able to obtain agreement with the observations in the solar system and terrestrial gravitational experiments for suitable values of the parameter $M_0$. This required that we scale $G$ and the parameter $\\sqrt{M_0}$ as functions of the mean orbital radius $\\langle r_{\\rm orb}\\rangle$ of bound systems with the behavior $\\alpha_i=(\\sqrt{M_0})_i=\\langle r_{\\rm orb}\\rangle_i^{3/2}$.  A numerical solution of the NGT field equations for cosmology must be implemented to see whether the theory can account for the large scale structure of the universe and account for galaxy formation and big bang nucleosynthesis, without requiring the existence of exotic, undetected dark matter and a positive cosmological constant to describe dark energy."
    },
    "0404/astro-ph0404388_arXiv.txt": {
        "abstract": "The realization that direct imaging of extrasolar planets could be technologically feasible within the next decade or so has inspired a great deal of recent research into high-contrast imaging.  We ourselves have contributed several design ideas, all of which can be described as {\\em shaped pupil coronagraphs}.  In this paper, we offer a complete and unified survey of asymmetric shaped pupils designs, some of which have been published in our previous papers.  We also introduce a promising new design, which we call {\\em barcode masks}.  These masks achieve the required contrast with a fairly large discovery zone and throughput but most importantly they are perhaps the easiest to manufacture and might therefore stand up best to a refined analysis based on vector propogation techniques. ",
        "introduction": "The discovery of more than 100 extrasolar Jupiter-sized planets in just the last decade has generated enormous interest, both among astronomers and the public, in the problem of discovering and characterizing Earthlike planets.  NASA is already planning its next large space-based observatory, the \\emph{Terrestrial Planet Finder (TPF)}, with a planned launch date toward the end of the next decade.  TPF's primary objective will be to discover Earthlike planets and characterize them for indications of life.  While the technical challenges for TPF are great,  foremost among them is the problem of high-contrast imaging.  In order to discover as many planets as possible, it is necessary to design an imaging system that achieves very high contrast between the parent star and the nearby planet.  An earlier study by \\cite{ref:Brown} indicates that a $D = 4$m class visible-light (i.e. $400\\text{nm} \\le \\lambda \\le 650\\text{nm}$) instrument ought to be able to discover about 50 extrasolar Earth-like planets if it can provide contrast of $10^{-10}$ at an angular separation of $3\\ld$ and that a $4 \\times 10$m class telescope ought to be able to discover about 150 such planets if it can provide the same contrast at a separation of $4\\ld$.  Many approaches have been examined for designing a telescope with this level of contrast.  The most promising fall into two broad categories---nulling interferometers (operating in the infrared) and coronagraphs (operating in the visible).  One important subset of coronagraphs, referred to as shaped pupils, has been gaining interest (\\cite{ref:Spergel, ref:KVSL, VSK02, VSK03}).  These are apodized entrance pupils that rely solely on one/zero binary openings.   In this paper we present a unified treatment of the one-dimensional shaped pupil designs, a few of which have been previously published, and we introduce a promising new design which we call {\\em barcode masks}.   The paper is organized as follows.  In the next section we briefly review the relationship between pupil-plane apodization and the corresponding image-plane electric field and point-spread function.  In the following section we discuss the various performance metrics that characterize a planet finding telescope.  In Section 4 we show how optimization problems can be formulated and used to find the ``best'' entrance pupil apodization for high contrast in one-dimensional square apertures as well as azimuthally symmetric (circular or elliptical) pupils.  In Section 5 we describe how these results can be used to find various families of assymmetric shaped pupils.  In section 6 we present a preliminary sensitivity analysis to determine the tolerances of the pupil approaches to manufacturing errors. ",
        "conclusions": ""
    },
    "0404/astro-ph0404341_arXiv.txt": {
        "abstract": "We report results from Exploratory Time observations of the double-pulsar system PSR~J0737$-$3039 using the Green Bank Telescope (GBT).  The large gain of the GBT, the diversity of the pulsar backends, and the four different frequency bands used have allowed us to make interesting measurements of a wide variety of phenomena. Here we briefly describe results from high-precision timing, polarization, eclipse, scintillation velocity, and single-pulse work. ",
        "introduction": "In December 2003, we proposed for and were awarded Exploratory Time to observe the spectacular double-pulsar system PSR~J0737$-$3039 \\citep[hereafter 0737;][and contributions to this volume from Burgay, Manchester, Kramer, and others]{bdp+03,lbk+04} as part of the NRAO Rapid Response Science program\\footnote{\\tt http://www.vla.nrao.edu/astro/prop/rapid/}.  We observed 0737 five times at four different frequencies (427, 2$\\times$820, 1400, and 2200\\,MHz) using the Berkeley-Caltech Pulsar Machine (BCPM) and the Green Bank Pulsar Processor (GBPP; which measures full polarization information).  For two of the observations (one each at 427 and 820 MHz) we also used the new GBT Spectrometer SPIGOT card (a correlator-based instrument that outputs lags at 25\\,MB/s) in some of its first scientific observations.  These data are public and can be obtained from NRAO. ",
        "conclusions": "The fantastic and highly-varied science to come from these five observations gives some indication of how important a role the GBT will play in future work on 0737.  These observations provide the best constraints yet available on the inclination of the orbits, the mass ratio of the pulsars, the geometry of \\A's radio emission, the asymmetries and achromaticity of \\A's eclipses (and the nature of \\B's magnetosphere that cause them), the orbital modulation of \\B's pulsed flux, and the systemic velocity of 0737 and the kick that caused it.  The high gain of the telescope, the availability of many sensitive and useful receivers, and the new pulsar backends (including the SPIGOT and two coherent de-dispersion systems) will insure that these are only the first of a long series of GBT-based results from the double-pulsar."
    },
    "0404/astro-ph0404177_arXiv.txt": {
        "abstract": "We present results of a self-consistent model of the spectral energy distribution (SED) of starburst galaxies. Two parameters control the IR SED, the mean pressure in the ISM and the destruction timescale of molecular clouds. Adding a simplified AGN spectrum provides mixing lines on IRAS color : color diagrams. This reproduces the observed colors of both AGNs and starbursts. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404611_arXiv.txt": {
        "abstract": "The \\lya\\ emitters found at $z \\sim 4.5$ by the Large Area Lyman Alpha (LALA) survey have high equivalent widths in the \\lya\\ line, which can be produced by either narrow-lined active galactic nuclei (AGNs) or by stellar populations with a very high proportion of young, massive stars. To investigate the AGN scenario, we obtained two deep $Chandra$ exposures to study the X-ray nature of the \\lya\\ emitters. The 172 ks deep $Chandra$ image on the LALA Bo\\\"{o}tes field was presented in a previous paper (Malhotra et al. 2003), and in this paper we present a new $Chandra$ deep exposure (174 ks) on the LALA Cetus field, which doubled our sample of X-ray imaged \\lya\\ sources, and imaged the brightest source among our \\lya\\ emitters. None of the 101 \\lya\\ sources covered by two $Chandra$ exposures was detected individually in X-ray, with a $3 \\sigma$ limiting X-ray flux of $F_{\\rm 0.5-10.0 keV} < 3.3 \\times 10^{-16}$ \\fluxunit\\ for on-axis targets. The sources remain undetectable in the stacked image, implying a 3$\\sigma$ limit to the average luminosity of $L_{\\rm 2-8keV} < 2.8 \\times 10^{42}$ ergs s$^{-1}$. The resulting X-ray to \\lya\\ ratio is $>$ 21 times lower than the ratios for known high redshift type-II quasars. Together with optical spectra obtained at Keck, we conclude that no evidence of AGN activity was found among our \\lya\\ emitters at $z \\sim 4.5$. ",
        "introduction": "The LALA survey comprises two primary fields, located in Bo\\\"{o}tes (J142557+3532) and in Cetus (J020520--0455). Each field is $36\\arcmin \\times 36\\arcmin$  in size, corresponding to a single field of the 8192$\\times$8192 pixel Mosaic CCD cameras at the National Optical Astronomy Observatory's 4 meter telescopes. For each field, 5 partially overlapping narrow band filters covering $4.37 < z < 4.57$ for \\lya\\ were utilized to obtain deep narrowband images. These filters have full width at half maximum (FWHM) $\\approx 80$\\AA\\ and are spaced at $\\approx 40$\\AA\\ intervals (central wavelengths $\\lambda_c \\approx 6559$, $6611$, $6650$, $6692$, and $6730$\\AA). For the Bo\\\"{o}tes field, 3 more narrowband images covering $z \\sim 5.7$ (Rhoads et al. 2003) and $z \\sim 6.5$ (Rhoads et al. 2004) are also available. Imaging data reduction followed the methods described in Rhoads et al (2000), and \\lya\\ candidates were selected using criteria described by Rhoads \\& Malhotra (2001). This resulted in a total of $\\sim$ 400 good candidates at $z \\sim 4.5$, which will be presented in a future paper (Rhoads et al. 2004, in preparation).  Two deep $Chandra$ ACIS images were obtained for the LALA Bo\\\"{o}tes and Cetus fields. The $Chandra$ observations were designed to maximize the number of large equivalent width sources within the ACIS-I field of view. The 172 ks (net exposure) $Chandra$ exposure of the  LALA Bo\\\"{o}tes field was taken in 2002, and details of the X-ray data analyses can be found in M03 and Wang et al. (2004). The new 174 ks (net exposure) $Chandra$ ACIS exposure on the LALA Cetus field was obtained in very faint (VFAINT) mode on UT 2003 June 13-15. Data reduction was done with the package CIAO 2.2.1 (see http://asc.harvard.edu/ciao), following the procedures described in M03. The average offset between the X-ray and optical images was obtained by comparing X-ray source positions with optical counterparts (whenever found). The derived 0.3$\\arcsec$ offset has been applied for all subsequent analyses. No obvious rotation and plate-scale effects were discovered. We ran WAVDETECT (Freeman et al. 2002) on the soft (0.5 -- 2.0 keV), hard (2.0 -- 7.0 keV), and total band (0.5 -- 7.0 keV) X-ray images. The detailed X-ray data reduction and the detected X-ray sources will be presented in a future paper (Wang et al. 2004, in preparation). ",
        "conclusions": ""
    },
    "0404/astro-ph0404427_arXiv.txt": {
        "abstract": "We present an investigation of satellite galaxies in the outskirts of galaxy clusters taken from a series of high-resolution \\nbody\\ simulations. We focus on the so-called ``backsplash population'', i.e. satellite galaxies that once were inside the virial radius of the host but now reside beyond it. We find that this population is significant in number and needs to be appreciated when interpreting the various galaxy morphology environmental relationships and decoupling the degeneracy between nature and nurture. Specifically, we find that approximately half of the galaxies with current clustercentric distance in the interval 1 -- 2 virial radii of the host are backsplash galaxies which once penetrated deep into the cluster potential, with 90\\% of these entering to within 50\\% of the virial radius. These galaxies have undergone significant tidal disruption, loosing on average 40\\% of their mass. This results in a mass function for the backsplash population different to those galaxies infalling for the first time. We further show that these two populations are kinematically distinct and should be observable within existent spectroscopic surveys. ",
        "introduction": "The relationship between galaxy morphology and local environment (i.e. the morphology-density relation) was first noticed by Hubble~\\& Humason (1931), where they reported that field and cluster galaxy populations differ. Oemler (1974) extended this finding by showing that the relationship held for differing clusters richness. The field truly emerged when Dressler (1980) demonstrated the strong relationship over five orders of magnitude between the local density of galaxies and the proportions of different morphological types. Bhavsar (1981), de Souza et al. (1982), and Postman~\\& Geller (1984) extended this work further to include the environments of both loose and compact groups. Recently, Aguerri~\\ea (2004) performed a thorough analysis of 116 bright galaxies in the Coma cluster, finding that bluer objects are located at larger projected radii while simultaneously showing a larger velocity dispersion than their red counterparts. Moreover, the bluest objects also host the most prominent disks contrary to systems observed close to the cluster centre or in high-density environments.  Environmental dependence of galactic stellar populations is also seen in the Butcher-Oemler effect (Butcher \\& Oemler 1978, Kodama \\& Bower 2001) with clusters at higher redshifts showing a greater fraction of blue objects than are seen at present.   The above observational work supports the idea that galaxies in clusters are substantially different from galaxies in the field.  But the origins of these morphology-density relationships are still not fully understood with several large and small scale mechanisms proposed to explain their existence, including ram pressure stripping (Gun \\& Gott 1972), tidal stripping/star formation (Merritt 1983, 1984), starvation (Larson, Tinsley \\& Caldwell 1980), galaxy merger and harassment (Icke 1985; Moore \\ea 1996) and dynamics (Tsuchiya \\& Shimada 2000).  Recent cosmological simulations (both hydrodynamical and \\nbody) have provided valuable insight into the mechanisms responsible for the morphology-density relationship (Springel \\ea 2001; Goto \\ea 2003; Okamoto \\& Nagashima 2003). In the analysis which follows, we focus on the dynamics of satellite galaxies taken from a series of high-resolution, fully self-consistent, cosmological simulations of eight galaxy clusters.  We concentrate on the outskirts of these clusters, i.e. distances in the range $[R_{vir},2.5 R_{vir}]$, which (observationally) have only recently been probed through wide-field optical imaging and spectroscopy (Miyazaki~\\ea 2002; Lewis~\\ea 2002). We will demonstrate that a rich population of galaxies exist beyond the virial radius most of which have previously spent time near the cluster centre and can be seen in Figure~\\ref{galback}. We characterise the spatial, velocity and mass properties of this population and contrast these with those of the spatially coincident newly infalling galaxies.  Our work complements the earlier studies of Balogh \\ea (2000) and Mamon \\ea (2004). Balogh \\ea investigated the \\textit{particle} backsplash from cosmological simulations and found that 50 $\\pm$ 20\\% of the particles within $[R_{200},2R_{200}]$ had passed through the $R_{200}$ radius. Mamon \\ea recently extended this work to calculate the maximum backsplash distance for particles to be $2.5 R_{100}$ and for galaxies $1.7 R_{100}$.  \\begin{figure} \\centerline{\\psfig{file=eps/backsplash.eps,width=\\hsize}} \\caption{A simulated cluster at z=0 with the virial radius indicated by the dark sphere. The other line in this figure represents the orbital path of a ``backsplash'' galaxy. Such a galaxy has previously spent time near the cluster centre but now lies outside the virial radius of the cluster.} \\label{galback} \\end{figure}  The outline of the paper is as follows. Section~\\ref{Computation} provides a description of the cosmological simulations employed.  We investigate the number distribution of galaxies in the cluster outskirts in Section~\\ref{number}, with the mass distribution discussed in Section~\\ref{mass} and the velocity distribution of the satellites investigated in Section~\\ref{velocity}. We finish with our summary and conclusions in Section~\\ref{Conclusions}. ",
        "conclusions": "\\label{Conclusions}   Observational data supports the idea that galaxies in clusters are substantially different from galaxies in the field. There is a clear correlation between galaxy morphology and density of the local environment. However, the origin of this relation is far from being understood.  In this study we have presented an analysis of satellite galaxies that once passed through the virial radius close to the centre of their respective host halo, but are now found outside the virial radius in the outskirts of the cluster. We have shown that this backsplash population is not negligible and needs to be accounted for when interpreting the various galaxy morphology relationships and decoupling the degeneracy between nature and nurture.  We must also appreciate that the infalling population is not expected to be pristine. Rather, we would expect that infall galaxies have undergone some sort of pre-processing in groups before entering the cluster too as indicated by the sub-subhalos in \\Fig{RminRz}.  Our results can be summarized as follows:  \\begin{itemize}  \\item 30\\% of all galaxies that ever came closer to the host than its virial radius are now located in the range [$R_{\\rm vir}$,2.5$R_{\\rm vir}$],  \\item 50\\% of all galaxies in the region [$R_{\\rm vir}$,2$R_{\\rm vir}$] are backsplash galaxies,  \\item 90\\% of the backsplash galaxies penetrated deeper than 50\\% of $R_{\\rm vir}$ into the host's potential,  \\item during their passage through the cluster, on average the backsplash galaxies lose 40\\% of their mass, thus  \\item the mass spectrum of the backsplash population has a steeper power-law slope than their infalling counterparts, thus it has fewer massive galaxies and more light ones,  \\item the velocities of the infalling satellites is too small to account for ram-pressure stripping in the cluster outskirts,  \\item the backsplash population has a factor of two smaller relative velocity than the infalling satellites, making it kinematically distinct.  \\end{itemize}  When transforming the last result into the observers plane though, the velocity separation between the infalling and backsplash population is removed. However, the backsplash population should still be detectable as it is responsible for a continuous rise in the distribution function towards low line-of-sight velocities.  Our results suggest that we not only expect the backsplash population to experience various large-scale transformation mechanisms, but also small-scale ones, undergoing starvation, ram pressure stripping, tidally triggered star formation and significant tidal stripping."
    },
    "0404/astro-ph0404561_arXiv.txt": {
        "abstract": "{ The 68~000 $I$-band light curves of variable stars detected by the \\OG\\ survey in the Large and Small Magellanic Clouds (MCs) are fitted by Fourier series, and also correlated with the \\DE\\ and \\M\\ {\\it all-sky release} databases and with lists of spectroscopically confirmed M-, S- and C-stars. Lightcurves and the results of the lightcurve fitting (periods and amplitudes) and \\DE\\ and \\M\\ magnitudes are presented for 2277 M-,S-,C-stars in the MCs. The following aspects are discussed: the $K$-band period-luminosity relations for the spectroscopically confirmed AGB stars, period changes over a timespan of about 17 years in a subset of about 400 LPVs, and candidate obscured AGB stars. The use of a sample of spectroscopically confirmed variables allows me to show specifically that almost all carbon stars are brighter than the tip of the RGB, and occupy sequences A+,B+,C and D. It is shown (for the LMC where there is a sufficient number of spectroscopically identified M-stars) that for sequences A+,B+,C the M-stars are on average fainter than the C-stars, as expected from an evolutionary point of view and previously observed in MC clusters. However, this is not so for sequence ``D'', suggesting that the origin of the so-called Long Secondary Periods is not related to an evolutionary effect. The fraction of objects that has a period on sequence ``D'' is also independent of chemical type. Three stars are identified that have been classified as oxygen-rich in the 1970s and carbon-rich in 1990s. Possibly they underwent a thermal pulse in the last 20 years, and dredged-up enough carbon to switch spectral type. The observations over almost two decades seem to suggest that up to 10\\% of AGB variables changed pulsation mode over that time span. More robust estimates will come from the ongoing and future (microlensing) photometric surveys. A sample of 570 variable red objects ($(J-K) >$ 2.0 or $(I-K) >$ 4.0) is presented in which most stars are expected to be dust obscured AGB stars. Estimates are presented for cut-offs in $(J-K)$ which should be applied to minimise dust obscuration in $K$, and based on this, C- and O-star $K$-band $PL$-relations for large amplitude variables in the SMC and LMC are presented. ",
        "introduction": "In the course of the micro lensing surveys in the 1990's, the monitoring of the Small and Large Magellanic Clouds has revealed an amazing number and variety of variable stars. A big impact was felt and is being felt in many areas of variable star research, like Cepheids and RR Lyrae stars. Also in the area of Long Period Variables (LPVs) and AGB stars there has been remarkable progress. Wood et al. (1999) and Wood (2000) were the first to identify and label different sequences ``ABC'' thought to represent the classical Mira sequence (``C'') and overtone pulsators (``A,B''), and sequence ``D'' which is still unexplained (Olivier \\& Wood 2003, Wood 2003, 2004). Stars on these sequence are sometimes referred to as having Long Secondary Periods--LSPs. This view has subsequently been confirmed and expanded upon by Noda et al. (2002), Lebzelter et al. (2002), Cioni et al. (2003), Ita et al. (2004a,b) and Kiss \\& Bedding (2003, 2004). These works differ in the source of the variability data ({\\sc macho}, \\OG, {\\sc eros}, {\\sc moa}), area (SMC or LMC), associated infrared data (Siding Spring 2.3m, \\DE, \\M, {\\sc sirius}), and selection on pulsation amplitude or infrared colours.  The present paper considers the \\OG\\ data for both LMC and SMC. Ita et al. (2004a) only consider the \\OG\\ data in overlap with their {\\sc sirius} IR observations in LMC and SMC, and Kiss \\& Bedding consider only stars in the SMC with \\M\\ data with $(J-K) > 0.9$.  Also in contrast to previous studies, emphasis is put on spectroscopically confirmed AGB stars (i.e. M-, S- and C-stars). In other studies M- and C-stars are usually identified photometrically by using a division at a $(J-K)$ colour of $\\sim$1.4 mag. This paper specifically addresses the properties of known carbon stars in relation to sequences ``ABCD''.  The paper is structured as follows. In Section~2 the \\OG, \\M\\ and \\DE\\ surveys are described. In Section~3 the model is presented, both in terms of the actual lightcurve fitting, and the post-processing. Section~4 presents the results. Discussed are the $K$-band $PL$-relation for the spectroscopically confirmed AGB stars, period changes over a timespan of almost 2 decades, and a sample of very red obscured AGB star candidates. The conclusions are summarised in Section~5.  Some of this work, and the star-to-star comparison of periods derived by me from {\\sc macho} and \\OG\\ data and literature values are described in Groenewegen (2004). ",
        "conclusions": "This paper addresses several aspects of the pulsational character of late-type stars in the Magellanic Clouds. The main focus is on the $K$-band $PL$-relation of almost 2300 spectroscopically confirmed M-,S- and C-stars. This sample avoids to make the clearly incorrect approximation made in other studies that M-stars and carbon stars can be separated at a colour $(J-K)$ = 1.4.  The present observations however do not allow the presentation a comprehensive picture of the evolution of pulsation periods of M- and C-stars. This would require a more detailed (AGB star) population synthesis study including pulsation properties. As previous MCs studies have found for AGB stars in general (Wood et al. (1999), Wood (2000), Noda et al. (2002), Lebzelter et al. (2002), Cioni et al. (2003), Ita et al. (2004a,b) and Kiss \\& Bedding (2003, 2004)), it is found specifically that both M-, and C-stars tend to occupy preferentially sequences B+ and then C for increasing amplitude.  For a given amplitude, the C-stars tend to have the longer period, and, for a every sequence they are more luminous. This effect has previously been observed in MC clusters (e.g. FMB) and is in qualitative agreement with evolutionary calculations that predict that C-stars evolve from M-stars.  Many objects have one period that falls in box ``D''. In the spectroscopically selected sample, 211 of 859 SMC stars (= 24.6\\%) have a period that falls in box ``D'', and 318 of the 1418 LMC stars (= 22.4\\%, namely 229/1064 = 21.5\\% of C-stars and 89/354 = 25.1\\% of M-stars). As will be discussed below, for at least some stars of the IR selected sample this is due to the fact that they are weaker in $K$ because of dust obscuration. For the overall majority of the stars in the spectroscopically selected sample this is not an issue. The reason why some late-type stars appear on that location of the $PL$-diagram is unexplained, see the discussion in Olivier \\& Wood (2003), Wood (2003, 2004). The classical large-amplitude Mira variables appear on sequence ``C'' (see the last panels in Figure~\\ref{fig:PL}) and are believed to be fundamental-mode pulsators, hence longer (radial mode pulsation) periods should not exist.  The present paper does not shed light into the nature of the LSP phenomenon, except that Figure~\\ref{Fig-HistoSeq-LMC} indicates that the $K$-band luminosity function for sequence ``D'' is essentially the same for the M- and C-stars, while for sequences ``A,B,C'' the C-stars are brighter, as expected from an evolutionary point of view. The luminosity function of objects on sequence ``D'' is fainter from those of the other sequences. This is due to the fact that periods longer than 800 days are underrepresented because of stricter selection rules. This affects predominatly sequence ``D'' objects with $K_0$ \\less 10.5 mag. The fact that the fraction of all object with periods on sequence ``D'' is essentially the same for C- and M-stars and that the $K$-band luminosity of C- and M-stars of sequence ``D'' objects is very similar suggests that the LSP phenomenon is unrelated to chemical type, and hence seems unrelated to a pulsation phenomenon.  For a few hundred variables it was possible to look for period changes over a timespan of typically 17 years. Almost all come from the studies by Hughes (1989) and Hughes \\& Wood (1990). They identified medium to large amplitude variables from photographic material using typically 21 observations in the time span 1977 to 1984. Out of 370 objects, 36 have been identified that seemingly changed pulsation mode (or at least changed ``box'') between $\\sim$1980 and the time of the \\OG\\ observations, and another 30 objects that changed pulsation period by more than 10\\%. This ratio of about 10\\% (36/370) is similar to the study of GLE03 who found large period changes in 3 out of 42 Mira variables studied. A caveat is that the original historical data points have never been published, and it would certainly be preferable to be able to phase the old data with the current period to see if in fact the period change is real. For the moment I consider the 10\\% change of pulsation sequence over $\\sim$17 years as an upper limit. By comparison, Zijlstra \\& Bedding (2003) find that only of order of 1\\% of well-known Miras show evidence for period changes. The understanding for MC objects may improve because of ongoing (e.g. {\\sc ogle-iii}) and future surveys.  Finally a sample of stars was studied selected on infrared colours, namely redder than the majority of the spectroscopically selected sample. It should be pointed out that there is no proof that these are AGB stars (except for the few ones in overlap with the spectroscopically selected sample) and they make suitable targets for spectroscopic follow-up to determine their spectral type. Many of these stars also have a period located in box ``D'' but in this case the effect of obscuration by dust must be considered. This also has implications when using samples of variables to determine $K$-band $PL$-relations.  In the last section one explicit example was shown, namely MSX 83, for which the SED was constructed and fitted with a dust radiative transfer model. Its period of 611 day, and $K$ magnitude of 10.58 would put it in ``box D''. However, when running the radiative transfer model without mass-loss the $K$-magnitude brightens to 9.17 mag (this value is somewhat dependent on the central star model atmosphere assumed) putting it on the extension of box ``C'', consistent with the expected location for a star with an pulsation amplitude of 0.66 mag.  In reverse this implies that when studying the $K$-band $PL$-relation, and when multi-colour data is available, a cut-off in colour should be applied in order to avoid a bias by including stars that are dust obscured in $K$.  Although this seems obvious, a quantification of where this cut-off should be placed and its actual application are rare in the literature; In fact I could only find one instance. Glass et al. (1995) mention they exclude some faint outliers with $(K-L) \\sim 2$ in the Sgr {\\sc i} bulge field (which corresponds to roughly $(J-K) \\sim$ 3.5, Glass, 1986), but did not impose a colour criterium a-priori. In other papers where dust obscuration in the studied variables should play a role the bolometric $PL$-relation is studied (e.g. WFLZ), which circumvents this problem in a natural way.  It is difficult to give an exact colour cut-off to apply, since this depends on the colours involved, the dust properties and the evolution of mass-loss on the AGB. Based on the calculations presented in the last columns in Table~\\ref{tab-models} stars with colours $(J-K)_0 > 2.0$ should certainly be avoided, and a stricter criterium would be to include M-stars only when $(J-K)_0 <$ 1.4 and C-stars only when $(J-K)_0 <$ 1.7.  Finally, applying these latter cut-offs to the spectroscopically selected stars in box ``C'' with amplitudes $>0.45$ mag and $>0.15$ mag, to improve the statistics, the $K$-band $PL$-relations listed in Table~\\ref{tab-pl} have been derived, where also some relations from the literature are listed. The period distribution of these samples is shown in Figure~\\ref{PerDist}. It would be interesting to fit the SEDs of all infrared selected stars to be able to include the dust-correced $K$-magnitudes in these $PL$-relations.  There is very good agreement between the new relations (including what appears to be the first Mira $K$-band $PL$-relation in the SMC) and previous works. The formal error on the zero point and slope have become smaller, because of the larger sample size, but the overall rms are still larger because the photometry used in FGWC and GLE03 are averages over the lightcurve while the present data are at best averages of two measurements.   \\begin{figure*}[ht] \\begin{minipage}{0.48\\textwidth} \\resizebox{\\hsize}{!}{\\includegraphics{PerDistLmc_C_0.15_9.99.ps}} \\end{minipage} \\hfill \\begin{minipage}{0.48\\textwidth} \\resizebox{\\hsize}{!}{\\includegraphics{PerDistSmc_C_0.15_9.99.ps}} \\end{minipage} \\caption{ Period distribution of the LMC (left panel) and SMC (right panel) variables in box ``C'' with $I$-band amplitudes larger than 0.15 magnitudes. Shown are the histograms for the M-stars (vertical lines), C-stars (hatched), and total. } \\label{PerDist} \\end{figure*}  \\begin{table*} \\caption{$K$-band $PL$-relations in box ``C'', of the form $m_{\\rm K}= a \\log P + b$ } \\begin{tabular}{ccrcl} \\hline $a$       &       $b$        &  N  &  rms & remarks  \\\\ \\hline $-3.78 \\pm 0.24$ & $20.17 \\pm 0.58$ &  44 & 0.26 & O-stars, LMC, Amplitude $> 0.45$ mag \\\\ $-3.50 \\pm 0.27$ & $19.42 \\pm 0.67$ &  58 & 0.22 & C-stars, LMC, Amplitude $> 0.45$ mag \\\\ $-4.14 \\pm 0.85$ & $21.34 \\pm 2.06$ &  14 & 0.25 & C-stars, SMC  Amplitude $> 0.45$ mag \\\\ $-3.52 \\pm 0.16$ & $19.56 \\pm 0.38$ &  83 & 0.26 & O-stars, LMC, Amplitude $> 0.15$ mag \\\\ $-3.56 \\pm 0.15$ & $19.57 \\pm 0.38$ & 181 & 0.25 & C-stars, LMC, Amplitude $> 0.15$ mag \\\\ $-3.67 \\pm 0.24$ & $20.22 \\pm 0.60$ &  92 & 0.24 & C-stars, SMC, Amplitude $> 0.15$ mag \\\\ \\\\ $-3.47 \\pm 0.19$ & $19.48 \\pm 0.45$ &  29 & 0.13 & M-Miras, LMC, FGWC \\\\ $-3.52 \\pm 0.21$ & $19.64 \\pm 0.49$ &  26 & 0.13 & M-Miras, LMC, GLE03 \\\\ $-3.30 \\pm 0.40$ & $18.98 \\pm 0.98$ &  20 & 0.18 & C-Miras, LMC, FGWC \\\\ $-3.56 \\pm 0.17$ & $19.64 \\pm 0.42$ &  54 & 0.25 & C-Miras, LMC, Groenewegen \\& Whitelock (1996)  \\\\ \\hline \\end{tabular} \\label{tab-pl} \\end{table*}"
    },
    "0404/astro-ph0404096_arXiv.txt": {
        "abstract": "{\\begin{center} {\\it \\large Abstract\\vspace{-.5em}\\vspace{0pt}} \\end{center}} \\def\\endabstract{\\par} \\usepackage{amsfonts} \\usepackage{amssymb} \\usepackage{verbatim} \\usepackage{graphicx} \\usepackage{hyperref}  \\ifPDF \\hypersetup{% pdftoolbar=treu,% pdfmenubar=true} \\else \\usepackage{epsfig} \\usepackage{pslatex} \\fi  \\setlength{\\hoffset}{-1cm} \\setlength{\\voffset}{-1cm} \\setlength{\\textheight}{23.6cm} \\setlength{\\textwidth}{18.1cm} \\setlength{\\parindent}{0.7cm} \\setlength{\\parskip}{0ex} \\setlength{\\columnsep}{5.4mm} \\setlength{\\topmargin}{1.3cm} \\setlength{\\oddsidemargin}{0.0cm} \\setlength{\\evensidemargin}{0.0cm}  \\renewcommand{\\thesection}{\\rm \\Roman{section}. } \\renewcommand{\\thesubsection}{\\it \\Alph{subsection}. } \\renewcommand{\\thesubsubsection}{\\it  \\arabic{subsubsection}) } \\renewcommand{\\thetable}{ \\arabic{table}} \\renewcommand{\\thefigure}{ \\arabic{figure}}  \\renewcommand\\floatpagefraction{.9} \\renewcommand\\topfraction{.9} \\renewcommand\\bottomfraction{.9} \\renewcommand\\textfraction{.1}  \\newcommand{\\ssection}[1]{% {\\begin{center} {\\it \\large Abstract\\vspace{-.5em}\\vspace{0pt}} \\end{center}} \\def\\endabstract{\\par}  \\setcounter{totalnumber}{50} \\setcounter{topnumber}{50} \\setcounter{bottomnumber}{50}    \\title{ \\vspace{-1cm} \\Large {The Baikal Neutrino Telescope: Results, Plans, Lessons}} \\author{\\large C. Spiering for the BAIKAL Collaboration \\\\ \\normalsize DESY Zeuthen, Platanenallee 6, D-15738 Zeuthen, Germany\\\\ \\normalsize christian.spiering@desy.de}   \\begin{document} \\date{}  We review recent results on the search for high energy extraterrestrial neutrinos, neutrinos induced by WIMP annihilation and neutrinos coincident with Gamma Ray Bursts as obtained with the Baikal neutrino telescope NT-200. We describe the moderate upgrade of NT-200 towards a $\\sim$10 Mton scale detector NT-200$+$. We finally draw a few lessons from our experience which may be of use for other underwater experiments. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404575_arXiv.txt": {
        "abstract": "We present long-slit HST/STIS measurements of the ionized-gas kinematics in the nucleus of three disk galaxies, namely NGC 2179, NGC 4343, NGC 4435. The sample galaxies have been selected on the basis of their ground-based spectroscopy, for displaying a strong central velocity gradient for the ionized gas, which is consistent with the presence of a circum nuclear keplerian disk (CNKD, \\cite{bert1998}; \\cite{funes2002}) rotating around a super massive black hole (SMBH). For each target galaxy we obtained the H$\\alpha$ and [NII] 6583 \\AA\\ kinematics along the major axis and two 0\\farcs25 parallel offset positions. Out of three objects only NGC 4435 turned out to have a disk of ionized gas in regular motion and a regular dust-lane morphology. Preliminary modeling indicates a SMBH mass ($M_\\bullet$) one order of magnitude lower than the one expected from the $M_\\bullet-\\sigma_c$ relation for galaxies (\\cite{ferr2000}; \\cite{geb2000}). ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404269_arXiv.txt": {
        "abstract": "{We show that the recently developed thermal model which successfully describes how jets are launched by young stellar objects, when applied to system containing disk-accreting white dwarfs naturally explain the otherwise surprising absence of jets in cataclysmic variable stars. Our main argument uses the crucial element of the thermal model, namely that the accreted material is strongly shocked due to large gradients of physical quantities in the boundary layer, and then cools on a time scale longer than its ejection time from the disk. In our scenario 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 two criteria in that model, for the shock formation and for the ejection of mass, we find the mass accretion rate above which jets could be blown from accretion disks around young stellar objects and white dwarfs. We find that these accretion mass rates are $\\dot M ({\\rm YSO}) \\ga 10^{-7} M_\\odot \\yr ^{-1}$ and $\\dot M ({\\rm WD}) \\ga 10^{-6} M_\\odot \\yr ^{-1}$ for young stellar objects and white dwarfs respectively. Considering the uncertainties of the model, these limits could overestimate the critical value by a factor of $\\sim 10$. } ",
        "introduction": "\\label{sec:intro}  It is widely believed that most astrophysical jets, and all massive jets (to distinguish from low density hot-plasma jets from radio pulsars) are launched from accretion disks (Livio 1999, 2000a). This belief is supported by observations of jets in Young Stellar Object (YSOs), Low-Mass and High-Mass X-ray Binaries (LMXBs \\& HMXBs) and Active Galactic Nuclei (AGNs) which all are systems containing accretion disks (or at least accretion flows with a considerable amount of angular momentum). The apparent universality of the accretion disk--jet relation is spoiled by one class of systems: Cataclysmic Variables (CVs). They are close binary systems in which a white dwarf accretes matter lost by its Roche-lobe filling low-mass companion (see Warner 1995 for a review). For weak enough white-dwarf's magnetic fields CVs posses accretion disks. But no jets have ever been observed from these numerous and extremely well observed binaries. The one reported occurrence (Shahbaz et al. 1997) failed to be confirmed (O'Brien \\& Cohen 1998) and the system itself is most probably not a CV but a Super Soft X-ray Source (SSXS; Knigge, King \\& Patterson 2000). However, jets have been observed emanating from some other SSXSs (see below) which shows that the absence of jets in CVs cannot be attributed to some special properties of white dwarfs in binary systems since white dwarfs are also present in SSXSs, and symbiotic systems, some of which blow jets.  One can expect that the absence of jets in CVs could tell us something about the still mysterious jets launching mechanism. Commonly it is assumed that magnetic fields play crucial roles in the formation of jets. Magnetic fields can appear in three types of roles:  (1) In triggering the jet ejection events, e.g., by causing instabilities in the disk. These types of ``magnetospheric\" MHD instabilities could exist in accretion disks even when the central star has no magnetic field (e.g., Li \\& Narayan 2004 and references therein). 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 De Villiers, Hawley, \\& Krolik (2004).  (2) In accelerating the jets. There are many models and countless of papers on this subject. Basically, most models 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 by K\\\"onigl and Pudritz (2000), Shu et al. (2000) and Ferreira (2002). 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, Priest \\& Bardou 1996). In addition, it seems that thermal pressure is needed in some of these models (e.g., Ferreira \\& Casse 2004).  (3) 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).  Magnetic jet launching models fail to account for the absence of jets in CVs despite of some interesting suggestions in Livio (2000b). It is therefore justified to consider jet launching mechanisms in which magnetic fields would be deprived of at least one of the three roles. In the present article we show that the absence of jets in CVs is naturally explained by the model of thermal pressure acceleration proposed by Torbett (1984) and Torbett \\& Gilden (1992), and which was developed and extended recently by Soker \\& Regev (2003; hereafter SR03) to explain collimated outflows in YSO.  SR03 interest in thermal pressure acceleration of jets was motivated by new results from recent X-ray observations of YSO. These show that there is essentially no difference between the properties of X-ray emission from YSO with and without outflows (Getman et al.\\ 2002), imposing quite severe constraints on models based on magnetic launching of jets. In the thermally-driven jet model, the magnetic fields are weak, and serve only to re-collimate the outflow at large distances from the source, (role (3) above), and possibly trigger disturbances in the boundary layer (BL), (role (1) above). The BL is the inner layer, where the disk adjust itself to the conditions of the accreting star. A crucial ingredient of the model is that the accreted material is strongly shocked, and that it cools down on a time scale longer than its ejection time from the inner disk. SR03 find that the thermal acceleration mechanism works only when the accretion rate in YSO accretion disk is large enough and the $\\alpha$ parameter of the disk small enough - otherwise the cooling time is too short and significant ejection does not take place. SR03 term these strong shocks `spatiotemporally localized (but not too small!) accretion shocks', or SPLASHes.  In the present paper we extend the analysis of SR03 to white dwarfs (WDs) accretors. We compare the derived conditions for thermally launching jets from accretions disks around YSO and around WDs. YSO refers also to main sequence stars accreting from mass-losing companions stars. We show that the thermally-driven jet model can be extended to jets blown by disks around WDs and explain the absence of jets in CVs. ",
        "conclusions": "\\label{sec:discuss}  In the previous two sections we found the constraints on the  mass accretion rate derived using the two conditions for slow radiative cooling derived by SR03. The first condition is that the strongly shocked gas in the BL will cool slowly, such that the thermal pressure will have enough time to accelerate the jet's material. The constraints on the accretion rate, $\\dot M_s$, are given by equations ({\\ref{accyso1}}) and ({\\ref{accwd1}}), for YSO and WD accretors, respectively. The second condition is that weakly shocked blobs in the BL will expand, and disturb the BL in such a way that a strong shock will develop. The constraints on the accretion rate, $\\dot M_b$, are given by equations ({\\ref{accyso2}}) and ({\\ref{accwd2}}), for YSO and WD accretors, respectively. We should stress the following in regard to these constraints. (1) The constraints are accurate to an order of magnitude. This is for several reasons, e.g.,  the values of $\\alpha$ and $\\epsilon$ are unknown; the demand on the ratio of the cooling time to acceleration time is given to an order of magnitude; the behavior of the gas, e.g., its opacity and pressure, should be treated more accurately with a full 3D numerical code. (2) The first constraint on $\\dot M_s$ is generic to the proposed thermally-accelerated jet model. (3) The second one, on the formation of the disturbances that lead to the formation of strong shocks in the BL, might be less important. This is because other types of disturbances can cause strong shocks to develop in the BL, e.g., MHD instabilities, magnetic eruptions from YSOs, and local thermonuclear events on a WDs.  With these, we note the following. For YSOs (and other main sequence stellar accretors)  the two requirements on the two cooling time-scales basically gives the same constraint on the mass accretion rate $\\dot M({\\rm YSO}) \\ga 10^{-7} M_\\odot \\yr^{-1}$ (eqs. {\\ref{accyso1}} and {\\ref{accyso2}}). This fits observations, which show jets from YSO accreting at such rates (e.g., Cabrit et al.\\ 1990).  For WDs the two constraints give $\\dot M_s \\ga 10^{-9} M_\\odot \\yr^{-1}$ and $\\dot M_b \\ga 10^{-6} M_\\odot \\yr^{-1}$ (eqs. {\\ref{accwd1}} and {\\ref{accwd2}}). The second one is very stringent. The highest accretion rates deduced from observations of nova-like systems or dwarf-novae at maximum are two orders of magnitude lower (see e.g. Warner 1995). Taking into account the order of magnitude uncertainty one could maybe get $\\dot M({\\rm WD}) \\ga 1-3 \\times 10^{-7}$. This is compatible with observations of SSXSs and symbiotic systems. SSXSs are thought to be white dwarfs accreting at rates of $3 \\times 10^{-8}-10^{-6} M_\\odot \\yr^{-1}$ from a companion, and sustaining nuclear burning on their surface (e.g., van den Heuvel et al.\\ 1992; Greiner 1996). To maintain a steady nuclear burning the mass accretion rate should be $3 \\times 10^{-8}-10^{-6} M_\\odot \\yr^{-1}$, where the upper range is for massive WDs (Nomoto 1982). Fast, $\\sim 1000-5000 \\km \\s^{-1}$, collimated outflows have been observed in some SSXSs, RX J0513.9-6951 (Crampton et al.\\ 1996; Southwell et al.\\ 1996), RX J0019.8+2156 (Becker et al. 1998; Quaintrell \\& Fender 1998; Tomov et al.\\ 1998), and RX J0925.7-4758 (Motch 1998). These systems teach us that WDs accreting mass at rates much higher  than those in cataclysmic variables can blow jets. In RX J0925.7-4758 the high luminosity (Motch 1998) implies an accretion rate of $\\sim 10^{-7} M_\\odot \\yr^{-1}$, and the WD radius is $\\sim 0.005 R_\\odot$. With $\\epsilon = 0.05$ the constraint (eq. {\\ref{accwd2}}) is $\\dot M_b \\ga 3 \\times 10^{-7} M_\\odot \\yr^{-1}$. The accretion rate onto RX J0513.9-6951 is even higher (Southwell et al.\\ 1996). Some symbiotic systems are also known to blow jets (Sokoloski 2004; Brocksopp et al.\\ 2004 and references therein). In these symbiotic systems, WD companions accrete from the wind of red giant branch stars or asymptotic giant branch stars, at relatively high rates. In some of the symbiotic systems which blow jets the WD sustains a quasi-steady nuclear burning, similar to SSXSs; in others, there is no nuclear burning (Brocksopp et al.\\ 2004). Still, the mass accretion rate is expected to be high in the later group as well. It is possible that in the later systems the WD are more massive; more massive WD stars require higher mass accretion rates to sustain nuclear burning (Nomoto 1982).  Although no jets have ever been observed in CVs some of them emit winds. 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 mechanisms. They are most probably radiation driven (see e.g., Murray, 2002; Proga 2002).  Finally, it is interesting to investigate what kind of condition one obtains when considering ultra-compact objects such as neutron stars and black holes. In such a case it is more convenient to scale Eq. ({\\ref{acc01}}) with the Schwarzschild radius $R_{\\rm G}=2GM/c^2$ and the Eddington accretion rate $\\dot M_{\\rm Edd}= L_{\\rm Edd}/0.1c^2= 2.3\\times 10^{-8} M_\\odot \\yr^{-1}$. One obtains then \\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}{R_{\\rm G}}, \\label{comp} \\end{equation} which is very close to the value at which low-mass X-ray transients enter (from below) the so called hard/low states associated with the appearance of steady jets (e.g., Fender 2001). Of course there are no boundary layers around accreting black holes so that our model cannot be directly applied to these objects. We will discuss this problem in a future paper."
    },
    "0404/astro-ph0404119_arXiv.txt": {
        "abstract": "We wish to estimate magnetic field strengths of radio emitting galaxy clusters by minimising the non-thermal energy density contained in cosmic ray electrons (CRe), protons (CRp), and magnetic fields. The {\\em classical} minimum energy estimate can be constructed independently of the origin of the radio synchrotron emitting CRe yielding thus an absolute minimum of the non-thermal energy density.  Provided the observed synchrotron emission is generated by a CRe population originating from hadronic interactions of CRp with the ambient thermal gas of the intra-cluster medium, the parameter space of the {\\em classical} scenario can be tightened by means of the {\\em hadronic} minimum energy criterion.  For both approaches, we derive the theoretically expected tolerance regions for the inferred minimum energy densities.  Application to the radio halo of the Coma cluster and the radio mini-halo of the Perseus cluster yields equipartition between cosmic rays and magnetic fields within the expected tolerance regions. In the hadronic scenario, the inferred central magnetic field strength ranges from $2.4~\\umu\\rmn{G}$ (Coma) to $8.8~\\umu\\rmn{G}$ (Perseus), while the optimal CRp energy density is constrained to $2\\% \\pm 1\\%$ of the thermal energy density (Perseus). We discuss the possibility of a hadronic origin of the Coma radio halo while current observations favour such a scenario for the Perseus radio mini-halo.  Combining future expected detections of radio synchrotron, hard X-ray inverse Compton, and hadronically induced $\\gamma$-ray emission should allow an estimate of volume averaged cluster magnetic fields and provide information about their dynamical state. ",
        "introduction": "\\label{sec:intro}  Clusters of galaxies harbour magnetised plasma. In particular, the detection of diffuse synchrotron radiation from radio halos or relics provides evidence for the existence of magnetic fields within the intra-cluster medium (ICM) \\citep[for a review, see][]{2002ARA&A..40..319C}. Since the detection rate of radio halos in galaxy clusters seems to be of the order of 30\\% for X-ray luminous clusters \\citep{1999NewA....4..141G}, the presence of magnetic fields appears to be common.  Based on these observations, \\citet{2002A&A...396...83E} developed a redshift dependent radio halo luminosity function and predicted large numbers of radio halos to be detected with future radio telescopes.  A different piece of evidence comes from Faraday rotation which arises owing to the birefringent property of magnetised plasma causing the plane of polarisation to rotate for a nonzero magnetic field component along the propagation direction of the photons \\citep{2001ApJ...547L.111C}.  However, the accessible finite windows given by the extent of the sources emitting polarised radiation are a limitation of this method.  The derived magnetic field strengths depend on the unknown magnetic field autocorrelation length which has to be deprojected from the observed two dimensional Faraday rotation measure maps using certain assumptions \\citep[see, however,][]{2003A&A...401..835E, 2003A&A...412..373V}.  A different approach is given by the energy equipartition argument if a particular cluster exhibits diffuse radio synchrotron emission.  The method assumes equal energy densities of cosmic ray electrons and magnetic fields in order to estimate volume averaged magnetic field strengths.  The minimum energy criterion is a complementary method.  It is based on the idea of minimising the non-thermal energy density contained in cosmic ray electrons (CRe), protons (CRp), and magnetic fields by varying the magnetic field strength. As one boundary condition, the implied synchrotron emissivity is required to match the observed value. Additionally, a second boundary condition is required mathematically which couples CRp and CRe.  For the classical case, a constant scaling factor between CRp and CRe energy densities is assumed. However, if the physical connection between CRp and CRe is known or assumed, a physically better motivated criterion can be formulated.  As such a case, we introduce the minimum energy criterion within the scenario of hadronically generated CRe.  Classically, the equipartition/minimum energy formulae use a fixed interval in radio frequency in order to estimate the total energy density in cosmic ray electrons (CRe), a purely observationally motivated procedure \\citep{1956ApJ...124..416B, 1970ranp.book.....P}. However, this approach has a drawback when comparing different field strengths between galaxy clusters because a given frequency interval corresponds to different CRe energy intervals depending on the magnetic field strengths \\citep{2001SSRv...99..243B}.  For this reason, variants of the minimum energy criterion have been studied in order to place the magnetic field estimates on more physical grounds, based then on assumptions such as the fixed interval in CRe energy \\citep{1993A&A...270...91P, 1996ARA&A..34..155B, 1997A&A...325..898B}. The modified classical minimum energy criterion does not specify a particular energy reservoir of the CRe.  However, this apparent model-independence is bought dearly at the cost of the inferred magnetic field strength depending on unknown parameters like the lower energy cutoff of the CRe population or the unknown contribution of CRp to the non-thermal energy density. In the following, we use the term {\\em classical minimum energy criterion} in its modified version, including e.g.~a fixed interval in CRe energy.  Natural candidates for acceleration mechanisms providing a highly-relativistic particle population are strong structure formation and merger shocks \\citep[e.g.,][]{1980A&AS...39..215H, 1999ApJ...520..529S} or reacceleration processes of 'mildly' relativistic CRe ($\\gamma_\\e\\simeq 100-300$) being injected over cosmological timescales into the ICM. Owing to their long lifetimes of a few times $10^9$ years, these mildly relativistic CRe can accumulate within the ICM \\citep[see][ and references therein]{2002mpgc.book....1S}, until they experience continuous in-situ acceleration via resonant pitch angle scattering by turbulent Alfv\\'en waves as originally proposed by \\citet{1977ApJ...212....1J} and reconsidered by \\citet{1987A&A...182...21S}, \\citet{2001MNRAS.320..365B}, \\citet{2002ApJ...577..658O}, \\citet{2002A&A...386..456G}, and \\citet{2003ApJ...594..732K}. However, this reacceleration scenario also faces challenges as recent results imply: \\citet{2003astro.ph.12482B} show, that if the CRp-to-thermal energy density ratio were more than a few percent, Alfv\\'en waves would be damped efficiently such that the reacceleration mechanism of the electrons is inefficient.  Because nearly all conceivable electron acceleration mechanisms produce a population of CRp which accumulates within the clusters volume, this represents an efficient damping source of Alfv\\'en waves.\\footnote{Indeed, there are first hints for the existence of a 10~MeV - 100~MeV CRp population deriving from the detection of excited gamma-ray lines from the clusters Coma and Virgo \\citep{2004A&A.413..817I}. If verified, that would make a high energy (GeV) CRp population very plausible.} \\citet{2004ApJ...604..108K} presented an interesting line of argumentation to investigate the nature of radio halos by comparing the observed and statistically predicted population. This approach might allow to measure the life time of radio halos and thus help to conclude their physical origin with a future flux-limited, controlled, and homogeneous radio halo sample.  In this work, we examine a minimum energy criterion within another specific model for the observed extended radio halos of $\\sim$~Mpc size: Hadronic interactions of CRp with the ambient thermal gas produce secondary electrons, neutrinos, and $\\gamma$-rays by inelastic collisions taking place throughout the cluster volume.  These secondary CRe would generate radio halos through synchrotron emission \\citep{1980ApJ...239L..93D, 1982AJ.....87.1266V, 1999APh....12..169B, 2000A&A...362..151D, 2001ApJ...559...59M, 2004A&A...413...17P}. This scenario is motivated by the following argument: The radiative lifetime of a CRe population in the ICM, generated by direct shock acceleration, is of the order of $10^8$ years for $\\gamma_\\e \\sim 10^4$. This is relatively short compared to the required diffusion timescale needed to account for such extended radio phenomena \\citep{2002BrunettiTaiwan}.  On the other hand, the CRp are characterised by lifetimes of the order of the Hubble time, which is long enough to diffuse away from the production site and to be distributed throughout the cluster volume to which they are confined by magnetic fields \\citep{1996SSRv...75..279V, 1997ApJ...477..560E, 1997ApJ...487..529B}.  The magnetic field strength within this scenario is obtained by analogy with the classical minimum energy criterion while combining the CRp and CRe energy densities through their physically connecting process. Apart from relying on the particular model, the inferred magnetic field strengths do not depend strongly on unknown parameters in this model.\\footnote{Likewise the minimum energy criterion within the reacceleration scenario of mildly relativistic CRe ($\\gamma_\\e\\simeq 100-300$) can be obtained by minimising the sum of magnetic, mildly relativistic CRe, and turbulent energy densities while allowing for constant synchrotron emission.}  The philosophy of these approaches is to provide a criterion for the energetically least expensive radio synchrotron emission model possible for a given physically motivated scenario.  To our knowledge, there is no first principle enforcing this minimum state to be realised in Nature. However, our minimum energy estimates are interesting in two respects: First, these estimates allow scrutinising the hadronic model for extended radio synchrotron emission in clusters of galaxies.  If it turns out that the required minimum non-thermal energy densities are too large compared to the thermal energy density, the hadronic scenario will become implausible to account for the extended diffuse radio emission. For the classical minimum energy estimate, such a comparison can yield constraints on the accessible parameter space spanned by lower energy cutoff of the CRe population or the contribution of CRp to the non-thermal energy density.  Secondly, should the hadronic scenario be confirmed, the minimum energy estimates allow testing for the realisation of the minimum energy state for a given independent measurement of the magnetic field strength.  This article is organised as follows: After introducing synchrotron radiation of CRe (Sect.~\\ref{sec:CRe}), analytic formulae for hadronically induced emission processes are presented (Sect.~\\ref{sec:CRp}). The classical and hadronic minimum energy criteria are then derived, the theoretically expected tolerance regions are given, and limiting cases are discussed (Sect.~\\ref{sec:MEC}). In Sect.~\\ref{sec:testing}, we examine whether future observations of inverse Compton emission and hadronically induced $\\gamma$-ray emission can serve as tests for the verification of the minimum energy criterion. Magnetic and cosmic ray energy densities and their tolerance regions are inferred from application of the minimum energy arguments in the Coma and Perseus cluster for both scenarios (Sect.~\\ref{sec:applications}).  Our article concludes with formulae which provide recipes for estimating the magnetic field strength in typical observational situations (Sect.~\\ref{sec:nutshell}).  Throughout this paper we use the present Hubble constant $H_0 = 70~h_{70}\\mbox{ km s}^{-1} \\mbox{ Mpc}^{-1}$, where $h_{70}$ indicates the scaling of $H_0$. ",
        "conclusions": "We investigated the minimum energy criterion of radio synchrotron emission in order to estimate the energy density of magnetic fields with the main focus on the underlying physical scenario.  The classical scenario might find application for cosmic ray electrons (CRe) originating either from primary shock acceleration or in-situ reacceleration processes while the hadronic model assumes a scenario of inelastic cosmic ray proton (CRp) interactions with the ambient gas of the intra-cluster medium (ICM) and thus leads to extended diffuse synchrotron and $\\gamma$-ray emission.  Generally, the hadronic minimum energy estimates allow testing the hadronic model for extended radio synchrotron emission in clusters of galaxies.  If it turns out that the required minimum non-thermal energy densities are too large compared to the thermal energy density, the hadronic scenario has to face serious challenges.  For the classical minimum energy estimate, such a comparison can yield constraints on the accessible parameter space spanned by the lower energy cutoff of the CRe population or the unknown contribution of CRp to the non-thermal energy density.  For the first time we examine the localisation of the predicted minimum energy densities and provide a measure of the theoretically expected tolerance regions of these energetically favoured energy densities. The tolerance regions of the particular energy densities inferred from the classical minimum energy criterion are approximately constant for varying magnetic field strength $\\eps_B$. On the contrary, the hadronic minimum energy criterion predicts constant energy densities for varying magnetic field strength in the case of low $\\eps_B$ compared to $\\eps_\\rmn{CMB}$, while the tolerance region of the CRp energy density decreases at the same rate as the tolerance region of $\\eps_B$ increases for high $\\eps_B$.  Future observations should shed light on the hypothetical realisation of such an optimal distribution of energy densities in Nature: Combining upper limits on the inverse Compton (IC) scattering of cosmic microwave background photons off CRe within the ICM provides lower limits on the magnetic field strength.  Unambiguous detection of the $\\pi^0$-decay induced $\\gamma$-ray emission owing to hadronic CRp interactions in the ICM together with the observed radio synchrotron emission yields strong upper limits on the magnetic field strength. These are only upper limits because the inevitably accompanying hadronically generated CRe could have a non-hadronic counterpart CRe population which also contributes to the observed synchrotron emission. A combination of IC detection in hard X-rays, radio synchrotron emission, and hadronically induced $\\gamma$-ray emission therefore simultaneously enables the determination of the CRp population as well as a bracketing of the total magnetic field strength and the CRe population. Applying the appropriate minium energy arguments would yield information about both the dynamical state as well as the fragmentation of the spatial distribution of the magnetic field.  Requiring the sum of cosmic ray and the magnetic field energy densities to be minimal for the observed synchrotron emission of the radio halo of the Coma cluster and the radio mini-halo of the Perseus cluster yields interesting results: Within the theoretically expected tolerance regions, equipartition is possible between the energy densities of CRp and magnetic fields, i.e.~the minimum energy criterion always seems to choose equipartition to be a quasi-optimal case.  Applying the hadronic minimum energy criterion to the diffuse synchrotron emission of the Coma cluster yields a central magnetic field strength of $B_\\rmn{Coma} = 2.4^{+1.7}_{-1.0}~\\umu\\rmn{G}$ while in the case of the cool core cluster Perseus we obtain $B_\\rmn{Perseus} = 8.8^{+13.8}_{-5.4}~\\umu\\rmn{G}$.  These values agree with magnetic field strengths inferred from Faraday rotation which range in the case of clusters without cool cores within $[3~\\umu\\rmn{G}, 6~\\umu\\rmn{G}]$ while cool core clusters yield values of $\\sim 12~\\umu\\rmn{G}$ \\citep{2003A&A...412..373V}. Within the hadronic model for the radio mini-halo in the Perseus cluster, this results in a confinement for the CRp energy density of $2\\%\\pm 1\\%$ of the thermal energy density while the magnetic energy density reaches only 0.4\\% of the thermal energy density within large uncertainties.  These energetic considerations show that the hadronic scenario is a very attractive explanation of cluster radio mini-halos.  In order to account for the radio halo of Coma in the hadronic scenario, the product of $\\eps_\\crp$ and $\\eps_B$ needs to increase by nearly two orders of magnitude relative to the square of the thermal energy density $\\eps_\\rmn{th}$ towards the outskirts of the halo.  Moving away from the minimum energy solution and adopting for instance a constant magnetic-to-thermal energy density, it is energetically possible to explain the observed synchrotron emission hadronically by only requiring the magnetic and CRp energy density to be a few per cent relative to the thermal energy density (and even less for the CRp in the cluster centre, provided $\\alpha_\\p \\sim 2.3$ and the cluster is isothermal). Such a magnetic energy density corresponds to a central magnetic field strength of $6~\\umu$G. Assuming a lower magnetic field strength of $3~\\umu$G corresponding to a magnetic-to-thermal energy density of approximately 0.5\\% requires the CRp energy density to be lower than 10\\% for the entire range of the radio halo.  The considered hadronic scenario assumes a CRp spectral index which is independent of position and thus the radio emission does not show any spatial variations over the clusters volume. In principle, one could allow for radial spectral variations of the CRp and thereby for the radio emission by adopting a particular history of this population.  For instance, one possible scenario would be given by continuous in-situ acceleration of CRp via resonant pitch angle scattering by turbulent Alfv\\'en waves. We discuss that a moderate radial steepening would not significantly modify the hadronic minimum energy condition while a confirmation of the strong steepening reported by \\citet{1993ApJ...406..399G} would seriously challenge the hadronic scenario.  As a caveat, it should be stressed that the inferred values for the particular energy densities only represent the energetically least expensive radio synchrotron emission model possible for a given physically motivated scenario. This minimum is not necessarily realised in Nature. Nevertheless, our minimum energy estimates are also interesting in a dynamical respect: Should the hadronic scenario of extended radio synchrotron emission be confirmed, the minimum energy estimates allow testing for the realisation of the minimum energy state for a given independent measurement of the magnetic field strength.  Within the tolerance regions, our minimum energy estimates in Perseus and Coma agree well with magnetic field strengths inferred from Faraday rotation.  Under the hypotheses of correctness of the hadronic scenario, such a possible realisation of the minimum energy state would seek an explanation of a first principle enforcing this extremal value to be realised in Nature."
    },
    "0404/astro-ph0404605_arXiv.txt": {
        "abstract": "{ LS~I~+61$^{\\circ}$303 is one of the most observed Be/X-ray binary systems because,  peculiarly, it has  periodical radio and X-ray emission together with  strong,  variable gamma-ray emission. This source remains,  however, quite enigmatic. Some properties of this system can be   explained assuming that the unseen companion is a non-accreting young pulsar with a  relativistic wind strongly interacting with  the wind of  the Be star. On the contrary, other properties of LS~I~+61$^{\\circ}$303 fit a model where the   companion  is accreting even with two events of  super-critical accretion  along the orbit. The very recent discovery  of a radio jet extending ca. 200 AU at both sides of a central core  has definitely proved the occurrence of accretion--ejection processes  in this system. Therefore it is of  great interest to combine  this result with previous observations  at other wavelengths within the framework  of the two-peak   accretion  (ejection) model. Concerning the first ejection, we show that the observed  gamma-rays variations   might be  periodic with outbursts   confined around the periastron passage (i.e. where the first accretion-rate peak occurs and high-energy emission but no radio emission is predicted). Concerning the second ejection, with  radio bursts, we point out that  it can be also traced  in the X-ray data, both  in episodes of hardening of the X-ray emission and  in  a transition from soft- to hard-states at the onsets of radiobursts. In fact, both hardening and  transitions between spectral states are related to the dramatic  change in the structure of the accretion disk  preceeding the ejection. Finally, we explore the nature of the accretor and  we conclude that on the basis of the present  optical data a black hole cannot be  ruled out. ",
        "introduction": "\\label{introduction}   The X-ray binary stellar system  LS~I~+61$^{\\circ}$303 has always  excited  particular interest because of its two peculiarities: On the one hand being the probable counterpart of the  variable gamma ray source \\object{2CG~135+01} (Gregory $\\&$ Taylor \\cite{gregory78}; Tavani et al. \\cite{tavani98}) and on the other hand  being a periodic radio source (Taylor \\& Gregory \\cite{taylor82}, \\cite{taylor84}).  The binary system   is composed of a compact object in an eccentric orbit around  a rapidly rotating B0-B0.5 main-sequence star, undergoing mass loss through an equatorial disk (Hutchings \\& Crampton \\cite{hutchings81}). The orbital period is assumed to be of 26.496 days determined on a radio data base of over 20 years (Hutchings \\& Crampton \\cite{hutchings81}; Gregory et~al. \\cite{gregory99}; Gregory \\cite{gregory02}).  High radio outbursts always peak near phase 0.6, where  $\\Phi$=0  has been set at Julian Date 2443366.775 (Taylor \\& Gregory \\cite{taylor82}; Paredes et al. \\cite{paredes90}). However, lower intensity radio outbursts occur within the   broad  distribution $\\Phi_{radio}$=0.45--0.95, due to variations of the equatorial disk in the Be star (Paredes et al. \\cite{paredes90}; Gregory et al. \\cite{gregory99}; Zamanov \\& Mart\\'i \\cite{zamanov00}; Gregory et al. \\cite{gregory02}; Gregory \\& Neish \\cite{gregory02b}). From optical and near-infrared observations  the periastron passage is estimated to be in the range $\\Phi_{periastron}$=0.2--0.5 (Hutchings \\& Crampton \\cite{hutchings81}; Mart\\'i \\& Paredes \\cite{marti95}). One of the fundamental questions concerning the periodic radio outbursts of \\object{LS~I~+61$^{\\circ}$303} therefore has been: Why are the radio outbursts shifted with respect to the periastron passage?  The X-ray emission is also periodic with period estimates of P= 26.7 $\\pm$ 0.2 by  Paredes et al. (\\cite{paredes97}) and    P=26.42$\\pm$ 0.05 by  Leahy (\\cite{leahy01}), clearly in agreement with the radio period. Quite surprising, however, is  that the X-ray outbursts are offset from the radio ones. In fact, the  two available  simultaneous X-ray and radio observations by Taylor et al.(\\cite{taylor96}) and by Harrison et al.  (\\cite{harrison00}) during an orbital period show that the  X-ray emission peaks in the phase interval $\\Phi_{X-ray}$=0.43--0.47 (recalculated  by Gregory \\cite{gregory02}) while  the radio outburst is offset to that by several days.  Taylor et al.  (\\cite{taylor92}) and Mart\\'{\\i} \\& Paredes (\\cite{marti95}) have modelled the properties of this system in terms of an accretion rate $\\dot{M} \\propto {\\rho_{wind}\\over v_{rel}^3}$, (where $\\rho_{wind}$ is the density of the Be star wind and $v_{rel}$ is the relative speed between the accretor and wind) which develops two peaks for  high eccentricities: the highest peak corresponds to the periastron passage (highest density), while the second peak occurs when the drop in the relative velocity $v_{rel}$ compensates (because of the inverse cube dependence) the decrease in density. From typical parameters of LS~I~+61$^{\\circ}$303 derived from near infrared data, Mart\\'{\\i} \\& Paredes (\\cite{marti95}) have shown that both peaks are above the Eddington limit and therefore one expects that matter is ejected perpendicular to the plane of the accretion disk. Near the periastron the short distance to the Be star enhances inverse Compton losses: X-ray or/and gamma-ray outbursts are expected but no radio bursts. At the second accretion peak, the compact object is much farther away from the Be star, so that the electrons can propagate out of the orbital plane:  an expanding double radio source should be observed with a radio interferometer (Taylor \\& Gregory  \\cite{taylor84}; Taylor et~al. \\cite{taylor92}; Massi et al. \\cite{massi93};  Mart\\'{\\i} \\& Paredes \\cite{marti95}; Gregory et~al. \\cite{gregory99}).  The main problem concerning this model is  that the luminosity of \\object{LS~I~+61$^{\\circ}$303} in the X-ray range is only $L_{\\rm X}\\simeq 10^{35}$~erg~s$^{-1}$ (Maraschi \\& Treves \\cite{maraschi81}). That is three orders of magnitude lower than the Eddington limit, even for  a neutron star, and in addition the bulk of the energy output seems to be shifted from X-ray to $\\gamma$-ray wavelengths  with  $L_{\\rm \\gamma}\\simeq10^{37}$erg~s$^{-1}$ (Hartman et~al. \\cite{hartman99}). The difficulty of interpreting these results in the context of a super-accretion model has led to an alternative young pulsar model where a population of  relativistic electrons are produced at the shock boundary between the relativistic wind of the young  pulsar and the wind of the Be star ( Maraschi \\& Treves \\cite{maraschi81}; Tavani \\cite{tavani95}; Harrison et al. \\cite{harrison00}; Hall et al. \\cite{hall03}). As a matter of fact  such a model  fits the time-variable high-energy emission observed near periastron from the Be/pulsar system PSR B1259-63  (Tavani \\& Aaron \\cite{tavani97}).  However, the recent discovery  of a radio emitting  jet extending ca. 200 AU at both sides of a central core (Massi et al. \\cite{massi04})  has shown the occurrence of accretion--ejection processes  in LS~I~+61$^{\\circ}$303. Therefore, this source seems  very similar to  the microquasar LS~5039 (Paredes et~al. \\cite{paredes00}) also subluminous in the X-ray range (even more than \\object{LS~I~+61$^{\\circ}$303}) and also having $L_{\\gamma} > L_{\\rm X}$. The quite stable gamma-ray emission in that case is explained due to upscattered stellar photons via inverse Compton from the relativistic electrons of the persistent  jet. If that is true for \\object{LS~I~+61$^{\\circ}$303} (Taylor et al.  \\cite{taylor96}; Massi et~al. \\cite{massi01}; Kaufman Bernad\\'o et~al. \\cite{kaufman02}), do the periodic outbursts of this source  imply  periodic gamma-ray bursts ? And in this case at which  $\\Phi_{\\gamma-ray}$ ? Moreover, recent developments in the theory of the accretion-ejection processes show that  magneto-rotational instabilities are able to accelerate and to collimate a part of the disk material into a double jet only after thermal instabilities in the accretion disk have inflated and transformed  it into a geometrically thick disk (Meier \\cite{meier01}). As shown by  Belloni et al. (\\cite{belloni97}) structural  changes of the disk are associated with  changes in the X-ray spectral states. Is it possible to discern a  change of state in available X-ray data ?  And in this case: are   changes of state and radio outburst related to each other as  expected in an accretion-ejection process ? The aim of this paper is to try to answer these questions. Section  2 analyses and discusses gamma-ray data while  Section 3 deals with  X-ray and optical results. The conclusions are given in Section 4. ",
        "conclusions": "\\label{discussion}  The  emission of  the gamma-ray source 2CG~135+01 is consistent with a periodic behaviour, with a period  similar to  the orbital/radio period of the system LS~I~+61$^{\\circ}$303. The  gamma-ray peaks, observed at two different epochs (Tavani et al. \\cite{tavani98}), remain   confined  around the periastron passage, a fact which may imply that the most of the seed photons for  Comptonization are indeed  stellar photons (Taylor et~al. \\cite{taylor96}).  The  X-ray observations along an orbital cycle (Greiner \\& Rau \\cite{greiner01}) show that the system remains always in a  high/soft state, characterized by a photon index $\\Gamma \\ge 2$. Only at the onset of radio emission there happens a  state transition, characterized by $\\Gamma  \\sim 1.6$,  to   the low/hard state  (Tanaka \\cite{tanaka97}, Fender et al. \\cite{fender99}). These  transitions between spectral states are related to the dramatic  change in the structure of the accretion disk  preceeding the ejection of part of it into a jet (Belloni et al. \\cite{belloni97}, Mirabel et al. \\cite{mirabel98}, Fender et al. \\cite{fender99}). Good sampled simultaneous X-ray and radio observations in different outbursts would be quite important  to trace (by monitoring the photon index $\\Gamma$ ) the two  accretion-ejection processes occurring  in this peculiar source.   Finally, we  suggest that LS~I~+61$^{\\circ}$30 may host a black hole, because   its high/soft state has  a power law component with $\\Gamma \\ge$2, which  is typical for systems hosting black holes. On the other hand the value of M$_x$=1.5 M$\\odot$ - always quoted in the literature - is related to average values for the  inclination of the orbit, mass of the Be star  and mass function. The ranges for these parameters  presently available are so large that a massive black hole in LS~I~+61$^{\\circ}$303 cannot  be excluded. A new determination of all these values is desirable."
    },
    "0404/astro-ph0404449_arXiv.txt": {
        "abstract": "Recently we have made measurements of thermonuclear burst energetics and recurrence times which are unprecedented in their precision, largely thanks to the sensitivity of the {\\it Rossi X-ray Timing Explorer}\\/ ({\\it RXTE}). In the \"Clocked Burster\", GS~1826$-$24, hydrogen burns during the burst via the rapid-proton (rp) process, which has received particular attention in recent years through theoretical and modelling studies. The burst energies and the measured variation of alpha (the ratio of persistent to burst flux) with accretion rate strongly suggests solar metallicity in the neutron star atmosphere, although this is not consistent with the corresponding variation of the recurrence time. Possible explanations include extra heating between the bursts, or a change in the fraction of the neutron star over which accretion takes place. I also present results from 4U~1746$-$37, which exhibits regular burst trains which are interrupted by ``out of phase'' bursts. ",
        "introduction": "Unstable thermonuclear ignition of accreted fuel on neutron stars (NSs) in low-mass X-ray binaries (LMXBs) is triggered once a critical column density is reached in the fuel layer (e.g. \\cite[]{bil00}).  Regular bursting is surprisingly uncommon, and only one source is known to consistently burst regularly (GS~1826$-$24; \\cite[]{gal03d}).  The conditions required for regular bursting likely include steady accretion and uniform spreading of the accreted fuel over the NS surface, as well as complete fuel consumption.  If the accretion rate is sufficiently steady the critical density for ignition will be reached after a fixed time. If, additionally, all the accreted fuel is burnt during each burst, then each burst will be ignited after the same interval, leading to regular bursting. Clearly, relaxation of any one of these conditions will lead to variations in the burst interval, and deviations away from regular bursting.  Here we present recent results obtained via measurement of thermonuclear burst properties in GS~1826$-$24 and the globular cluster source 4U~1746$-$37 with \\xte. In GS~1826$-$24 we found regular bursting at a range of accretion rates, which allows us to constrain the composition of the burning fuel. In 4U~1746$-$37 we found trains of regular bursts interrupted by bursts which were ``out of phase''. We discuss possible mechanisms for this phenomenon, as well as future observational tests to distinguish between them. ",
        "conclusions": "The precise measurements of burst properties possible with \\xte\\/ allow tests of burst theory to unprecedented levels of precision. For both sources discussed here, the bursting behaviour may be significantly altered because accretion is not taking place over the entire NS surface. Measurements of blackbody radii from burst spectra offer at best a qualitative way to measure the area of accretion, unless the deviations of the spectra from pure blackbodies can be accounted for and the effective temperature measured accurately. This can likely only be achieved with measurements by dedicated spectroscopic instruments like {\\it Chandra}\\/ or {\\it XMM-Newton}.  Such studies, when combined with the extensive archival observations of bursters accumulated by \\xte\\/ over its lifetime are thus an excellent way to improve our understanding of burst physics. Future observations in previously unexplored ranges of accretion rate will also further constrain the burst physics."
    },
    "0404/astro-ph0404480_arXiv.txt": {
        "abstract": "The Cyclic Model attempts to resolve the homogeneity, isotropy, and flatness problems and  generate a nearly scale-invariant spectrum of fluctuations during a period of slow contraction that precedes a bounce to an expanding phase. Here we describe at a conceptual level the recent developments  that have greatly simplified our understanding of the contraction phase and the Cyclic Model  overall. The answers to many past questions and criticisms are now understood. In particular, we show that the contraction phase has equation of state $w>1$ and that contraction with $w>1$ has a surprisingly similar properties to inflation with $w < -1/3$.  At one stroke, this shows how the model is different from inflation and why it may work just as well as inflation in resolving cosmological problems. ",
        "introduction": "Two years ago, the Cyclic Model \\cite{cyclic} was introduced as a radical alternative to the standard big bang/inflationary picture \\cite{inflation}. Its purpose is to offer  a new solution to the homogeneity, isotropy, flatness problems and a new mechanism for generating a nearly scale-invariant spectrum of fluctuations.  One might ask why we should consider an alternative when inflation \\cite{inflation,Bardeen,inflapert} has scored so many successes in explaining a wealth of new, highly precise data \\cite{MAPspergel}.   There are several reasons.  First, seeking an alternative is just plain good science.  Science proceeds most rapidly when there are two or more competing ideas.  The ideas focus attention on what are the unresolved issues theorists must address and what are the important measurements experimentalists must perform. Inflation has had no serious competition for several years, and the result has been that its flaws have been ignored.  Many cosmologists are prepared to declare inflation to be established even though crucial experimental tests remain.  Competition stimulates critical thinking and removes complacency.  A second reason to consider an alternative is that, even though inflationary predictions are in marvelous accord with the data thus far, the theoretical front has seen little progress.  In fact, if anything, there has been retrogress.  The main questions about inflation that were  cited twenty years ago remain today.  What is the inflaton and why are its interactions finely-tuned?  How did the universe begin and why did it start to inflate?   With the advent of string theory, these issues have become severe problems. Despite heroic efforts to construct stringy inflation models with tens or hundreds of moving parts (fluxes, branes and anti-branes) and examining a complex landscape of (at least) $10^{500}$ vacua, even a single successful inflationary model is  difficult to construct \\cite{Malda}.  The notion that there is a landscape of $10^{500}$ or more string vacua has suggested to some that, if there is  an acceptable vacuum somewhere, inflation makes it possible to populate all vacua; and that the ultimate explanation for our universe is anthropic \\cite{Susskind}. However, this cannot be the whole story since it begs the question of how the universe started in the first place.  No matter where you lie in the landscape, extrapolating back in time brings you to a cosmic singularity in a finite time. The issue of the beginning remains  unresolved.  Furthermore, relying on the anthropic principle is like stepping on quicksand.  The power of a theory is measured by the ratio of its explanations/predictions to assumptions.  A good scientific theory is observationally testable.  An anthropic explanation is based upon considerations involving regions of space that are causally disconnected from us and that will, in many cases, never be observed by us. What parameters and properties can vary from region to region?  What is the probability distribution?  In models such as eternal inflation,  the relative likelihood of our being in one region or another is ill-defined since there is no unique time slicing and, therefore, no unique way of assessing the number of regions or their volumes. Brave souls have begun to head down this path, but it seems likely to us to drag a beautiful science towards the darkest depths of metaphysics.  Another unresolved issue is trans-Planckian effects on the production of density perturbations \\cite{trans}. In inflationary cosmology, the fluctuations observed in the cosmic microwave background had wavelengths at the beginning of inflation that were smaller than the Planck scale. The standard approximation is to assume the initial distribution of sub-horizon and, hence, sub-Planckian fluctuations corresponds to quantum fluctuations on an empty, Minkowski background.  However, quantum gravity effects may cause the distribution to be different on sub-Planckian wavelengths.  The unknown distortion would be inflated and produce an uncertain correction to inflationary predictions for the cosmic microwave background anisotropy.   Finally, the big bang/inflationary picture is still reeling from the recent shock that most of the universe consists of dark energy \\cite{SN}.  The concept had been that, once conditions are set in the early universe, the rest of cosmic evolution is simple.  Dark energy has shattered that dream.  Dark energy was not anticipated and plays no significant role in the theory.  Observations have forced us to add dark energy {\\it ad hoc.}  The current approach in big bang/inflationary model-building has been to treat the key issues --  the bang, the creation of homogeneity and density fluctuations, and dark energy --  in a modular way. Separate solutions with separate ingredients are sought for each.  Perhaps this approach will work, all the problems cited above will be resolved, and a simple picture will emerge.  But, perhaps the time has come to consider a different, holistic approach.  The cyclic model has an ambitious manifesto. Its goal is  to address the entire history of the universe, past and future, in an efficient, unified approach.  There is one essential ingredient -- branes in the higher-dimensional picture or a scalar field in the four-dimensional effective theory --  that is simultaneously responsible for explaining the big bang; the solution to the homogeneity, isotropy, flatness,  and monopole problems \\cite{cyclic}; the generation of nearly scale-invariant fluctuations that seed large-scale structure \\cite{ekperts,newtolley};  and, the source of dark energy \\cite{cyclic}.  Simplicity and parsimony are essential elements.  The range of acceptable parameters is broad \\cite{design}.  Over the past two years, the Cyclic Model has progressed  remarkably.  The concept has been examined by numerous groups, and many, many useful criticisms and questions have been raised \\cite{lyth,linde,contrascale,durrer,ambig,string}. As we and our collaborators have tried to address these issues, the results have been interesting. First, we have discovered that the Cyclic Model already contained the answers.  Not a single new ingredient has had to be added thus far.  Rather, we have learned to recognize fully the physical properties of the components the model contained at the outset \\cite{reply,bcbb,duality,dual2,chaos}.    That is, we have been discovering new physical principles stemming from the original model rather than adding new ingredients and patches.  Second, as we have come to understand the Cyclic Model better, the picture has become much, much simpler.  We believe we can stick by our manifesto: If the model is going to work, it will be because of basic ideas as simple and compelling as inflation.  In fact, we find that there are remarkable, unanticipated parallels between inflationary expansion and the contracting and bounce phases of the Cyclic Model \\cite{duality,dual2}. There remain important open issues about the bounce itself, but, now we can confidently say that many of the issues that plagued previous attempts at contracting cosmological models have been cleared away and there are solid reasons for optimism about resolving the remaining issues.   The purpose of this essay is to present the simplified view of the Cyclic Model, focusing on the stages that are most novel and controversial: the contraction and bounce. We focus on the two key ingredients needed to understand the contracting phase: branes and the equation of state $w>1$.  As we explain, the two features lead to a series of novel physical effects that  solve the homogeneity, isotropy, and flatness problems and ensure a nearly scale-invariant spectrum of density perturbations following the big bang. ",
        "conclusions": ""
    },
    "0404/astro-ph0404163_arXiv.txt": {
        "abstract": "{  We present N-body/SPH simulations of the evolution of an isolated dwarf galaxy including a detailed model for the ISM, star formation and stellar feedback. Depending on the strength of the feedback, the modelled dwarf galaxy shows periodic or quasi-periodic bursts of star formation of moderate strength. The period of the variations is related to the dynamical timescale, of the order of $1.5~10^8$ yr. We show that the results of these simulations are in good agreement with recent detailed observations of dwarf irregulars (dIrr) and that the peculiar kinematic and morphological properties of these objects, as revealed by high resolution HI studies, are fully reproduced. We discuss these results in the context of recent surveys of dwarf galaxies and point out that if the star formation pattern of our model galaxy is typical for dwarf irregulars this could explain the scatter of observed properties of dwarf galaxies. Specifically, we show that the time sampled distribution of the ratio between the instanteneous star formation rate (SFR) and the mean SFR is similar to that distribution in observed sample of dwarf galaxies. ",
        "introduction": "The nature of the processes regulating star formation in irregular galaxies is poorly understood. Whereas there is at least some understanding of star formation in regular spiral galaxies, this is less so for irregulars. For spiral galaxies the guiding observation that the star formation rate (SFR) is related to the gas surface density by the Schmidt law has given rise to a number of competing theories that reproduce the general features of star formation in large spiral galaxies (Elmegreen~\\citeyear{E02}, Dopita \\& Ryder~\\citeyear{DR94}). These systems seem to be regulated by large scale gravitational instabilities. Star formation in irregular galaxies has proven to be more difficult to understand. Irregulars have a widely varying SFR, spanning 4 orders of magnitude for the normalized SFR/area (Hunter~\\citeyear{H97}), possibly due to the fact that gas thermodynamics, governed by varying heating and cooling processes, plays the decisive role (Elmegreen~\\citeyear{E02}). But why do some irregulars have very high SFRs relative to their mass, while others hardly show any activity? Are there any intrinsic properties of the galaxies that can explain this disparity between SFRs or do all dwarf galaxies exhibit episodes of high star formation?  In recent years a number of studies have highlighted these questions by investigating samples of dwarf galaxies and comparing their properties as derived from photometry, HI and H$\\alpha$ observations. Van Zee (\\citeyear{Z00}, \\citeyear{Z01}) investigated a sample of isolated dwarf galaxies and found no strong correlation between star formation and independent physical parameters. Hunter et al. (\\citeyear{HEB98}) tested different regulating processes, amongst which disk instabilities, thermal and shear regulated star formation, but found that none could explain patterns of star formation. Stil (\\citeyear{S99}) investigated the relation between star formation and HI gas kinematics.  On the other hand, detailed studies of a number of nearby dwarf galaxies have fully revealed the complex structure of the ISM in these systems. High resolution aperture synthesis mapping (e.g. Kim et al.~\\citeyear{KSSD98}, Wilcots \\& Miller~\\citeyear{WM98}, Puche et al.~\\citeyear{PWBR92}, Walter \\& Brinks~\\citeyear{WB01}) of their HI has shown the interstellar medium (ISM) of these dwarfs to be a frothy structure, with holes of varying sizes, shells and filaments, even extending far beyond the optical radius. From velocity dispersion studies (Young et  al.~\\citeyear{YL97}) the presence of cold and warm neutral components predicted by the two phase model for the ISM (Field~\\citeyear{F65}) has been deduced. Comparison with H$\\alpha$ and UV observations shows that the dense walls of these holes are the sites of star formation (Walter et al.~\\citeyear{WT01}), and suggest 'chains' of successive star forming sites (Stewart et al.~\\citeyear{SFB00}). The cause of the holes in the HI distribution seems to be the energy input from ionizing radiation, stellar winds and supernovae, although Rhode et al.~(\\citeyear{RSWR99}) and Efremov et al.~(\\citeyear{EEH98}) discuss other possible scenarios.  Together these two types of observations have painted a picture of the complex interaction between star formation and the ISM of these systems that is formidable to capture theoretically. Some early attempts have been made to understand star formation qualitatively by the application of stochastic self propagating star formation (SSPSF) models to dwarf galaxies (Gerola \\& Seiden~\\citeyear{GSS80}, Comins~\\citeyear{C83}). While they probably capture some general characteristics of star formation, they are phenomenological and do not include the underlying physics of the ISM.  Efforts to investigate the influence of star formation on the ISM of dwarf galaxies have mainly concentrated on the effects of large (central) bursts and on questions concerning the ejection of gas and the distribution of metals(e.g. Mac-Low \\& Ferrara~\\citeyear{MF99}, Mori et al.~\\citeyear{MYTN97}). Recently there also  have been some simulations adressing the question of survival of small galaxies (Mori et al.~\\citeyear{MFM02}). Generally these simulations have not tried to set up a self consistent model for the ISM and star formation, but prescribed a certain SFR.  The importance of a good model for the ISM and feedback has been recognized by a number of authors. Andersen \\& Burkert~(\\citeyear{AB00}) formulated an extensive model for the ISM in terms of a phenomenological model for the interstellar clouds. Their model showed self regulation of the SFR and they found moderate fluctuations in SFR. Berczik \\& Hensler~(\\citeyear{BHTS03}) incorporated such a cloud model into a chemodynamical galaxy evolution code. Semelin and Combes~(\\citeyear{SC02}) formulated a model with similar characteristics, representing clouds by sticky particles, but did not apply these to dwarf galaxies. Springel \\& Hernquist~(\\citeyear{SH03}) formulated a subgrid model for the multiphase interstellar medium, producing a quiescent self-regulating ISM. However, relatively little effort has been directed towards resolving the normal evolution of irregular dwarfs and providing the connection with detailed studies of single systems and extensive unbiased samples. Nevertheless dwarf galaxies are good test systems for exploring star formation in galaxies: they are dynamically simple systems in the sense that they do not exhibit spiral density waves or shear. Furthermore their small size means that simulations can follow the various physical processes at finer linear and density scales.\\emph{ As the small scale physics of star formation and feedback presumably do not differ between normal and dwarf galaxies, we can use results obtained from these simulations and apply the same methods to larger systems}. Here we present results of a simulation of the evolution of a normal dwarf irregular galaxy including a detailed model for the ISM, star formation and feedback. The distinguishing characteristics of this work are that the model for the ISM we employ does not explicitly postulate the presence of a two phase medium, rather it forms it as a result of the physics of the model. Furthermore we take special care in formulating a star formation model that is solely based on the Jeans instability, and we formulate a feedback scheme that gives us unambiguous control over the strength of the feedback. We will discuss the results of the simulation both in relation to detailed observations of comparable single systems, as well as in the context of recent surveys of dwarf galaxies. ",
        "conclusions": "The model we employed to calculate the evolution of a single dwarf galaxy is quite extensive and complicated and necessarily contains a number of free parameters, some of which we tried to constrain using observational data or theory( as in the case of the supernova energy), yet some remain only constrained by physical intuition. Therefore, it is worthwhile to summarize again the features of observed dwarf systems successfully reproduced by the model: 1) the morphology and kinematics of the HI distribution, 2) the two phase structure of the ISM, 2) the spatial pattern of star formation, 3) the star formation rate, 4) consistency with the observed distribution of Scalo $b$ parameter. These provide strong and independent checks of our modelling approach as they result from the intricate interplay of the model for the ISM we employed, the scheme for star formation and the method of returning mechanical energy to the ISM. These successes give some confidence that the model captures the essentials of dwarf galaxy evolution and to its predictive power.  Comparing our model with recent simulations of dwarf galaxy evolution( e.g. Berzcik et al.~\\citeyear{BHTS03}, Mayer et al.~\\citeyear{MGC01}, Andersen \\& Burkert~\\citeyear{AB00}, Pasetto et al.~\\citeyear{PCC03}, Mori et al.~\\citeyear{MYTN97},~\\citeyear{MYN99}, Carraro et al.~\\citeyear{CCGL01}) we see that our work differs from previous work in two crucial aspects: 1) the modelling of the physics of the neutral ISM, 2) the treatment of supernova feedback. We follow the evolution of the WNM and CNM of the ISM explicitly, this process being the upper part of a cascade leading down to star formation. Star formation in our model happens in cold( $\\rm T \\la 200~K$) and dense ($\\rm n=1-10~cm^{-3}$) gas. These are realistic sites for star formation as we know that molecular clouds are embedded in neutral envelopes. The further stages of star formation, molecular cloud formation and collapse, are not included as they require prohibitive resolution and the inclusion of additional physics. These processes are only implicit in our star formation recipe, but at least our methods follows collapsing gas to the point that it has experienced a transition to a cold, dense state from which it may be trusted to form stars with rates and efficiencies that are constrained by observations. Some authors have tried to bypass this problem by formulating a phenomenological ISM model in terms of 'sticky' particles representing molecular clouds (Andersen \\& Burkert~\\citeyear{AB00}), sometimes in addition to a smooth SPH component representing the WNM (Semelin \\& Combes~\\citeyear{SC02}, Berczik et al.~\\citeyear{BHTS03}). They also succesfully reproduce a selfregulated ISM, including some effects, like evaporation of molecular clouds that are probably not well represented in our model. In these models stars form from the cloud particles, imposing a Schmidt law, not, as in our model, from the consideration of the the instabilities in the ISM. Furthermore, our model has a consistent representation of the ISM linking phases by physical processes, rather than prescriptions.  Simulations of dwarf galaxy evolution that include supernova feedback have typically included this as a local heating term. This method for implementing supernova feedback is not effective in forming the structures associated with stellar feedback and simulations using it do not show any effect of SN feedback. This is a well understood numerical artefact and some authors have devised methods to prevent the radiative loss of mechanical energy that is the root of the problem (Springel \\& Hernquist~\\citeyear{SH03}, Thacker \\& Couchman~\\citeyear{TC00}, Gerritsen \\& Icke~\\citeyear{GI97}, the problem is also not exclusive to SPH type codes see e.g. Fragile et al.~\\citeyear{FMAL03}). They have not applied their methods to the evolution of dIrrs so a direct comparison is not yet possible. However, the notion that \\emph{the energy input from stellar winds and SN of young stellar clusters is crucial for the understanding the structure and kinematics of the ISM of dIrrs} is borne out by our simulations: if feedback is not included, the ISM stays in a smooth disk with very low random motions( 2-3 km/s). The self propagating mode of star formation is absent in that case. The dwarf galaxy will only poorly resemble a real dIrr.  Chemical enrichment is not yet included in our code. We can estimate the importance of a changing chemical composition for our simulation. The total amount of metals $\\rm M_z$ produced during the simulation is $\\rm M_Z \\approx 0.015 \\Delta M_{\\star} = 1.6 \\times 10^5 \\msun$. This will raise the metallicity  of the galaxy, under the assumption that the metals will be well mixed, with at most $25\\% $, which is not entirely insignificant, although to first approximation the thermal evolution is independent of metallicity as both the cooling and UV heating scale with Z. Note also that metals may be lost from the galaxy in galactic winds. It does mean that we should include chemical enrichment to follow the evolution for longer timescales and especially if we want to investigate extremely metal poor galaxies. In principle the inclusion of chemical enrichment will also add further constrains comparing with observations, very important to asses the long term evolution of the model. Both these points, however, do not alter the conclusion for simulations as presented here.  The methods we use can be applied more general. Questions that we plan to address are for example the following: \\begin{itemize} \\item{ In view of the recent realization that there is a deficit of observed small galaxies as compared to predictions of $\\Lambda$CDM models of galaxy formation (Klypin et al.~\\citeyear{KKVP99}) it is interesting to consider what happens for galaxies of progressively smaller mass. We are in the process of running a grid of models exploring this question, but some effects may be clear from the preceding: for smaller galaxies variations will be on longer time scales (scaling as $\\rho_{halo}^{-.5})$ and bigger in amplitude (because the feedback becomes relatively stronger as the escape velocity decreases). Ultimately halos will be too small to retain their ISM, leaving them bare. The mass ranges for which this happens, the influence of other galactic parameters and the timescales on which these processes take place will be of interest to validate cosmological and galaxy formation models. } \\item{ We also plan to investigate a possible relation between dwarf irregulars and dwarf ellipticals. Although recent simulations(Mayer et al.~\\citeyear{MGC01}, Pasetto et al.~\\citeyear{PCC03}) have shown for satellite galaxies that a transition from dIrr to dE or dwarf Spheroidal (dSph) is plausible as a result of the action of tidal fields, it has not been conclusively determined whether a transition from dIrr to dE is to be expected in general. Various people have put forward arguments in favour (Davies \\& Phillipps~\\citeyear{DP88}) and against (Bothun et al.~\\citeyear{BMCM86}, Marlowe et al.~\\citeyear{MMH99}) such a descendancy for the dE. Although the results presented here do not immediately illuminate this question, we think that additional simulations of the sort presented here, testing a wider range of galactic properties and following the evolution for longer time scales may answer whether this is a viable scenario and whether we can put the various classes of small galaxies into an unified evolutionary framework. } \\end{itemize}  In summary, our model suggests that it is possible that a large part of the current dIrr population is in a quasi-periodic burst mode of star formation. The scatter in observed properties in this picture is mainly due to the fact that galaxies are observed at different phases of their evolution. The main difference between our model and classical SSPSF models is that in our models variations are due to the interplay of stellar feedback and gasdynamics, the galaxy being periodically stirred by bursts of star formation after which a quiescent period occurs during which gas falls back to the disk. In classical SSPSF models the variations are due to the fact that star formation is described by a correlated stochastic process and the small size of the galaxies, which induces large statistical variations. Our model predicts that the amplitude of the variations depends on the strength of the feedback and that the period depends on the dynamical time scale."
    },
    "0404/astro-ph0404355_arXiv.txt": {
        "abstract": "{ The structure and the evolution of Pulsar Wind Nebulae (PWNe) are studied by means of two-dimensional axisymmetric relativistic magnetohydrodynamic (RMHD) simulations. After the first imaging of the Crab Nebula with {\\em Chandra}, a growing number of objects has been found to show in the X-rays spatial features such as rings and jets, that clearly cannot be accounted for within the standard framework of one-dimensional semi-analytical models. The most promising explanation suggested so far is based on the combined effects of the latitude dependence of the pulsar wind energy flux, shaping the wind termination shock and naturally providing a higher equatorial emission, and of the wind magnetization, likely responsible for the jet collimation by hoop stresses downstream of the shock. This scenario is investigated here by following the evolution of a PWN interacting with the confining Supernova Remnant (SNR), from the free expansion to the beginning of the reverberation phase. Our results confirm the oblate shape of the wind termination shock and the formation of a polar jet with supersonic velocities ($v\\approx 0.5 - 0.7 c$) for high enough values of the equatorial wind magnetization parameter ($\\sigma\\gsim 0.01$). ",
        "introduction": "Pulsar Wind Nebulae (PWNe, or plerions) arise from the confinement of pulsar winds by the surrounding medium, usually an expanding Supernova Remnant (SNR). The relativistic magnetized pulsar wind is slowed down to non-relativistic velocities at a termination shock, where the magnetic field is compressed and the bulk energy of the outflow is converted into heat and acceleration of particles. These then give rise to the synchrotron and Inverse Compton emission observed from plerions in a very wide range of frequencies, extending from radio wavelengths to X-rays and even $\\gamma$-rays.  The best studied plerion is the Crab Nebula, whose emission has been extensively investigated in all frequency bands and for which most models have been proposed. New light on the spatial structure of the Crab Nebula emission at high frequencies has been shed by observations made with the {\\em Chandra} X-ray satellite (Weisskopf et al. \\cite{weiss00}), which, thanks to the unprecedented spatial resolution, has revealed a number of intriguing features in the inner part of the nebula (see also Hester et al. \\cite{hester95,hester02}). The new details highlighted strengthen the view of the Crab Nebula as an axisymmetric object. In what is thought to be the equatorial plane of the pulsar rotation, {\\em Chandra} observations show the presence of a bright ring of emission, lying at a much closer distance to the pulsar than the already identified X-ray torus (e.g. Hester et al. \\cite{hester95}). The most puzzling discovery, however, is probably the presence of two opposite jet-like features oriented along an axis perpendicular to the plane of the torus and emerging from the very close vicinity of the pulsar. Similar features have been observed also in a number of other objects, namely around the Vela pulsar (Helfand et al. \\cite{helfand01}; Pavlov et al. \\cite{pavlov03}), PSR 1509-58 (Gaensler et al. \\cite{gaensler02}) and in the supernova remnants G0.9+01 (Gaensler et al. \\cite{gaensler01}) and G54.1+0.3 (Lu et al. \\cite{lu02}).  While the presence of a X-ray bright torus may be at least qualitatively explained within the framework of standard 1-D RMHD models (Kennel \\& Coroniti \\cite{kc}, KC84 hereafter; Emmering \\& Chevalier \\cite{ec87}), if we further assume that either the energy flux emerging from the pulsar or the termination shock dissipation efficiency is higher at low latitudes around the equator, the presence of jets that seem to emanate directly from the pulsar poses severe theoretical problems in its interpretation (Lyubarsky \\& Eichler \\cite{lyueic}), given the difficulties at explaining self-collimation of ultra-relativistic flows. A recent suggestion for an answer to this puzzle (Bogovalov \\& Khangoulian \\cite{bk1,bk2}; Lyubarsky \\cite{lyu02}) is that the jets are actually originating downstream of the pulsar wind termination shock, where the flow is only mildly or non-relativistic. If this is the case, the fact that they are observed starting from a much closer distance from the pulsar than where the shock in the equatorial plane is thought to be, has to be interpreted assuming that the given degree of anisotropy in the energy flow from the pulsar also causes the shock front to be highly non-spherical in shape, much closer to the pulsar along the rotation axis than in the equatorial plane. Moreover, even if the pulsar wind is weakly magnetized just upstream of the termination shock, the magnetic field inside the plerion can become as high as to reach equipartition. Therefore, collimation of the downstream flow may be easily achieved there by magnetic hoop stresses (Lyubarsky \\cite{lyu02}; Khangoulian \\& Bogovalov \\cite{kb}), resulting in plasma compression toward the axis and eventually in a polar jet-like outflow.  Thanks to the recent progress in numerical relativistic fluid dynamics and MHD (see Del Zanna \\& Bucciantini \\cite{luca1}; Del Zanna et al. \\cite{luca2}, and references therein), we are now able to start a more quantitative investigation of this problem by means of computer simulations. Our aim is to clarify whether a given latitude dependence of the pulsar wind energy flux may actually explain the jet-torus morphology observed at X-ray frequencies for the Crab Nebula and other plerions, and, if this is the case, what are the conclusions that one may infer on the structure and magnetization of the unshocked pulsar wind. Here we present the results of a first series of long-term 2-D axisymmetric RMHD simulations, from which some general conclusions on the physical mechanisms at work and useful scalings may already be derived (see Amato et al. \\cite{amato} for preliminary results). A similar numerical investigation has been recently carried out (Komissarov \\& Lyubarsky \\cite{kl}, KL03 hereafter), confirming the basic physical picture as viable for explaining the main observational features, as strongly suggested also by the close resemblance, at least at a qualitative level, between the map of simulated emission and {\\em Chandra} images of the Crab Nebula.  The paper structure is as follows. In Sect.~\\ref{sec:model} the pulsar wind model adopted at large distances from the light cylinder is sketched. In Sect.~\\ref{sec:setup} the numerical details of the simulations and the initial conditions are reported. Sect.~\\ref{sec:results} deals with the results of the simulations, split in three sub-sections for convenience. Finally the results are summarized in Sect.~\\ref{sec:final}, where conclusions are drawn for this preliminary work. ",
        "conclusions": "\\label{sec:final}  In this paper the structure and evolution of PWNe interacting with the surrounding SN ejecta has been analyzed by means of relativistic MHD axisymmetric simulations. Our main goal has been here the investigation of the mechanism originating the polar jets which are observed in X-rays in a growing number of PWNe.  The most recent and promising analytical studies (Bogovalov \\& Khangoulian \\cite{bk2}, Lyubarsky \\cite{lyu02}) start with the assumption that the pulsar wind is highly anisotropic, with a much larger energy flux at the equator than at the pole. Such a wind, interacting with the expanding SN ejecta produces a hot magnetized bubble with a torus-jet structure, as observed. The polar jet is originated by the magnetic hoop stresses in the PWN that, for high enough values of the wind magnetization parameter $\\sigma$, divert part of the equatorial flow toward the axis, collimating and accelerating it. The first numerical RMHD simulations (Amato et al. \\cite{amato}; KL03) confirm this scenario, and the simulated synchrotron emission map in the latter cited work strikingly resembles, at least qualitatively, the X-ray images of the Crab Nebula.  Here we have made an effort to improve on those preliminary simulations. The equatorial relativistic wind has a Lorentz factor as high as $\\gamma=100$, the magnetization parameter $\\sigma$ goes from $0.003$ up to $0.1$, which leads to magnetic fields in the PWN well beyond equipartition. The magnetic field shape assumed in the wind is that proper for an aligned or weakly oblique rotator, while in KL03 the assumed field is far from the $\\sin\\theta$ dependence, with a very broad region of low magnetization around the equator. The evolution is followed up to the beginning of the reverberation phase and comparison with previous models is made whenever possible.  The results show that the predicted self-similar evolution of the external PWN boundary is well reproduced at low magnetizations, and before reverberation effects begin. The expected elongation of the PWN due to magnetic pinching is also recovered, and we show how this effect increases in time and with $\\sigma$. The elongation of the nebula appears to be independent on the wind anisotropy, measured in our model by the parameter $\\alpha$.  In the inner part of the PWN the termination shock assumes an oblate shape, as expected for an anisotropic wind energy flux. An equatorial supersonic flow is produced in a complex shock structure. At intermediate latitudes, where the TS is highly oblique, the downstream flow is still supersonic. The plasma moves along the front gradually focusing toward the equator. This pattern holds in hydrodynamical simulations as well, since it is due only to the wind flux anisotropy. The equatorial flow that is driven by this focusing mechanism is supersonic, with typical velocities of $v\\approx 0.5 - 0.7c$, consistent with the values inferred for motion in the equatorial plane of the Crab Nebula (Hester et al. \\cite{hester02}).  When the wind magnetization is high enough ($\\sigma\\gsim 0.01$ in our simulations) the magnetic hoop stresses in the equatorial plane are so strong as to suppress completely this flow after a few termination shock radii. The plasma is then completely diverted toward the axis, where the magnetic compression finally drives a supersonic polar outflow with velocities that are once again in agreement with the values observed in the Crab and Vela PWNe ($v\\approx 0.3 - 0.7c$, see Hester et al. \\cite{hester02}; Pavlov et al. \\cite{pavlov03}). At later times the interaction of the polar jets with the contact discontinuity density jump may cause additional effects. The flow circulates all the way along the CD from high latitudes down to the equatorial plane. Here, these large scale vortexes drag some dense material from the ejecta toward the origin, with the effect of confining the equatorial outflow to the very inner parts of the nebula. Notice that the circulation is the opposite (from the equator to the pole) in the hydro and low $\\sigma$ cases. The value of $\\sigma$ that distinguishes the two regimes that we call highly and lowly magnetizatized depends on the expansion velocity of the CD: for nebulae expanding at a larger rate we expect the transition to occur for higher wind magnetizations.  Finally, for the high $\\sigma$ cases, the flow pattern strongly depends on the Poynting flux distribution in the wind. In particular, if in a narrow region around the equator the magnetic field vanishes, an additional vortex appears in the circulation pattern. The equatorial supersonic flow is never suppressed completely: it reaches the CD and then it circulates back toward the axis. The polar jet is still present and drives a vortex circulating at higher latitudes in the opposite direction than the former.  In conclusion, axisymmetric relativistic MHD simulations of the interaction of pulsar winds with expanding SNRs are able to reproduce at least qualitatively most of the structures seen in X-ray images: the overall toroidal structure of the plerion, the supersonic motions in the equatorial plane and, if the wind magnetization is high enough, also the presence of polar jets with supersonic velocities. A more detailed study, with a larger sampling of the parameter space would be necessary for a quantitative comparison with the observations. An interesting perspective that the present study suggests is the inference of the wind magnetic field structure from direct comparison between simulated synchrotron maps and X-ray observations. Preliminary results are encouraging, although we prefer to leave quantitative comparisons as future work.  Some of the observed features are anyway impossible to reproduce within the present axisymmetric RMHD framework. Emission structures like knots and sprites might well be related to non-ideal effects, like magnetic reconnection and dissipation, that are non-trivial to deal with. Remaining within the RMHD approximation, a full 3-D setting would allow to deal with some non-axisymmetric instabilities that may play an important role in PWNe: let us mention as an example the kink instability, which is likely to be at the origin of the bending of the jet (observed in both the Crab and Vela nebulae). We also leave the study of these important physical processes as future work."
    },
    "0404/astro-ph0404025_arXiv.txt": {
        "abstract": "We study the properties of the diffuse light in galaxy clusters forming in a large hydrodynamical cosmological simulation of the $\\Lambda$CDM cosmology.  The simulation includes a model for radiative cooling, star formation in dense cold gas, and feedback by SN-II explosions.  We select clusters having mass $M>10^{14} h^{-1} M_\\odot$ and study the spatial distribution of their star particles. While most stellar light is concentrated in gravitationally bound galaxies orbiting in the cluster potential, we find evidence for a substantial diffuse component, which may account for the extended halos of light observed around central cD galaxies.  We find that more massive simulated clusters have a larger fraction of stars in the diffuse light than the less massive ones.  The intracluster light is more centrally concentrated than the galaxy light, and the stars in the diffuse component are on average older than the stars in cluster galaxies, supporting the view that the diffuse light is not a random sampling of the stellar population in the cluster galaxies.  We thus expect that at least $\\sim 10\\%$ of the stars in a cluster may be distributed as intracluster light, largely hidden thus far due to its very low surface brightness. ",
        "introduction": "The presence of diffuse `intracluster light' in galaxy groups and clusters is now well established; observations by several groups provide estimates of the fraction of diffuse light and its distribution, using different techniques (see Arnaboldi 2003 for a review). The fraction of stars contained in this space-filling component seems to increase strongly with the density of the environment: from loose groups ($ < 2\\%$, Castro-Rodriguez et al. 2003; Durrell et al. 2003) to Virgo-like (~10\\%; Feldmeier et al. 2003; Arnaboldi et al. 2003) and rich clusters ($ \\sim 20\\%$ or higher; Gonzalez et al. 2000, Feldmeier et al. 2002; Gal-Yam et al. 2003). This correlation may represent an important clue for understanding the mechanisms that produce intracluster (IC) light and drive its evolution in the cluster environment.  Cosmological simulations of structure formation facilitate studies of the diffuse light and its expected properties. Dubinski (1998) constructed compound models of disk galaxies and placed them into a partially evolved simulation of cluster formation, allowing an evolutionary study of the dark matter and stellar components independently.  Using an empirical method to identify stellar tracer particles in high-resolution dark matter (DM) simulations, Napolitano et al. (2003) studied a Virgo-like cluster, finding evidence of a young dynamical age of the intracluster component.  The main limitations in these approaches is the restriction to collisionless dynamics.  In this Letter, we analyze for the first time the IC light formed in a cosmological hydrodynamical simulation including a self-consistent model for star formation.  In this method, no assumptions about the structural properties of the forming galaxies need to be made, and the gradual formation process of the stars, as well as their subsequent dynamical evolution in the non-linearly evolving gravitational potential can be seen as a direct consequence of the $\\Lambda$CDM initial conditions. It is therefore of immediate interest whether this theoretical formation scenario makes predictions for IC light consistent with observations. Using a large volume of $192^3\\, h^{-3}{\\rm Mpc}^3$, we can furthermore study a statistically significant sample of clusters at $z=0$, and analyze the correlations of properties of diffuse light with, e.g., cluster mass and X-ray temperatures. ",
        "conclusions": "We used a cosmological simulation of $(192 h^{-1}{\\rm Mpc})^3$ to study the statistical properties of the IC light in clusters of galaxies and the dependence of its physical properties on cluster mass and X-ray temperature. These predictions can be tested against known properties of cD halos and used to plan observational tests to understand the physical properties of IC light.  The presence of the IC component is evident when the whole distribution of stars in the simulated clusters is analysed in a way similar to Schombert's (1986) photometry of BCGs. Galaxies at the center of our simulated clusters have surface-brightness profiles which turn strongly upward in a $(\\mu,R^{1/\\alpha})$ plot. This light excess can be explained as IC stars orbiting in the cluster potential. Integrating its density distribution along the LOS, the slopes from our simulations are in agreement with those observed for the surface brightness profiles of the diffuse light in nearby clusters. In the Coma cluster, Bernstein et al. (1995) parametrize the surface brightness as $r^\\beta$ and find that the diffuse light is best fit by $\\beta=-1.3\\pm0.1$. In the Fornax cluster, the surface brightness profile of the cD envelope of NGC 1399 follows a power law of the form $\\propto r^\\beta$ with $\\beta=-1.5$ (Bicknell et al. 1989).  At large cluster radii, the surface brightness profile of the IC light appears more centrally concentrated than the surface brightness profile of cluster galaxies (see Fig.~\\ref{fig1} and Fig.~\\ref{fig2}). From the simulations we also obtained the redshifts $z_{form}$ at which the stars formed: those in the IC component have a $z_{form}$ distribution which differs from that in cluster galaxies, see Fig.~\\ref{fig3}. The ``unbound'' stars are formed earlier than the stars in galaxies. The prediction for an old stars' age in the diffuse component agrees with the HST observation of the IRGB stars in the Virgo IC field, e.g. $t> 2$Gyr (Durrell et al. 2002), and points toward the early tidal interactions as the preferred formation process for the IC light.  The different age and spatial distribution of the stars in the diffuse component indicate that it is a stellar population that is not a random sampling of the stellar populations in cluster galaxies.  The more massive clusters have the largest fraction $f$ of diffuse light (Fig.~\\ref{fig3}). It is $f>0.1$ for cluster masses $M > 10^{14}h^{-1}M_\\odot$.  Our simulations may thus explain the low inferred star-formation efficiency in clusters vs. less massive structures (David 1997). If only the {\\it bound} stellar mass is accounted for in the ratio of the total cluster stellar mass vs.\\ cluster gas mass in our LSCS, then this ratio decreases from groups to rich clusters. The observational trend would then be reproduced in the simulation.  Similarly, the disagreement found between the amount of stars produced in clusters in our LSCS and in observed clusters (see B04) is less severe, if an IC component is present in real clusters and has been systematically neglected when evaluating their internal stellar mass budget.  The main result of this work is that large cosmological hydrodynamical simulations are in qualitative agreement with the observed properties of diffuse light in galaxy clusters. A quantitative assessment will require additional numerical efforts and more observations.  A detailed study of the dynamical history of the unbound stellar population in our simulation will be presented in a forthcoming work."
    },
    "0404/astro-ph0404213_arXiv.txt": {
        "abstract": "We have been conducting deep searches at $\\sim$20\\,cm of $>$30 globular clusters (GCs) using the 305-m Arecibo telescope in Puerto Rico and the 100-m Green Bank telescope (GBT) in West Virginia. With roughly 80\\% of our search data analyzed, we have confirmed 13 new millisecond pulsars (MSPs), 12 of which are in binary systems, and at least three of which eclipse.  We currently have timing solutions for five of these systems and basic orbital and spin parameters for six others. ",
        "introduction": "The number of known MSPs in GCs has increased significantly in the last few years due to a number of targeted surveys (see the review by F.~Camilo and the contributions by A.~Possenti and B.~Jacoby in this volume) that have benefitted from increased computational resources, new large-bandwidth pulsar backends (primarily at 20\\,cm, and improved search algorithms.  Approximately half of the currently known $\\sim$80 pulsars in 23 GCs\\footnote{\\url{http://www.naic.edu/~pfreire/GCpsr.html}} were found in the past four years.  During the last three years we have been searching more than 30 GCs with the Arecibo and Green Bank telescopes. The high time and frequency resolution of these data, along with newly developed search algorithms \\citep*{rem02,rce03}, makes us significantly more sensitive than past surveys to sub-millisecond pulsations as well as to pulsars in ultra-compact binary systems.  To date, we have confirmed 13 new MSPs in six GCs. ",
        "conclusions": "Our on-going 20-cm survey of clusters using the GBT and Arecibo has been very successful, resulting in the discovery of 13 new MSPs. These detections have benefitted greatly from the upgraded Arecibo telescope, the WAPP pulsar backends, and perhaps most importantly, a prodigious amount of available computing power.  The vast majority of these systems would not have been discovered without both computationally expensive acceleration searches and repeated observations of the clusters.  We anticipate determining timing solutions for the majority of the systems --- or at least reliable orbital parameters --- within the next year.  We also believe that future multi-wavelength campaigns will uncover many more pulsars in these clusters.\\\\ {\\em Acknowledgements:} We wish to thank the Canadian Foundation for Innovation for the grant that purchased ``The Borg'' and made this project possible."
    },
    "0404/astro-ph0404135_arXiv.txt": {
        "abstract": "In this paper, I present results from theoretical and numerical (Monte Carlo) {\\it N-particle\\/} fully relativistic 4-D analysis of Penrose scattering processes (Compton and \\gggg) in the ergosphere of a supermassive or stellar mass Kerr (rotating) black hole.  Specifically, the escape conditions and the escaping orbits of the Penrose pair production (\\gggg)  electrons are analyzed, revealing that these particles escape along collimated, jet geodesic trajectories encircling the polar axis.  Such collimated vortical tightly wound coil-like  trajectories of relativistic particles are inherent properties of rotating black holes. The helical polar angles of escape for these $e^-e^+$ pairs range from $\\sim 40^o$ to $\\sim 0^o.5$ (for the highest energy particles).   These jet distributions appear to be  consistent with the astrophysical jets of active galactic nuclei (AGNs) and galactic black holes, and suggest a mechanism for precollimation within the inner radius of the dynamically stable accretion disk. ",
        "introduction": "\\label{sec:intro}  We now have observational evidence that black holes indeed exist in nature.  They are at the cores of quasars and other active galactic nuclei (AGNs) as well as sources in our Galaxy, commonly referred to as microquasars or galactic black holes. Many of these sources are associated with polar jets emanating from their cores. Black holes were theoretically predicted from Einstein's theory of general relativity.  Theoretical and numerical calculations (Williams 1991, 1995, 1999,  2001,  2002a, 2002b, 2003), described briefly below, show that Penrose (1969) gravitational-particle scattering processes are sufficient to describe energy-momentum extraction from a rotating Kerr (1963) black hole, from radii within the marginal stable orbit ($r\\la r_{\\rm ms}\\simeq 1.2M$, in gravitational units: $c=G=1$; ${a/M}=0.998$, where $a$ is the angular momentum per unit mass parameter and $M$ is the mass of the black hole), while electromagnetic interactions or magnetohydrodynamics (MHD) appear to govern the ``flow'' of polar jets of the extracted particles, escaping away from the central source, out to the observed distances, as suggested by observations (Junor, Biretta, \\& Livio 1999).  In the primary paper cited above (Williams 1995), theoretical model calculations involving Monte Carlo computer simulations of Compton scattering and electron-positron ($e^- e^+$) pair production processes in the ergosphere  of a supermassive ($\\sim 10^8 M_\\odot$) rotating black hole are presented. Particles from an accretion disk surrounding the rotating black hole fall into the ergosphere and scatter off particles that are in bound equatorially and nonequatorially confined orbits. The accretion flow is assumed to be of the form of the so-called {\\it two-temperature} bistable disk model, in which the disk can in principle exist in two phases: a thin disk (Novikov \\& Thorne 1973) and/or ion torus (Lightman \\& Eardley 1974; Shapiro, Lightman, \\& Eardley 1976; Eilek 1980; Eilek \\& Kafatos 1983), where the electrons and protons (or ions) can have separate temperatures of up to $\\sim 10^9$~K and $\\sim 10^{12}$~K, respectively. Note that disk models of this sort are also referred to as thin disk/ion corona models, and more recently, the ion corona has been called an advection dominated accretion flow (ADAF; Mahadevan, Narayan,  \\&  Krolik 1997). The Penrose mechanism, in general,  allows rotational energy of a Kerr black hole (KBH) to be extracted by scattered particles escaping from the ergosphere to large distances from the rotating black hole. The results of these model calculations show that the Penrose mechanism is capable of producing the astronomically observed high energy particles ($\\sim $~GeV) emitted by quasars and other AGNs.  This mechanism can extract hard X-ray to $\\gamma$-ray photons from Penrose Compton scatterings of initially low energy soft X-ray photons by target orbiting electrons in the ergosphere. The Penrose pair production (\\gggg) processes allow  relativistic $e^- e^+$ pairs to escape with energies up to $\\sim 4$~GeV, or greater depending on the form of the accretion disk; these pairs are produced when infalling low energy photons collide with bound, highly blueshifted photons at the $\\it photon~ orbit$. This process may very well be the origin of the relativistic electrons inferred  from observations to emerge from the cores of AGNs. Importantly, these model calculations show that the Lense-Thirring effect (Thirring \\& Lense 1918), i.e., the dragging of inertial frames into rotation, inside the ergosphere, caused by the angular momentum of the rotating black hole, results in a gravitomagnetic force being exerted on the scattered escaping particles. This force (which is the gravitational analog or resemblance of a magnetic  force) produces asymmetrical particle emissions in the polar direction, above and below the equatorial plane, consistent with the asymmetrical or one-sided jets observed in radio strong AGNs (Williams 2002a, 1999)\\footnote{Figures~1(c) and~1(d) of Williams (1999) are incorrect.  Figures~4(e) and~4(h), respectively, of Williams (2002a) are the correct figures, where in  Figure~4(h) the target electrons have both positive and negative equal absolute value polar coordinate angular momenta.}.   The dragging of inertial frames also causes the Penrose escaping particles to escape along vortical trajectories (as discussed in the following paragraphs).  These Penrose processes can apply to any size rotating black hole and, in general, to any type of relativistic elementary particle scattering energy-momentum exchange  process, inside the ergosphere, allowing particles to escape with rotational energy-momentum from the KBH.  Even in the context of MHD, according to the guiding center (Landau \\& Lifshitz 1975) approximation, the single-particle approach is essential close to the black hole (de Felice \\&  Carlotto 1997; Karas \\& Dov\\v{c}iak 1997; de Felice \\& Zonotti 2000), i.e., the behavior of individual particles moving along geodesics in the strong central gravitational force field is also that of the bulk of fluid elements. This suggests that even though a MHD simulation is not performed in this present paper, the results will be valid for the trajectories of particle flows and $e^-e^+$ pairs produced near the black-hole event horizon.  Nevertheless, the Penrose analysis presented here should be considered in the context of a  full-scale relativistic MHD simulation.  However, because of the proximity of the Penrose processes to the horizon, and the dominating effect of gravity,  the resulting trajectories are expected to be approximately the same in the MHD regime.  In the model calculations summarized above, in which energy-momentum is extracted from a rotating black hole, it is found that particles escape with relativistic velocities along vortical trajectories, above and below the equatorial plane, with small helical angles of escape, implying strong coil-like vortex plasma collimation (Williams  2001, 2002b). Thus, from these model calculations, it appears that the rotating black hole naturally produces particle trajectories collimated about the axis of symmetry, i.e., the polar axis. Such vortical orbits or trajectories have been discussed by some other authors (de Felice \\& Calvani 1972; de Felice \\& Curir 1992; de Felice \\& Carlotto 1997;  de Felice \\& Zanotti (2000); see also Bi\\v{c}\\'{a}k, Semer\\'{a}k, \\& Hadrava 1993; Karas \\& Dov\\v{c}iak 1997).  Their independent findings, deduced from the geodesic properties of the Kerr metric, and referred to as geometry induced collimation, can serve as confirmation of  Williams' (1991, 1995) theoretical and numerical calculated results, or vice versa.  In this paper, I examine the escape conditions and the resulting four-momentum vectors of the  Penrose pair production (\\gggg) processes, showing that the particles indeed escape to infinity in the form of vortical jets intrinsically collimated about the polar axis, because of the frame dragging of spacetime inside the ergosphere of the KBH. Importantly, we shall see that these $e^-e^+$ pairs escape without any appreciable interaction with the dynamically stable accretion disk particles. Now, although most of the Penrose Compton scattered photons escape along vortical trajectories as well, we are concerned, in this paper, only with the Penrose pair production (\\gggg) electrons, because these are the particles that compile the main constituents of the observed jets, being responsible for the synchrotron radiation and Doppler boosting (giving rise to superluminal motion). I refer the reader to the above references, particularly Williams (1995), for a thorough description of the ``Penrose-Williams'' (Williams 2002b) processes discussed in this present paper. ",
        "conclusions": "In this paper, I have shown that the Penrose pair production (\\gggg) electrons, of processes occurring at the photon orbit $r_{\\rm ph}$, escape to infinity along vortical trajectories about the polar axis, without any  appreciable interaction with the stable accretion disk particles (which are located at $r\\ga r_{\\rm ms}$). This is also expected to be true for the Penrose Compton scattered escaping photons that have inward directed radial momenta with turning points at $r\\sim r_{\\rm ph}$, escaping along vortical orbits concentric to the polar axis (Williams 2001, 2002b).  On the other hand, since the Penrose Compton scattering processes occur at radii $r_{\\rm mb}\\la r\\la r_{\\rm ms}$ the scattered photons with positive radial momenta probably have appreciable interaction with the inner region of the accretion disk, as suggested by the broad Fe K$\\alpha$ emission line at $\\sim 6$~keV observed in the bright Seyfert~1 galaxy MCG---6-30-15 (Wilms et al.~2001); see the qualitative description in Williams (2002b). Details concerning MCG---6-30-15 will be investigated in a future paper.  Overall, these vortical trajectories of escaping particles suggest that the KBH is responsible for the precollimation (de Felice \\& Curir 1992; de Felice \\& Carlotto 1997;  de Felice \\& Zanotti 2000; see also Bi\\v{c}\\'{a}k, Semer\\'{a}k, \\& Hadrava 1993; Karas \\& Dov\\v{c}iak 1997) of the observed jets of relativistic particles, emanating from the cores of objects  powered by black holes (compare Junor, Biretta, \\& Livio 1999). Note, the coil-like collimation of the vortical orbits, presented here, and those of the Penrose Compton scattering processes, are investigated in further details elsewhere.  Further investigations of this precollimation will include the self-induced dynamo magnetic field associated with the vortical orbiting escaping Penrose produced (\\gggg) pairs; such a field could be an important contribution to the observed synchrotron emission as well as  in maintaining collimation.  Recent polarization measures (e.g., see Homan 2004 and references therein) appear to be consistent with the dynamics and kinematics of these vortical polar trajectories of escaping electrons. The degree to which the above statement applies consistently, with the interpretation that the observed transverse rotation-measure gradients across the jets (Gabuzda \\& Murray 2003) are due to an intrinsic helical magnetic field structure associated with the accretion disk, must be looked at in detail.  The best-case scenario would be for the magnetic field of the accretion disk to assist in further collimating and accelerating the Penrose produced relativistic jet particles as they escape from the black hole out to observed distances.  In comparison with some MHD energy extraction models, of the Blandford-Znajek (1977) and Blandford-Payne (1982) type, we find the following. Firstly, because of the proximity of the Penrose pair production (\\gggg) processes (Williams 1995) to the event horizon, i.e., occurring at the photon orbit $r_{\\rm ph}$ inside the marginally bound orbit $r_{\\rm mb}$ (the last radius for any bound material particle before it falls directly into the black hole), the accretion disk magnetic field, proposed to be frozen to the plasma particles (e.g., Koide et al. 2002;  Koide  2003; Meier, Koide, \\& Uchida 2001), is expected to have a negligible effect on these pair production processes, even without considering that the existences of such magnetic fields or ``flux tubes'' near the event horizon are inconsistent with general relativistic findings by Bi\\v{c}\\'{a}k (2000) and Bi\\v{c}\\'{a}k \\& Ledvinka (2000).  These findings reveal that such fields, if aligned with the rotation axis, will be expelled from the horizon or redshifted away, for a rapidly rotating black hole (i.e., a near extreme KBH).  Now based on these findings, it is probably safe to say that the disk magnetic field may have little, if any, effect on the Penrose Compton scattering processes (Williams 1995), particularly those occurring at $\\sim r_{\\rm mb}$, the radii first to be populated by the target electrons (Williams 2002b) and where the most energy would be extracted for a specific $Q_e$ value [compare the bound electron orbits displayed in Figs.~1$c$ and~3$c$; see also Fig.~1(b) of Williams 1995]. The radius of marginally bound orbit $r_{\\rm mb}$ is closer to the event horizon ($\\Delta r\\equiv r_{\\rm mb}-r_+\\simeq 0.026M$) than it is to $r_{\\rm ms}$ ($\\Delta r\\equiv r_{\\rm mb}-r_{\\rm ms}\\simeq 0.111M$), where $r_+\\simeq 1.063M$ for $a=0.998M$ (Bardeen et al. 1972).  Secondly, the main advantage of the Penrose-Williams processes, over such Blandford-Znajek (1977) and Blandford-Payne (1982) type models, is that the Penrose-Williams gravitational energy-momentum extraction processes occur independently of a magnetic field, while producing highly relativistic particles that escape to infinity along intrinsically collimated vortical trajectories in the form of symmetrical or asymmetrical jets, because of the inertial frame dragging.  Another feature of these processes is that most of the energy extracted is gravitational binding energy with only a small fraction of the scattered particles having final negative energies (up to $\\sim 30$\\% of the Penrose Compton scattered photons and up to $\\sim 10$\\% of the Penrose pair produced electrons, occurring only in the lowest energy-momentum scattering processes), which means that the angular momentum of the black hole will not decrease significantly in these processes.  The lifetime of these processes is expected to be indefinite as long as there is matter to be accreted, since positive energy-momentum particles are also scattered into the KBH. The resulting escaping plasma jets are expected to generate a self-induced dynamo magnetic field, as stated above, in the form of a polar solenoid-like field, which could possibly magnetically confine and further assist in acceleration, collimation processes, as well as producing observed synchrotron radiation.  Importantly, as soon as the inner region unstable disk particles reach a temperature $T\\ga 17$~keV---the energy needed to populate the target particle orbits for Penrose Compton scattering, assuming, of course, the probable existence of turning points (compare eq.~[\\ref{eq:escape1}] and discussion in \\S~\\ref{sec:orbits}), this Penrose process can ``turn on'' (Williams 2002b), irrespective of the mass of the rotating black hole.  Moreover, the general observed high energy luminosity spectra of quasars and microquasars can be reproduced by these processes (see Williams 2002b, 2003).  The above features are general characteristics of observed black hole sources.  None of the 3-D MHD energy extraction jet models at the present (e.g., Koide et al. 2000; Meier et al. 2001; Koide et al. 2002; Koide 2003, 2004) exhibit such characteristic features in such details as does the Penrose-Williams mechanism.  Nevertheless, 3-D MHD models can possible achieve the necessary powers (Blandford \\& Znajek 1977; Meier et al. 2001), but, in some cases, only after assuming an unrealistically large strength magnetic field.  Lastly, it appears that the modern-day Blandford-Znajek (1977) and Blandford-Payne (1982) type models are essentially faced with the age-old problem of converting from electromagnetic energy to particle energy and thus the inability of generating the highly relativistic particles needed to be consistent with observations.  The problems associated with such models in the direct extraction of energy near the event horizon suggest that these models may be important in the weak gravitational field limit, serving perhaps the same purpose they do in the jets of protostars, i.e., appearing to have a dominant role on a large scale at distances outside the strong effects of general relativity.  The recent model of Koide (2004) and its application to $\\gamma$-ray bursts, assuming a large scale, superstrong radial magnetic field ($\\sim 10^{15}$~G) down to the event horizon of a rapidly rotating KBH ($a=0.99995$), produces a mildly relativistic outflow.  Not only, at least at the present, does this model not collimate, but it may not be consistent with the general relativistic determination that any ``radiatable'' multipole field gets radiated away completely as a star collapses to a KBH (de la Cruz, Chase, \\& Israel 1970; Price 1972), which includes the magnetic field, leaving only nonzero monopole parameters $M$, $J$, and $q$, where $J$ ($=Ma$) is the black-hole angular momentum, and $q$ is the electric charge, hence the ``no-hair theorem'' (Misner, Thorne, \\& Wheeler 1973; Carter 1973). Also, the problem remains in the model of Koide (2004) as to how such a superstrong field is created in the popular (or generally accepted) model of stellar evolution.  In any case, the Penrose-Williams processes described in this paper (see also Williams  2002a, 2002b, 2003) may be related to the beamed energy output observed in $\\gamma$-ray bursts.  It is suspected that these processes, which could ``quickly'' turn on and off under suitable conditions, might be important.  Particularly, these processes might prove to be invaluable in explaining the jets in so-called collapsars, i.e, black hole formation in massive stars with rotation (MacFadyen, Woosley, \\& Heger 2001), without the needed of a superstrong magnetic field.  Equally, considering an ``inactive'' and/or perhaps ``isolated'' rotating  black hole, if its tidal forces encounter and destroy an object of sufficient density, then as the debris is infalling, these Penrose processes could produce characteristics associated with some $\\gamma$-ray bursts: an energetic short-lived burst, collimation, synchrotron emission, as well as the afterglows in the X-ray (Costa et al. 1997), optical (van Paradijs et al. 1997), and radio (Frail et al. 1997) regimes.  That is, $\\gamma$-ray and X-ray jets can  be produced possibly from Penrose Compton scattering, and radio to optical synchrotron jet emission from subsequent Penrose pair production (\\gggg) electrons, interacting with the expected intrinsically induced magnetic field (mentioned above). Details of the application of these Penrose processes to $\\gamma$-ray bursts await future investigations.  Note, in general, evaluation in a full-scale relativistic MHD regime is needed to follow the evolution of the trajectories of the escaping Penrose particles and their interactions with the intrinsically induced magnetic field and that of the surrounding accretion disk.  This indeed will be a challenge. Yet, with existing MHD 3-D simulations, and since we have analytical expressions for the trajectories of the escaping particles (see Williams 2002a), such a task could be accomplished.  Finally, for completeness, even though the Penrose processes present here are quite efficient (Williams 1995) without a magnetic field, the presence of a disk magnetic field, however, inside the ergosphere, might increase the efficiency (see Wagh \\& Dadhich 1989 and references therein).  The effects of such a magnetic field, if sufficiently small, might be represented by a random motion superimposed on the orbital velocities of the charged particles, say the target electrons in Penrose Compton scattering (see e.g. Williams 1995)."
    },
    "0404/astro-ph0404303_arXiv.txt": {
        "abstract": "We study the evolution of a low-mass X-ray binary by coupling a binary stellar evolution code with a general relativistic code that describes the behaviour of the neutron star. We find that non-conservative mass transfer scenarios are required to prevent the formation of submillisecond pulsars and/or the collapse to a black hole. We discuss the sweeping effects of an active magneto-dipole rotator on the transferred matter as a promising mechanism to obtain highly non-conservative evolutions. ",
        "introduction": "The ``classical'' scenario for the formation of millisecond radio pulsar binaries with low-mass companions envisages four main stages (see {\\it e.g.} Bhattacharya \\& van den Heuvel 1991 for a review): (i) The magnetic moment $\\mu$ of the newly formed radio pulsar ($\\mu = B R^{3}$, where $B$ and $R$ are the surface magnetic field and the neutron star, NS, radius\\footnote{Actually, since the NS is a relativistic object, $R$ is the NS circumferential equatorial radius, i.e.\\ the proper circumference in the equatorial plane divided by $2 \\pi$.}, respectively) decays spontaneously on a e-folding timescale $t_{\\mu} \\sim 10^{7} - 10^{8}$ yr (e.g.\\ Lyne, Anderson \\& Salter 1982) from an initial value $\\sim 10^{31}$ G cm$^3$ to a final value $\\sim 10^{26} - 10^{27}$ G cm$^3$, when the decay probably stops (e.g.\\ Bhattacharya \\& Srinivasan 1986). (ii) Simultaneously, the pulsar spin period $P$ increases under magnetic dipole emission according to $L_{\\rm PSR} = (2/3c^{3}) \\mu^{2} (2 \\pi/P)^{4}$ (where $L_{\\rm PSR}$ is the bolometric magneto--dipole luminosity and $c$ is the speed of light). The pulsar swithces off when eventually crosses the ``death line'' in the $\\mu-P$ plane, {\\it i.e.} the line defined by the relation $\\mu_{26}/P^2 = 2 \\times 10^{3}$ (where $\\mu_{26} = \\mu / 10^{26}$ G cm$^3$) below which it is believed that the radio pulsar phenomenon does not take place (see {\\it e.g.} Ruderman \\& Sutherland 1975). (iii) The companion star overflows its Roche lobe and transfers mass with angular momentum to the NS {\\it via} a Keplerian accretion disc, thereby spinning it up to millisecond periods (close to the Keplerian period at inner rim of the accretion disk, see below) and back across the death line (recycling). During this phase the system is visible as a low-mass X-ray binary (LMXB). (iv) Mass transfer ceases: the NS is again visible as a radio pulsar whose spin rate decays under magnetic dipole emission very slowly as $L_{\\rm PSR} \\propto \\mu^{2}$ and $\\mu$ is reduced by three or four orders of magnitude. The end point is therefore a millisecond pulsar orbiting a low mass companion ($< 0.3$ \\msun) that is the remnant of the $\\sim 1$ \\msun\\ mass donor.   \\subsection*{Recycling in Transient Systems} Most LMXBs are NS Soft X-ray Transients (NSXT), i.e.\\ transient systems harboring a NS (see Campana \\etal 1998 for a review). Adopting the same conversion efficiency of the accreting matter energy into X-rays during the outbursts and the quiescent states (but see Barret \\etal 2000 for a different explanation), the inferred variations in the accretion rate are a factor $\\sim 10^{5}$.  NSXTs can provide a direct evidence of the recycling scenario, since during stage (iii) the mass transfer rate varies up to five orders of magnitude. During the LMXB phase the accretion disk is truncated because of one of the following reasons: (i) the interaction with the magnetic field of the NS, which truncates the disc at the magnetospheric radius $R_{\\rm M}$, at which the accretion flow is channeled along the magnetic field lines towards the magnetic poles onto the NS surface; (ii) the presence of the NS surface itself at $R$; and (iii) the lack of closed Keplerian orbits for radii smaller than the marginally stable orbit radius, $R_{\\rm MSO}$ (at few -- depending on the mass and spin of the compact object -- Schwarzschild radii from the NS centre). The position of $R_{\\rm M}$ is determined by the istantaneous balance of the pressure exerted by the accretion disc and the pressure exerted by the NS magnetic field: \\begin{equation} R_{\\rm M} = 1.0 \\times 10^{6} \\, \\phi\\ \\mu_{26}^{4/7} m^{-1/7} R_6^{-2/7} \\dot{m}^{-2/7} \\; \\; \\; {\\rm cm}\\ , \\label{eq:rma} \\end{equation} where $\\phi \\le 1$, $m$ is the NS gravitational mass in \\msun\\footnote{Actually, since the NS is a relativistic object, it is important to distinguish between the baryonic mass, roughly speaking a measure of the amount of matter, and the gravitational mass that is smaller by a factor of $\\sim (3/5) G M^2/R \\sim 0.1 M c^2$, corresponding to the binding energy of the NS.}, $R_{6}$ is the NS radius in units of $10^{6}$ cm, and $\\dot{m}$ is the baryonic mass accretion rate$^2$ in Eddington units (the Eddington accretion rate is $1.5 \\times 10^{-8} R_{6} \\; \\; \\; {\\rm M}_{\\odot} {\\rm yr}^{-1}$). Eq. 1 % shows that as $\\dot{m}$ decreases, $R_{\\rm M}$ expands.  Accretion onto a spinning magnetized NS is centrifugally inhibited once $R_{\\rm M}$ expands beyond the corotation radius $R_{\\rm CO}$, at which the Keplerian angular frequency of the orbiting matter is equal to the NS spin: $R_{\\rm CO} = 1.5 \\times 10^{6} \\, m^{1/3} P_{-3}^{2/3} \\; \\; \\; {\\rm cm}$ where $P_{-3}$ is the NS spin period in milliseconds. In this case the accreting matter could in principle be ejected from the system: this is called propeller phase (Illarionov \\& Sunyaev 1975). Finally, if $R_{\\rm M}$ further expands beyond the  light-cylinder radius (where an object corotating with the NS attains the speed of light, $R_{\\rm LC} = 4.8 \\times 10^{6} \\, P_{-3}  \\; \\; \\; {\\rm cm}$), the NS becomes generator of magnetodipole radiation and relativistic particles. Indeed,  a  common requirement  of  all  the  models of  the  emission mechanism  from   a  rotating  magnetic  dipole  is   that  the  space surrounding the NS  is free of matter up to $R_{\\rm LC}$.  Let us consider the behaviour of a NSXT at the end of an outburst. Adopting $\\dot m \\sim 1$ in outburst, eq. 1 % gives $R_{\\rm M} \\sim R \\sim 10^{6} \\; \\; {\\rm cm}$. In quiescence $\\dot{m} \\sim 10^{-5}$ and $R_{\\rm M}= (10^{-5})^{-2/7} \\times 10^6 = 2.7 \\times 10^{7} \\; \\; {\\rm cm}\\ > R_{\\rm LC}$, for spin periods up to few milliseconds. Therefore it is likely that, during the quiescent phase, a magneto-dipole emitter switches on (see {\\it e.g.} Stella \\etal 1994; Burderi \\etal 2001). In this case, as the NS $\\mu$ and $P$ place such a system above the ``death line'', it is plausible to expect that the NS turns-on as a millisecond radio pulsar until a new outburst episode pushes $R_{\\rm M}$ back, close to the NS surface, quenching radio emission and restoring accretion. ",
        "conclusions": ""
    },
    "0404/astro-ph0404559_arXiv.txt": {
        "abstract": "We report the detection, with \\hst, of an optical counterpart to the transient supersoft X-ray source 1E~1339.8+2837, in the globular cluster M3. The counterpart is found near the faint end of the subgiant branch in the $V$ vs $V-I$ color magnitude diagram, but is extremely bright in $U$. Variability is detected over a range of timescales suggesting the presence of an accretion disk and perhaps also ellipsoidal variations of the subgiant secondary. The optical colors of the binary are similar to those of cataclysmic variables recently discovered in 47~Tucanae and NGC~6397.  We suggest that magnetically channeled accretion may explain the relatively low X-ray luminosity of this source's supersoft state. ",
        "introduction": "\\label{sect.intro}  Supersoft X-ray sources (SSSs) are characterized by high bolometric luminosities ($10^{36}-10^{38}$ \\ergs) and extremely soft X-ray spectra (blackbody temperatures of 15--80 eV). These characteristics are best explained by white dwarfs (WDs) that are burning hydrogen-rich matter accreted onto their surface from a companion star at the rate of about $10^{-7}$ \\mdot\\ per year (Kahabka \\& van den Heuvel 1997). These accretion rates are about a hundred times higher than those typically found in cataclysmic variables (CVs), the more common form of mass accreting binary containing a WD. Such high accretion rates can be supplied by mass transfer on a thermal time scale from a slightly evolved companion star (mass 1.3--2.5 \\mdot) that is more massive than the WD (Kahabka \\& van den Heuvel 1997).  Of the nine known SSSs in the Milky Way (Kahabka \\& van den Heuvel 2003), only one is found in a globular cluster.  A low luminosity X-ray source (\\lx\\ $\\sim 7 \\times 10^{33}$ \\ergs) named 1E~1339.8+2837 was discovered just inside the 33$''$ core radius of M3 (NGC~5272) by {\\it Einstein} (Hertz \\& Grindlay 1983; we assume a distance of 10.1 kpc to M3, using Djorgovski 1993). This source had a much higher luminosity of $2 \\times 10^{35}$ \\ergs\\ when observed by ROSAT in January 1991 and January 1992 (Hertz, Grindlay, \\& Bailyn 1993) and a very soft spectrum ($kT \\approx 20$ eV, Hertz et al. 1993; $kT = 36 \\pm 20$ eV, Dotani et al. 1999). This supersoft outburst is 10-100 times less luminous than most other SSSs (see Kahabka \\& van den Heuvel 2003 and Kahabka \\& van den Heuvel 1997). The ROSAT luminosity and temperature imply a WD radius of only $9 \\times 10^{7}$ cm, suggesting that only part of the WD surface is undergoing burning. The source has been seen in quiescence (\\lx\\ $\\sim 10^{33}$ \\ergs) with a hard X-ray spectrum several times since the 1991/1992 outburst (6/1992, 7/1993, 7/1994, 7/1995 with ROSAT, and 1/1997 with ASCA; Dotani et al. 1999).  Because 1E~1339 is located in a globular cluster its distance and reddening are well known, helping in modeling of the system. Also, excellent constraints on the binary parameters are possible if the secondary in the system is found, combined with good optical photometry. No successful identification of an optical companion has yet been made, but very few searches have been reported in the literature. Hertz et al. (1993) report a star in the ROSAT error circle with UV magnitudes consistent with those of a horizontal branch star, and Geffert (1998) reports a blue star and a variable star within a 5$''$ error circle, but provides no further details.  Here, we report the first detection of an optical companion to 1E~1339 in M3 using the ROSAT X-ray position determined by Verbunt (2001). \\hst\\ imaging is used to combat crowding near the center of the cluster.  We briefly discuss the implications of this detection, addressing the origin of the supersoft X-ray emission, and the parameters of the binary system. More details about the implications of these results will be given in Kahabka et al. (2004; in preparation). ",
        "conclusions": "\\label{sect.disc}   To summarize, the properties of Star A (blue $U-V$ color, variability, and the close astrometric match with the ROSAT position of the SSS) mean that it is almost certainly the optical counterpart of 1E~1339.  The optical colors of the SSS are similar to those of the CVs in 47~Tuc and NGC~6397, suggesting that 1E~1339 may be a type of CV, consistent with the scenarios predicted by Hertz et al. (1993) and Dotani et al. (1999).  One popular scenario to allow the high mass transfer rates required for steady nuclear burning on the surface of the WD is for mass transfer to occur on a thermal timescale. This is believed to occur in binaries with a slightly evolved companion star that is more massive than the WD primary (Kahabka \\& van den Heuvel 1997).  With a secondary mass for the SSS of no more than 0.9 \\mdot\\ (less depending on how much mass has been lost), it should be easy to fulfill this requirement by normal stellar evolution in a primordial binary or an exchange collision between a primordial binary and a WD.  WDs with masses of $\\sim$0.55 \\mdot\\ are currently being produced in the cluster (applying the initial-to-final mass relation of Weidemann 2000), and WDs with similar masses will be common.  Therefore, our results are broadly consistent with expectations of high mass transfer rates to explain nuclear burning on the surface of the WD primary.  However, a slightly evolved companion star with a mass between 1.3 and 2.5 \\mdot\\ is ruled out by these observations, since the only globular cluster stars in this mass range (besides neutron stars) are blue stragglers.  A useful diagnostic for this system is the \\fxfopt\\ ratio. We used the 0.5-2.5 keV X-ray flux values of Dotani et al. (1999) from April and June 1995, and January 1997, and \\fopt\\ $=10^{-0.4V-5.43}$ to derive \\fxfopt=0.9--2.8. These large \\fxfopt\\ values are consistent with those found for magnetic (DQ~Her) systems by Verbunt et al. (1997). We speculate that magnetic behavior might explain why only part of the WD surface appeared to undergo hydrogen burning in 1991/1992. It has already been argued that large numbers of magnetic CVs may be found in globular clusters (Grindlay et al. 1995 \\& Edmonds et al. 1999).  We note that studies of field CVs have shown that some CVs can turn on as SSSs (Greiner et al. 1999 \\& Patterson et al. 1998), with luminosities ($10^{35}-10^{36}$ \\ergs) that are lower than those of the original set of Milky Way and Magellanic Cloud SSSs. One of these, T~Pyx, might be a magnetic CV (Patterson et al. 1998).  \\cha\\ observations, led by C.O.H., will provide new constraints on the X-ray spectrum and variability of 1E~1339, and, with a larger sample of optically identified X-ray sources, should give us excellent (0.1-0.2\") relative astrometry between \\hst\\ and \\cha.  One dataset was obtained in November 2003 and two will follow in 2004. The second of these \\cha\\ observations will be coordinated with optical and UV observations with \\hst/ACS in 2004.  Spectroscopic studies of this relatively bright binary should be possible with \\hst/STIS, but will be a challenge for ground-based telescopes because of the nearby giant star."
    },
    "0404/astro-ph0404590_arXiv.txt": {
        "abstract": "We revisit the computation of a ``snow line'' in a passive protoplanetary disk during the stage of planetesimal formation. We examine how shadowing and illumination in the vicinity of a planet affects where in the disk ice can form, making use of our method for calculating radiative transfer on disk perturbations with some improvements on the model.  We adopt a model for the unperturbed disk structure that is more consistent with observations and use opacities for reprocessed dust instead of interstellar medium dust.  We use the improved disk model to calculate the temperature variation for a range of planet masses and distances and find that planets at the gap-opening threshold can induce temperature variations of up to $\\pm30\\%$.  Temperature variations this significant may have ramifications for planetary accretion rates and migration rates. We discuss in particular the effect of temperature variations near the sublimation point of water, since the formation of ice can enhance the accretion rate of disk material onto a planet. Shadowing effects can cool the disk enough that ice will form closer to the star than previously expected, effectively moving the snow line inward. ",
        "introduction": "The concept of a ``snow line'' was introduced by \\citet{hayashi} and refers to the distance from the Sun at which the midplane temperature of the preplanetary solar nebula drops to the sublimation temperature of ice.  The presence of ice beyond the snow line, which Hayashi calculated to be at 2.7 AU, ought to enhance planet formation and explains the existence of the gas and ice giants in their present locations.  \\citet{snowline} revisited this issue, adopting an updated protoplanetary disk model in hydrostatic and radiative equilibrium, with little or no accretion heating (passive disk). They find that the snow line can be as close as $\\sim1$ AU, depending on disk parameters.  However, if the temperature of the disk varies with height \\citep[e.g.,][]{vertstruct}, different parts of the disk will reach the sublimation temperature of water, 170 K, at different radii, so rather than a snow ``line'' in the disk, there will be a snow ``transition.''  In this paper, we show that the presence of planets themselves can affect where in the disk the snow transition occurs.  We have previously shown that the presence of a protoplanet can affect the temperature structure in the protoplanetary disk \\citep[hereafter \\paperone{}]{paper1}.  These temperature variations affect where in the disk ice can form.  The goal of this study is to quantify this effect by undertaking a parameter study in which we calculate the temperature structure for a variety of planet masses and distances.  This work is a step toward bridging the gap between simulations of disk-protoplanet interactions and analytical models based on observations of protoplanetary disks.  High resolution two- and three dimensional simulations help us to understand the hydrodynamic and tidal interactions between protoplanets and disks \\citep[e.g.,][]{kley99,lsa99,bryden99, kdh01, bate}. However, a major shortcoming of all these codes is that they assume a very simple equation of state and include no radiative transfer effects. \\citet{boss01} does consider radiative transfer in the diffusion approximation, but these are simulations of relatively massive disks with high accretion rates and include only compressional and viscous heating, whereas our models concern passive disks in which stellar irradiation is the primary source of heating. Generally speaking, simulated protoplanetary disks are typically vertically isothermal and do not include heating from the central star. While they can probe gravitational and tidal effects of a planet in a disk, they cannot account for effects of shadowing and illumination on the temperature structure as a gap opens in a disk.  Conversely, the analytic disk models self-consistently calculate effects of radiative transfer, which is important because a major source of heating in circumstellar disks is irradiation from the central star \\citep[e.g.,][]{calvet,CG,vertstruct,DDN}.  Analytical models show that disk temperature structure can vary greatly with disk height due to heating at the surface from stellar irradiation and viscous heating at the midplane.  However, these models cannot account for perturbations in the disk such as those imposed by the presence of a planet because they only consider radiative transport in one or two dimensions. Monte Carlo simulations of radiative transfer on a disk with a gap created by a planet have been done, but these models are essentially two-dimensional as well \\citep{rice2003,WWBW}.  These models are interesting from an observational point of view because large planets are able to significantly change the disk structure, but they do not address what happens to planets below the gap-opening threshold, where planet growth and migration are poorly understood.  The temperature structure of a disk, particularly in the vertical direction, can have an important effect on the dynamics of disk-protoplanet interactions since waves in disks do not necessarily carry energy evenly with height if it is not vertically isothermal \\citep{98lubowogilvie}.  The local temperature structure can significantly affect how waves are dissipated in the disk, which will in turn affect tidal torques and migration rates.  In this paper, we make a number of modifications to the model presented in \\paperone{} for calculating radiative transfer on perturbed disks, primarily in the calculation of the disk properties. These changes are described in detail in \\S\\S\\ref{diskstruct} and \\ref{planet_in_disk}.  The algorithm for calculating radiative transfer remains essentially the same.  In \\S\\ref{diskstruct}, we calculate the structure of the unperturbed disk, in \\S\\ref{planet_in_disk} we calculate the effect of a protoplanet on the disk structure, and in \\S\\ref{radtrans} we review the method of calculating radiative transfer on a three-dimensional perturbation in the disk.  In \\S\\ref{results} we apply the revised method and analyze the results. We compare the results with those previously obtained in \\paperone{}, calculate the effect on the temperature of the disk photosphere over varying planet masses and distances, and discuss how these temperature variations change the locations where ice can form in the disk. Section \\ref{discussion} is a discussion of the results and implications. ",
        "conclusions": "\\label{discussion}  We have improved on the disk model from our previous paper \\citep{paper1} by updating the opacities and calculating the vertical temperature structure more self-consistently.  As a result, temperature perturbations in the disk's photosphere due to the influence of a protoplanet are greater in magnitude.  For planets at the gap opening threshold, temperature variations can be up to $\\pm30\\%$.  While these temperature perturbations are unlikely to be observed with even the most sensitive instruments, they may have significant effects on planet building.  The temperature variations are large enough to affect the composition and dynamics of the disk material near the planet, which can have consequences for planetary accretion and migration.  If the temperature changes enough to drop above or below the condensation or sublimation temperature for ice formation, this will change the size distribution of dust grains.  This will also change the composition of the gas as molecules freeze out onto the dust. The disk temperature can also change the time scale for the dust settling to the midplane, so that the dust-to-gas ratio may vary with height.  In particular, shadowing and illumination effects can change the locations in the disk where water ice can form.  We have shown that ice can form inward of 0.5 AU in the presence of a protoplanet, whereas without the protoplanet the minimum distance at which ice can form is at 0.57 AU. This may mean that accretion rates can be enhanced closer to the star than previously expected.  The temperature may also affect the movement of disk material near the planet by shifting the streamlines along which gas and dust move past the planet.  This is because of the pressure gradient that the temperature perturbation imposes on the disk.  Thus, accretion onto the planet may preferentially come from one side of the planet or the other.  This, along with the change in composition of disk material, can affect the growth rate and eventual composition of the planet.  In addition, the temperature perturbation may change the migration rate of the planet under Type I migration. \\citet{wardA} has demonstrated that the pressure gradient caused by typical disk temperature profiles contributes to increasing (decreasing) torques from the outer (inner) disk so that the total net torque will almost certainly cause inward migration of the planet. Therefore changing the temperature profile in the vicinity of the planet can change migrations rates.  In future work, we will address the questions raised about planet growth and migration in the light of the temperature variations that we have studied in this paper."
    },
    "0404/astro-ph0404245_arXiv.txt": {
        "abstract": "We present a new (semi-)analytic model for feedback in galaxy formation.  The interstellar medium (hereafter ISM) is modeled as a two-phase medium in pressure equilibrium, where the cold phase is fragmented into clouds with a given mass spectrum.  Cold gas infalls from an external halo.  Large clouds are continually formed by coagulation and destroyed by gravitational collapse.  Stars form in the collapsing clouds; the remnants of exploding type II supernovae (hereafter SNe) percolate into a single super-bubble (hereafter SB) that sweeps the ISM, heating the hot phase (if the SB is adiabatic) or cooling it (in the snowplow stage, when the interior gas of the SB has cooled).  Different feedback regimes are obtained whenever SBs are stopped either in the adiabatic or in the snowplow stage, either by pressure confinement or by blow-out.  The resulting feedback regimes occur in well-defined regions of the space defined by vertical scale-length and surface density of the structure.  In the adiabatic blow-out regime the efficiency of SNe in heating the ISM is rather low (\\circa5 per cent, with \\circa80 per cent of the energy budget injected into the external halo), and the outcoming ISM is self-regulated to a state that, in conditions typical of our galaxy, is similar to that found in the Milky Way.  Feedback is most efficient in the adiabatic confinement regime, where star-formation is hampered by the very high thermal pressure and the resulting inefficient coagulation.  In some significant regions of the parameter space confinement takes place in the snowplow stage; in this case the hot phase has a lower temperature and star formation is quicker.  In some critical cases, found at different densities in several regions of the parameter space, the hot phase is strongly depleted and the cold phase percolates the whole volume, giving rise to a burst of star formation.  While the hot phase is allowed to leak out of the star-forming region, and may give rise to a tenuous wind that escapes the potential well of a small galactic halo, strong galactic winds are predicted to happen only in critical cases or in the snowplow confinement regime whenever the SBs are able to percolate the volume.  This model provides a starting point for constructing a realistic grid of feedback solutions to be used in galaxy formation codes, either semi-analytic or numeric.  The predictive power of this model extends to many properties of the ISM, so that most parameters can be constrained by reproducing the main properties of the Milky Way. ",
        "introduction": "Galaxy formation is an open problem.  This is due to the complexity of the feedback processes that arise from the energetic activity of massive or dying stars, taking place through winds, ionizing photons and SN explosions (not to mention AGN).  These feedback processes involve a large range of scales and masses, from the sub-pc scale of star formation to the $\\ga$10 \\kpc\\ scale of galactic winds, and from 1 to $10^{12}$ \\msun\\ or more.  It is useful at this stage to identify ranges of scales in which different processes are dominant.  On $\\ga$1 \\kpc\\ spatial and $\\ga$$10^6$ \\msun\\ mass scales the dominant processes such as shock heating of gas, radiative cooling, disc formation, galaxy merging and tidal or ram-pressure stripping are closely related to the dark-matter halo hosting the galaxy and to its hierarchical assembly.  On scales ranging from $\\sim$1 \\pc\\ to $\\sim$1 \\kpc, or from $\\sim$1000 to $\\sim$$10^6$ \\msun, cool gas reaches suitable conditions for collapse and star formation, and the energy input from massive stars (through winds, UV photons and SNe) acts in shaping and sustaining the multi-phase structure of the ISM.  At smaller scales star formation takes place; it is most likely driven and self-limited by magneto-hydro-dynamical (MHD) turbulence.  This division is obviously meant to be only a rough approximation of reality.  Numerical simulations of whole galaxies are still limited to space and mass resolutions not much smaller than $\\sim$1 kpc and $10^6$ \\msun\\ respectively (see, e.g., Weinberg, Hernquist \\& Katz 2002; Steinmetz \\& Navarro 2002; Mathis et al. 2002; Lia, Portinari \\& Carraro 2002; Recchi et al. 2002; Pearce et al. 2001; Toft et al. 2002; Springel \\& Hernquist 2003; Tornatore et al. 2003; Governato et al. 2004).  They can address effectively the processes dominant in the large-scale range identified above, but the feedback processes acting on intermediate and small scales are ``sub-grid'' physics and are treated with simple heuristic models that require the introduction of free parameters.  Current models of semi-analytic galaxy formation treat feedback at a similar, phenomenological level (see, e.g., Cole et al. 2000; Somerville, Primack \\& Faber 2001; Diaferio et al. 2001; Poli et al. 2001; Hatton et al. 2003); they typically connect the efficiency of feedback to the circular velocity of the dark matter halo, with the aid of free parameters.  Models of galaxy formation that include a more detailed description of feedback have been presented, e.g., by Silk (1997, 2001), Ferrara \\& Tolstoy (2000), Efstathiou (2000), Tan (2000), Lin \\& Murray (2000), Hirashita, Burkert \\& Takeuchi (2001), Ferreras, Scannapieco \\& Silk (2002) or Shu, Mo \\& Mao (2003).  In this framework (semi-)analytic work can give a very useful contribution in selecting the physical processes that are most likely to contribute to feedback.  The focus of this paper is on modeling the intermediate range of scales defined above, where the physics of the ISM is in act.  The standard picture of the ISM is that of a multi-phase medium in rough pressure equilibrium; the reference model is that of McKee \\& Ostriker (1977), who considered a medium composed by cold, spherical clouds with temperature and density \\tc\\circa 100 \\K\\ and \\nc\\circa 10 \\cmt, kept confined by a hot phase with \\th\\circa $10^6$ \\K\\ and \\nh\\circa $10^{-3}$ \\cmt.  A warm phase of $T_{\\rm w}$\\circa 10$^4$ \\K\\ and $n_{\\rm w}$\\circa 10$^{-1}$ \\cmt\\ was produced at the interface.  This vision is partially confirmed by multi-wavelength observations (see, e.g., Heiles 2001), although reality appears more complex, suggesting the presence of at least 5 different phases.  This picture is challenged by the results of many simulation programs, aimed to the numerical modeling of the ISM (see, e.g., Mac Low et al. 1998; Ostriker, Gammie \\& Stone 1999; Avila-Reese \\& Vazquez-Semadeni 2001; Kritsuk \\& Norman 2002; see Mac Low 2003 and Vazquez-Semadeni 2002 for reviews).  In this context the ISM is dominated by compressible, supersonic, MHD turbulence.  These groups are still struggling to tame the full complexity of the problem, so that these simulations are not directly aimed to or easily usable by modeling of galaxy formation.  For our purposes it is worth mentioning some results.  The distributions of temperature and density of the simulated gas particles show a wide range of values without any strong multi-modality, but some broad peaks are anyway present. The distribution of pressure shows a much more limited range of values. Structures defined as overdensities are not static clouds but transient features of an overall fractal distribution (which is consistent with observations, see Chappell \\& Scalo 2001) that do not last more than a sound crossing time, unless they are gravitationally bound.  Thus, the ``classical'' picture of the ISM is not validated, but a model with multiple phases in rough pressure equilibrium can still be used, though with care, as a useful first-order approximation, able to catch some significant elements of the dynamics of the ISM.  The motivation for the present work is to investigate the kind of physical processes that arise in galaxy formation, in order to provide a grid of solutions for the behaviour of feedback in a wide range of realistic cases, to be used in simulations or semi-analytic models of galaxy formation.  We restrict to a two-phase medium in pressure equilibrium, composed by cold clouds embedded in a diffuse hot phase. The dynamics of the ISM is at present assumed to depend only on its ``local'' properties, leaving thus out ``large scale'' events like differential rotation, spiral arms, mergers, galactic winds and so on. These events will be introduced once the global characteristics of the galaxy are specified.  This paper is the first of a series aimed to modeling feedback in galaxy formation.  It presents a minimal feedback model with its main properties and results.  Preliminary results were presented by Monaco (2002; 2003).  An upcoming paper will focus on the destruction of collapsing, star-forming clouds (Monaco 2004, hereafter paper II).  The paper is organized as follows.  Section 2 describes the physical ingredients of the model, Section 3 introduces the system of equations used, Section 4 the main solutions.  Section 5 is devoted to a discussion of the results, and Section 6 gives the conclusions. Finally, three appendices give a list of frequently used symbols, a determination of the time scales of coagulation of cold clouds and a study of the fate of SBs in the \\nh--\\lmech\\ plane. ",
        "conclusions": "We have presented a model for feedback in galaxy formation, based on a two-phase ISM, that does not restrict to self-regulated, equilibrium solutions and neglects (for simplicity) the global structure of the galaxy, apart from its density, vertical scale-height and velocity dispersion of clouds.  From the dynamics of the SBs that arise from the collapsing ``molecular'' clouds, we have identified four possible regimes of feedback, depending on whether SBs blow out of the ``disc'' or remain pressure-confined, and whether they have time to enter the PDS stage.  For a reference set of parameter values we have studied the dynamics of the system in the vertical scale-height -- surface density plane, identifying the regions of the plane corresponding to different regimes.  Both blow-out and confinement mostly take place in the adiabatic regime.  In a Milky Way-like adiabatic blow-out case, the main characteristics of the ISM of the Galaxy are broadly reproduced.  In the adiabatic confinement regime the ISM is predicted to have higher pressure, temperature of hot phase and densities of both phases, and smaller collapsing clouds; in some cases the density of the cold phase could be high enough to trigger diffuse star formation.  PDS confinement is found for high-density, thick structures in significant regions of the parameter space.  In this case feedback is less effective, the hot phase cooler and star formation quicker.  In many cases the system becomes critical, in the sense that the hot phase is severely depleted and the cold phase percolates the whole volume.  This happens for very low-density thin systems (that would however be kept ionized by the cosmological UV background), in some regions of the parameter space also for low-density thick systems in adiabatic blow-out (that may correspond to some gas-rich dwarf galaxies) and for high-density thick systems in PDS confinement (that may correspond to high-redshift galaxies).  The most likely result of this critical behaviour is the sudden consumption by star formation of the cold gas accumulated by the galaxy; the dynamics switches from a ``candle''- to a ``bomb''-like solution.  The porosity of SBs is usually found to be much lower than unity. However, in some cases unit porosity is found while SBs are in the PDS stage.  This corresponds to the formation of a super-SB that sweeps the whole galaxy, removing most ISM from it.  These events, together with the critical solutions, are likely connected to the triggering of galactic winds.  With respect to previous models of feedback, the main parameters that are typically present, as the efficiency of feedback, the Schmidt law with its normalization, or the rate of blow-out and leak-out of gas from a star-forming galaxy, are predictions of the present model.  The parameter space is connected to the properties of the ISM, and can thus be constrained by observations of the Milky Way and nearby galaxies; most parameters can be fixed in principle by reproducing only the Milky Way.  Moreover, the mass flows used in this model can be fine tuned by comparing with future detailed simulations of the ISM in a forming galaxy that include all the main physical processes though to be at work.  This model does not restrict to self-regulated ISM, and presents a rich variety of solutions with a relatively limited set of parameters. Although the turbulent nature of the ISM is not explicitly taken into account, the model is thought to give a good approximation to the solution of the feedback problem.  The feedback regimes found here can be used, together with the refinements of the model that will be given in upcoming papers, to construct a realistic grid of feedback solutions to be used in galaxy formation codes, either semi-analytic or numeric."
    },
    "0404/astro-ph0404523_arXiv.txt": {
        "abstract": "}[2]{{\\footnotesize\\begin{center}ABSTRACT\\end{center}  \\vspace{1mm}\\par#1\\par \\noindent {~}{\\it #2}}}  \\newcommand{\\TabCap}[2]{\\begin{center}\\parbox[t]{#1}{\\begin{center} \\small {\\spaceskip 2pt plus 1pt minus 1pt T a b l e} \\refstepcounter{table}\\thetable \\\\[2mm] \\footnotesize #2 \\end{center}}\\end{center}}  \\newcommand{\\TableSep}[2]{\\begin{table}[p]\\vspace{#1} \\TabCap{#2}\\end{table}}  \\newcommand{\\FigCap}[1]{\\footnotesize\\par\\noindent Fig.\\  % \\refstepcounter{figure}\\thefigure. #1\\par}  \\newcommand{\\TableFont}{\\footnotesize} \\newcommand{\\TableFontIt}{\\ttit} \\newcommand{\\SetTableFont}[1]{\\renewcommand{\\TableFont}{#1}}  \\newcommand{\\MakeTable}[4]{\\begin{table}[htb]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newcommand{\\MakeTableSep}[4]{\\begin{table}[p]\\TabCap{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newcommand{\\TabCapp}[2]{\\begin{center}\\parbox[t]{#1}{\\centerline{ \\small {\\spaceskip 2pt plus 1pt minus 1pt T a b l e} \\refstepcounter{table}\\thetable} \\vskip0mm \\centerline{\\footnotesize #2}} \\vskip0mm \\end{center}}  \\newcommand{\\MakeTableSepp}[4]{\\begin{table}[p]\\TabCapp{#2}{#3} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newenvironment{references}% { \\footnotesize \\frenchspacing \\renewcommand{\\thesection}{} \\renewcommand{\\in}{{\\rm in }} \\renewcommand{\\AA}{Astron.\\ Astrophys.} \\newcommand{\\AAS}{Astron.~Astrophys.~Suppl.~Ser.} \\newcommand{\\ApJ}{Astrophys.\\ J.} \\newcommand{\\ApJS}{Astrophys.\\ J.~Suppl.~Ser.} \\newcommand{\\ApJL}{Astrophys.\\ J.~Letters} \\newcommand{\\AJ}{Astron.\\ J.} \\newcommand{\\IBVS}{IBVS} \\newcommand{\\PASP}{P.A.S.P.} \\newcommand{\\Acta}{Acta Astron.} \\newcommand{\\MNRAS}{MNRAS} \\renewcommand{\\and}{{\\rm and }} {We present new version of the OGLE-II catalog of eclipsing binary stars detected in the Small Magellanic Cloud, based on Difference Image Analysis catalog of variable stars in the Magellanic Clouds containing data collected from 1997 to 2000.  We found 1351 eclipsing binary stars in the central 2.4 square degree area of the SMC. 455 stars are newly discovered objects, not found in the previous release of the catalog. The eclipsing objects were selected with the automatic search algorithm based on the artificial neural network. The full catalog is accessible from the OGLE {\\sc Internet} archive.}{Keywords: binaries: eclipsing -- Magellanic Clouds -- Catalogs} ",
        "introduction": "Precise determination of distances to nearby galaxies is still one of the main goals of modern astrophysics. Eclipsing binary stars were used for this purpose for almost hundred years and in the last decade we witnessed their great ``comeback'' for two main reasons. First, very large telescopes with mirror diameter more than 6 meters can provide accurate spectroscopy of such faint stars as eclipsing binaries in nearby galaxies. Secondly, long time-base photometry and precise light curves of eclipsing binaries, mostly in Magellanic Clouds and Galactic bulge, are supplied as a by product of microlensing searches, \\eg MACHO (Alcock \\etal 1997) and OGLE (Udalski \\etal 1997a, Udalski \\etal 1998, Wyrzykowski \\etal 2003).  To date several attempts of distance determination with eclipsing binary method were presented (\\eg Fitzpatrick \\etal 2003, Fitzpatrick \\etal 2002, Ribas \\etal 2002), mostly to the LMC because this value is crucial for the distance scale. The extragalactic distance scale is tied to the LMC distance.  In last few years, the distance to the Small Magellanic Cloud was also determined with eclipsing binary method by several authors: Wyithe and Wilson (2001), Wyithe and Wilson (2002), Harries, Hilditch and Howarth (2003). Their papers were based on the photometry obtained by the OGLE collaboration (Udalski \\etal 1998), which contained data from the first 1.5 year of observations of the second phase of the OGLE survey (Udalski, Kubiak and Szyma{\\'n}ski 1997b). However, OGLE-II continued collecting data until the end of 2000. A much larger and almost complete subset of the OGLE-II images was reanalyzed with the image subtraction technique -- Difference Image Analysis ({\\.Z}ebru{\\'n}, Soszy{\\'n}ski and Wo{\\'z}niak 2001a). Variable stars detected in that study were presented in the catalog of variable stars in the Magellanic Clouds ({\\.Z}ebru{\\'n} \\etal 2001b).  The main aim of this paper is to provide a catalog of eclipsing binary stars in the SMC based on the DIA photometry. The catalog contains 1351 stars, from which only 896 were cross-identified in previous version of the catalog indicating that 455 stars are newly discovered eclipsing binary stars.  The search algorithm and classification method were identical with those used in the catalog of eclipsing binary stars in the LMC (Wyrzykowski \\etal 2003). We used artificial neural network for recognition of the variability type and divided discovered eclipsing binaries into three classical types: EA (Algol type), EB ($\\beta$~Lyr type) and EW (W~UMa type). The sample is reasonably complete at the level of the DIA catalog of variable stars in the SMC although the completeness of the latter drops rapidly for fainter objects. The sample allows statistical analysis of eclipsing binaries in the SMC and should provide a good material for testing theory of evolution of binary systems as well as for studying the evolution of the SMC, star formation or other projects. ",
        "conclusions": "We present 1351 eclipsing binary stars located in the central regions of the SMC found in the OGLE-II data collected during four observing seasons. The number of stars is, however, smaller than found in the previous release of the catalog (Udalski \\etal 1998) based on the first 1.5 year of OGLE-II observations. This is likely due to incompleteness of the DIA catalog of variable stars in the Magellanic Clouds ({\\.Z}ebru{\\'n} \\etal 2001b), from which data for the present search were taken. \\begin{figure}[htb] \\vglue-2mm \\centerline{\\includegraphics[width=10.3cm,height=10.3cm]{fig1.ps}} \\vspace*{-3mm} \\FigCap{Comparison of light curves of the eclipsing binary stars from the previous OGLE-II SMC catalog (Udalski \\etal 1998) (left column) and present version (right column).} \\end{figure}  Only 896 stars were cross-identified with the previous catalog. It means, that 455 stars presented here are newly discovered eclipsing binary stars. The main improvements compared to the previous catalog include much longer time-base of observations (4 years) and use of the DIA photometry instead of convensional PSF photometry.  In the dense stellar fields the former technique is superior to the latter. Overall, present light curves have lower photometric scatter, better phase coverage and they yield much more accurate periods.   A sample of light curves of eclipsing binary stars from both catalogs is presented in Fig.~1. Stars from the catalog of Udalski \\etal (1998) are in the left column while the stars from the current catalog are in the right column.  For some objects a small systematic shift up to about 0.05 mag  between light curves in the two catalogs can be seen. Most likely it is due to differences in the zero point of calibrations or, as mentioned above, inaccuracies in the conversion of the DIA flux differences to the magnitude scale.  Fig.~2 presents histogram of the DIA $I$-band brightness for all eclipsing binary stars found in the SMC (solid lines) and for those, which were cross-identified with the previous edition of the catalog (dotted line). The number of stars grows up to about ${I\\approx18}$~mag and then falls down to zero at ${I\\approx20}$~mag. Newly discovered eclipsing binary stars (difference of solid and dotted lines on the histogram) are distributed more or less proportionally to the number of stars at a given brightness. \\begin{figure}[p] \\centerline{\\includegraphics[width=10cm, height=8cm]{fig2.ps}} \\FigCap{Histogram of the DIA $I$-band brightness in 0.2~mag bins for all eclipsing binary stars (solid line) and only those, which were cross-identified with the previous edition of the catalog (dotted line).} \\vskip5mm \\vskip6cm fig3.jpg \\vskip6cm \\FigCap{OGLE-II fields in the SMC. Dots indicate positions of eclipsing stars. North is up and East to the left in the DSS image.} \\end{figure} Fig.~3 presents a picture of the SMC from the Digitized Sky Survey (DSS) with contours of the OGLE-II fields. Positions of the eclipsing binary stars are marked with dots. The stars are distributed proportionally to the density of the SMC stars, with the largest concentration in the fields SMC\\_SC4--SMC\\_SC6.  \\begin{figure}[htb] \\centerline{\\includegraphics[width=11cm,height=8.8cm]{fig4.ps}} \\FigCap{Histogram of periods of eclipsing binaries in 0.25~day bins. Dashed (red), dot-dashed (green) and solid (blue) lines correspond to classes EA, EB and EW respectively. Doted (black) line corresponds to all eclipsing objects found in the SMC. Additional 235 objects have periods longer than 10 days.} \\end{figure}  Fig.~4 shows the histogram of orbital periods of the SMC eclipsing stars in 0.25~day bins from 0 to 10~days. Dashed (red), dot-dashed (green) and solid (blue) lines correspond to classes EA, EB and EW, respectively, and dotted (black) line corresponds to all eclipsing objects. Additional 235 objects with periods longer than 10 days are distributed more or less uniformly and their number falls to zero at longer periods. The majority of stars are short period systems with the most frequent period of about 1~day. The longest period equals to 632.615 days (SMC\\_SC3 OGLE004402.68-725422.5), but both eclipses of this star are very similar and it is possible, that the real period is twice that long if the star has very faint secondary minima, invisible in our data. Another star (SMC\\_SC4 OGLE004617.60-731859.0) has the period equal to 580.5 days.  \\begin{figure}[htb] \\vglue-3mm \\centerline{\\includegraphics[width=13cm]{fig5.ps}} \\FigCap{Color-magnitude diagram of eclipsing binaries in the SMC. Red dots, green crosses and blue triangles mark EA, EB and EW type objects, respectively.} \\end{figure} Fig.~5. presents {\\it I} \\vs ${V-I}$ color-magnitude diagram for all eclipsing  binary stars from the catalog. EA, EB and EW classes are marked with different symbols. Fig.~5 indicates, that most of the eclipsing stars belong to the SMC, but there are also some foreground stars, mostly EW class objects.  \\begin{figure}[htb] \\vglue-3mm \\centerline{\\includegraphics[width=13cm]{fig6.ps}} \\FigCap{Color-magnitude diagram of eclipsing binaries in the SMC. Different symbols mark position of stars with short, medium, long and very long periods.} \\end{figure} Another CMD diagram is presented in Fig.~6. Eclipsing binary stars are divided into 4 groups depending on their periods: short (less than 3 days), medium (from 3 to 7 days), long (from 7 to 15 days) and very long (more than 15 days). Each group is marked on the CMD with different symbol and color. The majority of short and medium period eclipsing stars are located on the main sequence and belong to the young population of stars. Part of the long period stars are located also on the main sequence, but some of them lie on the lower giant branch. Very long period eclipsing stars are mostly concentrated on the red giant branch.   The new version of the catalog of eclipsing binary stars in the SMC based on the OGLE-II DIA catalog of variable stars in the Magellanic Clouds contains 1351 objects of three classical types: EA, EB and EW. 455 stars are newly discovered eclipsing binaries, not found in the previous edition of the catalog (Udalski \\etal 1998).  The exceptional good quality of the DIA photometry and  very long time-base of the OGLE-II observations  enabled construction of a uniform sample of eclipsing binaries  with high quality light curves and very accurate periods. The catalog provides observational material  for a variety of astrophysical studies in the SMC.  \\Acknow{We would like to thank Prof.~Bohdan Paczy{\\'n}ski for his encouragements and discussions about this work. This work was partly supported by the KBN grant 2P03D02523 to {\\L}.~Wyrzykowski, 2P03D02124 to A. Udalski and NASA grant NAG5-12212 and NSF grant AST-0204908 to B.~Paczy\\'nski. We acknowledge usage of the Digitized Sky Survey which was produced at the Space Telescope Science Institute based on photographic data obtained using the UK Schmidt Telescope, operated by the Royal Observatory Edinburgh.}"
    },
    "0404/astro-ph0404009_arXiv.txt": {
        "abstract": "We present three~methods for measuring the slope of the Galactic dust extinction law, $R_V$, and a method for measuring the fine-scale structure of dust clouds in the direction of differentially-reddened globular clusters.  We apply these techniques to $BVI$ photometry of stars in the low-latitude Galactic globular cluster NGC~4833 which displays spatially-variable extinction/reddening about a mean $\\langle{A_V}\\rangle\\approx1$.  An extensive suite of Monte Carlo simulations is used to characterize the efficacy of the methods.  The essence of the first two~methods is to determine, for an assumed value of $R_V$, the {\\it relative\\/} visual extinction $\\delta{A_V}$ of each cluster horizontal branch (HB) star with respect to an empirical HB locus; the locus is derived from the color-magnitude diagram (CMD) of a subset of stars in a small region near the cluster center for which differential extinction/reddening are relatively small.  A star-by-star comparison of $\\delta{A_V}$ from the ($B-V$,~$V$) CMD with that from the ($V-I$,~$V$) CMD is used to find the optimal $R_V$.  In the third method, $R_V$ is determined by minimizing the scatter in the HB in the ($B-V$,~$V$) CMD after correcting the photometry for extinction and reddening using the Schlegel, Finkbeiner, \\& Davis (1998) dust maps.  The weighted average of the results from the three~methods gives $R_V=3.0\\pm0.4$ for the dust along the line of sight to NGC~4833.  The fine-scale structure of the dust is quantified via the difference, $(\\Delta{A_V})_{ij}\\equiv (\\delta{A_V})_i-(\\delta{A_V})_j$, between pairs of cluster HB stars ($i$,\\,$j$) as a function of their angular separation $r_{ij}$.  The variance (mean square scatter) of $(\\Delta{A_V})_{ij}$ is found to have a power-law dependence on angular scale: ${\\rm var}(r)\\propto{r}^\\beta$, with $\\beta=+0.9\\pm0.1$.  This translates into an angular power spectrum $P(\\kappa)\\propto\\kappa^\\alpha$, with the index $\\alpha=-1.9\\pm0.1$ for $r\\sim1'$--$5'$, where $\\kappa\\equiv1/r$.  The dust angular power spectrum on small scales (from optical data) matches smoothly onto the larger-scale power spectrum derived from Schlegel et~al.'s far-infrared map of the dust thermal emission. ",
        "introduction": "Scattering, absorption, and reradiation of photons by dust grains affect the propagation of starlight and play a key role in regulating the energy balance in the interstellar medium.  Studies of these processes have provided a great deal of insight into the physical and chemical properties of dust grains (for reviews of dust properties, see Mathis 1990a; Witt 2003).  It is also important to disentangle the effects of intervening dust from many kinds of astronomical observations, ranging from photometry and spectroscopy of Galactic stars and external galaxies to mapping small-scale anisotropy in the Cosmic Microwave Background Radiation (CMBR).  \\subsection{Reddening Law and Dust Grain Properties}  The degree of absorption and scattering by dust depends, in general, on the wavelength of the incident radiation (Whitford 1958).  This dependence can be quantified in terms of the normalized dust extinction law $A_\\lambda/A_V$ (or reddening law), where extinction is the sum of absorption plus scattering. The ratio of total to selective extinction, $R_V\\equiv{A_V}/E(B-V)$, is a commonly-used measure of the slope of the extinction/reddening law.  Using data along many different Galactic sight-lines, Cardelli, Clayton, \\& Mathis (1989, hereafter CCM) found a tight correlation between the overall shape of the reddening law and the slope $R_V$.  CCM devised an $R_V$-based, one-parameter family of empirical fitting functions to characterize the observed range of extinction law shapes; this parameterization was later refined by Fitzpatrick (1999).  Galactic dust displays a wide range of behavior.  While diffuse interstellar dust is observed to have $R_V=3.1$ {\\it on average\\/} (Savage \\& Mathis 1979; CCM), this value is by no means universal.  Significant deviations from this canonical $R_V=3.1$ value (and corresponding differences in the overall shape of the extinction/reddening law) are known to exist for a variety of interstellar dust clouds, and these variations appear to be generally correlated with environment (Fitzpatrick 1999).  For example, $R_V$ has been found to range from 4 to 5 in dense molecular clouds (Mathis 1990b; Larson, Whittet, \\& Hough 1996; Whittet et~al.\\ 2001), whereas there are indications that more diffuse, high-latitude cirrus clouds may have $R_V$ values as small as $\\sim2$ (Fitzpatrick \\& Massa 1990; Larson et~al.\\ 1996; Szomoru \\& Guhathakurta 1999).  Studies based on stellar photometry from the Optical Gravitational Lens Experiment (OGLE) and MAssive Compact Halo Objects (MACHO) projects point to $R_V$ being substantially smaller than~3 in the direction of the Galactic bulge (Popowski 2000; Popowski, Cook, \\& Becker 2003; Udalski 2003; Sumi 2004; Popowski 2004).  Barbaro et~al.\\ (2001) find departures from the CCM parameterization in the ultraviolet portion of the Galactic extinction law.  It is worth noting that the above studies sample dust at a wide range of distances from the Sun: some probe the diffuse interstellar medium in its immediate vicinity whereas others, especially those at low Galactic latitudes, can probe dust at much larger distances well outside the Solar neighborhood (e.g.,~the RCrA cloud studied by Szomoru \\& Guhathakurta 1999).  Some studies have suggested that $R_V$ variations are tied to variations in the size distribution of dust grains from one line of sight to another.  In diffuse cirrus clouds, a relative abundance of small grains might explain the steep rise of the extinction curve into the ultraviolet (CCM; Larson et~al.\\ 1996).  By contrast, a high abundance of large grains, which grow readily by coagulation in dense molecular clouds, may explain the larger than average values of $R_V$ observed in these regions (Whittet et~al.\\ 2001).  Whittet et~al.\\ suggest that, in dense molecular clouds, $R_V$ remains close to the standard value of~3.1 except for lines of sight with unusually high extinction ($A_V\\ge3$).  This may indicate a more complicated relationship between $R_V$ and $A_V$.  Moreover, it has been argued that chemical composition can also play a role in determining the shape of the extinction law (Rhoads, Malhotra, \\& Kochanski 2004).  To improve our understanding of the dependence of $R_V$ on environment, measurements are necessary along many sight-lines through a broad range of cloud types.  Clusters with differential extinction/reddening have long been used to measure the $R_V$ of intervening dust (Mihalas \\& Routly 1968).  The traditional method of spectral typing individual cluster stars (e.g.,~on the Morgan-Keenan system), in conjunction with photometric measurements, allows for a direct and precise measurement of extinction and reddening.  However, this method of measuring $R_V$ requires high-quality, flux-calibrated spectra and a good understanding of the effects of metallicity on the energy output of stars.  Our study focuses instead on the use of broad-band photometry in light of the fact that an extensive suite of photometric data sets is currently available.  In this paper, we propose three~methods for measuring the $R_V$ of dust in the foreground of differentially-reddened globular clusters.  The first two~methods are closely related to each other and rely on accurate three-band optical photometry of cluster stars.  The third relies on two-band optical photometry and a map of the dust thermal emission (Schlegel, Finkbeiner, \\& Davis 1998, hereafter SFD) and is particularly well suited to wide-field data.  All three~methods are based on the premise that the width of the blue horizontal branch (HB) is minimized when the appropriate $R_V$ value is used to correct the photometry for extinction and reddening.  The methods are applied to NGC~4833, a low-latitude Galactic globular cluster with variable extinction across its face.  Realistic Monte Carlo simulations of the cluster data set are used to estimate the accuracy with which $R_V$ can be recovered.  \\subsection{``Cirrus''---The Rich Texture of Interstellar Dust}  Nearly two decades ago, Low et~al.\\ (1984) noticed diffuse background emission in the InfraRed Astronomical Satellite (IRAS) 60 and 100$\\,\\mu$m maps and termed it `infrared cirrus' due to its complex texture.  The cirrus has been associated with thermal emission from dust grains in the diffuse interstellar medium (Beichman 1987).  Combining the angular resolution of the IRAS observations with the photometric accuracy of the COsmic Background Explorer (COBE)/Diffuse InfraRed Background Experiment (DIRBE) data, SFD derived all-sky maps of the dust column density and mean temperature.  One of the primary uses of these maps has been to correct extragalactic sources for extinction and reddening.  The SFD maps show structure in the cirrus down to the smallest angular scales resolved by IRAS ($\\approx6'$).  Gautier et~al.\\ (1992) described the structure in the IRAS data in terms of its angular power spectrum, \\begin{equation} P(\\kappa)\\propto\\kappa^\\alpha~~~, \\end{equation} where $P$ is the Fourier power, $\\kappa$ is the reciprocal of the angular scale, and $\\alpha\\sim-3$ is the spectral index.  Similarly, Guhathakurta \\& Cutri (1994) examined both IRAS 100$\\,\\mu$m maps and optical CCD surface photometry of reprocessed starlight from dust grains: they found that power on small scales was dominated by stars and galaxies, and since it was impossible to completely separate out these components, they merely cleared out the obvious compact ``objects'' (stars and galaxies) and smoothed the residual image over arcminute scales; the resulting power spectrum had an index $\\alpha=-3$.  More recently, Kiss et~al.\\ (2003) observed 13~fields with Infrared Space Observatory (ISO)/ISOPHOT and found that the power spectrum index varies from field to field over the range $-5.3\\le\\alpha\\le-2.1$ for angular scales larger than $3'$.  This work extends the finding of $-3.6\\le\\alpha\\le-0.5$ by Herbstmeier et~al.\\ (1998) in their earlier ISOPHOT-based study.  These last two~studies indicate that $\\alpha\\sim-3.0$ is not universal.  In this study, the scatter in the HB of the differentially-reddened globular cluster NGC~4833 is used to measure the angular power spectrum of dust on scales smaller than the angular resolution of the thermal emission maps (e.g.,~IRAS, DIRBE, ISOPHOT).  For uncrowded data such as ours, the smallest angular scale down to which this method can be applied is set by the surface density of HB tracers which in turn determines the typical nearest-neighbor separation; in crowding-limited situations, the smallest angular scale is a few times the size of the stellar point spread function: few arcseconds for ground-based images and sub-arcsecond for {\\it Hubble Space Telescope\\/} images.  Because interstellar dust causes extinction and reddening of starlight and reprocesses the energy, dust patchiness has a profound impact on the accuracy of many astronomical measurements even at high Galactic latitudes.  Thermal emission from cold dust can confuse CMBR anisotropy measurements. Measurements of large-scale structure in galaxy surveys may be influenced by spatial non-uniformities in the foreground dust.  Studies of stellar populations are at risk because measurements of $T_{\\rm eff}$ can be impacted by reddening.  Extinction can affect the photometry of distance indicators and thereby produce systematic errors in the distance-scale ladder.  It is therefore important to quantify extinction/reddening variations on the smallest scales possible.  While the SFD study does an excellent job of characterizing moderate- to large-scale dust variations, their maps miss power on scales smaller than a few~arcminutes.\\\\\\\\  The photometry of NGC~4833 is presented in \\S\\,2.  The construction of simulated cluster data sets in discussed in \\S\\,3.  The $R_V$ measurement methods are outlined in \\S\\,4, along with their application to NGC~4833 and simulated data sets to estimate the accuracy of each method.  In \\S\\,5, we describe the angular power spectrum of cirrus in the direction of NGC~4833. Possible future extensions of this work are discussed in \\S\\,6.  The main conclusions of this paper are summarized in \\S\\,7. ",
        "conclusions": "\\begin{itemize} \\item[$\\bullet$]{We have demonstrated the use of three~methods, ``$A_V$~RMS'', ``$A_V$~Slope'', and ``Optical/IR'' methods, for determining the dust extinction law slope $R_V$ in the direction of differentially-reddened Galactic globular clusters, and have tested the methods on an extensive suite of simulated cluster data sets.}  \\item[$\\bullet$]{For cluster data sets with low photometric error, $\\sigma_{\\rm phot}\\lesssim0.01$~mag, the ``$A_V$~RMS'' and ``$A_V$~Slope'' methods can be used to determine $R_V$ to an accuracy of $\\sigma(R_V)=0.1$--0.3.  For cluster data sets with $\\sigma_{\\rm phot}\\gtrsim0.03$~mag, the two~methods yield $R_V$ to within $\\pm0.7$, with some systematic biases.}  \\item[$\\bullet$]{The ``Optical/IR'' method generally provides relatively imprecise estimates of $R_V$, $\\sigma(R_V)\\sim0.6$, over the full range of photometric errors explored.}  \\item[$\\bullet$]{Combining the results from all three~methods gives a mean extinction law slope of $R_V=3.0\\pm0.4$ for the line of sight towards the low-latitute Galactic globular cluster NGC~4833.}  \\item[$\\bullet$]{The scatter in the cluster HB is used to estimate the amount of small-scale structure in the dust complex in the foreground of NGC~4833. The Schlegel et~al.\\ (1998) IRAS+DIRBE-based map of the dust thermal emission averages over small-scale reddening variations that are $\\approx6$\\% of the mean reddening value.}  \\item[$\\bullet$]{Star-to-star variations in relative visual extinction across the face of NGC~4833 provide a measure of the foreground dust angular power spectrum for projected separations in the range $r\\sim1'$--$5'$.  This small-scale power spectrum derived from cluster optical data matches smoothly onto the larger-scale power spectrum derived from the SFD reddening map of the region.  The overall power spectrum is well fit by a power law: $P(\\kappa)\\propto\\kappa^\\alpha$, where $\\kappa$ is the reciprocal of the angular scale $r$, with spectral index $\\alpha\\approx-2.0\\pm0.1$.} \\end{itemize}"
    },
    "0404/astro-ph0404465_arXiv.txt": {
        "abstract": "We investigate the relation between circular velocity $\\vc$ and bulge velocity dispersion $\\sigma$ in spiral galaxies, based on literature data and new spectroscopic observations. We find a strong, nearly linear \\vcsigma\\ correlation with a negligible intrinsic scatter, and a striking agreement with the corresponding relation for elliptical galaxies. The least massive galaxies ($\\sigma<80~\\kms$) significantly deviate from this relation. We combine this \\vcsigma\\ correlation with the well-known \\MBHsigma\\ relation to obtain a tight correlation between circular velocity and supermassive black hole mass, and interpret this as observational evidence for a close link between supermassive black holes and the dark matter haloes in which they presumably formed. Apart from being an important ingredient for theoretical models of galaxy formation and evolution, the relation between $M_{\\text{BH}}$ and circular velocity has the potential to become an important practical tool in estimating supermassive black hole masses in spiral galaxies. ",
        "introduction": "The existence of supermassive black holes (SMBHs) in the nuclei of galaxies has been suspected for almost half a decade, as accretion onto SMBHs seemed the only logical explanation for the existence of quasars. HST observations have provided evidence that SMBHs with masses ranging from $10^6$ to $10^9$~\\Msun\\ are present in the centre of a few dozens of nearby (quiescent) galaxies. Be this sufficient evidence for the existence of SMBHs, we can now tackle more fundamental questions concerning their formation and evolution. An obvious way to proceed is the study of the relation between SMBHs and the galaxies that host them. It was found that black hole masses are correlated with parameters of the hot stellar components of their host galaxies. The tight \\MBHsigma\\ relation (Gebhardt et al.~2000; Ferrarese \\& Merritt~2000) is now the preferred paradigm to study SMBH demographics in galactic nuclei.  This apparently tight link between bulges and SMBHs reflects an important ingredient that should be reproduced (and thus hopefully explained) by theoretical models of galaxy formation. In fact, the tightness of the \\MBHsigma\\ correlation is somewhat surprising. In most of the state-of-the-art models, the total galaxy mass (or dark matter mass $\\MDM$), rather than the bulge mass, plays a fundamental role in shaping the SMBHs. A close correlation could therefore be expected between $\\MBH$ and $\\MDM$, rather than between $\\MBH$ and the bulge properties. Establishing whether the \\MBHsigma\\ or the \\MBHMDM\\ relation reflects the fundamental mode by which SMBHs form and evolve will ultimately rely on a comparison of the intrinsic scatter of the two correlations.  Unfortunately, a direct observational characterization of the \\MBHMDM\\ relation is currently impossible. Ferrarese~(2002b) first argued that a correlation between $\\MBH$ and $\\MDM$ should be reflected in an \\MBHvc\\ correlation, where $\\vc$ is the circular velocity in the flat part of the rotation curve of spiral galaxies. Indeed, in most of the state-of-the-art galaxy formation models, there is a one-to-one correspondence between the circular velocity and the mass of the dark matter halo. Unfortunately, there are (presently) only a handful of spiral galaxies with secure SMBH masses, and only two of them have a well-measured extended rotation curve. A way to avoid this problem is to adopt the tight \\MBHsigma\\ correlation in order to estimate black hole masses in a larger sample of galaxies. A tight correlation between SMBH mass and dark matter halo mass should thus appear in the form of a correlation between central velocity dispersion and circular velocity. Ferrarese~(2002b) presented a first attempt at establishing such a correlation. Baes et al.~(2003) significantly improved on these results by almost doubling the sample size. The present contribution is focused on the latter results. ",
        "conclusions": ""
    },
    "0404/astro-ph0404186_arXiv.txt": {
        "abstract": "We present the preliminary result of our project, consisting in studying the properties of a large sample of galaxy clusters. The M\\\"{u}nster Red Sky Survey, which is a large galaxy catalogue covering an area of about 5000 square degrees on the southern hemisphere serves as our observational basis. It is complete up to $r_F=18^m.3$. Creation of a cluster catalogue is the first step of our investigation. We propose to use the 2D Voronoi tessellation technique for identifying galaxy clusters in this 2D catalogue. Points with high values of the inverse Voronoi tessel area will be regarded as galaxy cluster centroids. We show that this approach works correctly. ",
        "introduction": "In this paper we present the first element of our project. It consists in studying the properties of a large sample of galaxy clusters. In order to perform such studies, we need a sample of clusters extracted in a uniform manner from a homogeneous set of data. Therefore, we chose the M\\\"{u}nster Red Sky Survey as our observational basis. The first step of investigation is to create the catalogue of galaxy clusters. There are three basic cluster detection algorithms: the matched filter algorithm (Postman et al. 1996), the adaptive matched filter algorithm (Kepner et al. 1999) and the Voronoi tessellation technique (Icke \\& van de Weygaert 1987, Zaninetti 1989, Ramella et al. 1999, 2001). Kim et al. (2002) made a comparison of these cluster-finding algorithms, using a Monte Carlo experiment with simulated clusters. We decided to apply the Voronoi tessellation technique for cluster detection. The Voronoi tessellation technique is completely non-parametric, and therefore sensitive to both symmetric and elongated clusters, allowing correct studies of non-spherically symmetric structures. For a distribution of seeds, the Voronoi tessellation creates polygonal cells containing one seed each and enclosing the whole area closest to its seed. This is the definition of a Voronoi cell in 2D. This natural partitioning of space by the Voronoi tessellation has been used to model the large-scale distribution of galaxies. ",
        "conclusions": "We show that the Voronoi tessellation allows us to find the overdense regions. The main point of a further step in the search of galaxy clusters is the determination of the contrast level. This is crucial for finding real clusters. This can be done using various methods. One of them consist in using the number overdensity in respect to the average background. Another possibility is to use the colour-magnitude relation combining the data from the M\\\"{u}nster Red Sky Survey with those from the APM catalogues. The determination of photometric redshifts can be considered as well. At present, we are able to conclude that the applied analysis allows us to find correctly the overdense regions in two-dimensional data."
    },
    "0404/astro-ph0404379_arXiv.txt": {
        "abstract": "We have measured the central structural properties for a sample of S0-Sbc galaxies down to scales of $\\sim$10~pc using {\\it Hubble Space Telescope} NICMOS images. We find that the photometric masses of the central star clusters, which occur in 58\\% of our sample, are related to their host bulge masses such that $\\MassPt = 10^{7.75\\pm0.15}\\,(\\MassBul/10^{10}\\MassSun)^{0.76\\pm 0.13}$.  Put together with recent data on bulges hosting supermassive black holes, we infer a {\\it non-linear} dependency of the `Central Massive Object' mass on the host bulge mass such that $\\MassCMO/\\MassSun = 10^{7.51\\pm 0.06} (\\MassBul/10^{10}\\,\\MassSun)^{0.84 \\pm 0.06}$.  We argue that the linear relation presented by Ferrarese et al.\\ is biased at the low-mass end by the inclusion of the disc light from lenticular galaxies in their sample. Matching our NICMOS data with wider-field, ground-based $K$-band images enabled us to sample from the nucleus to the disk-dominated region of each galaxy, and thus to perform a proper bulge-disk decomposition.  We found that the majority of our galaxies ($\\sim$90\\%) possess central light excesses which can be modeled with an inner exponential and/or an unresolved point source in the case of the nuclear star clusters.  All the extended nuclear components, with sizes of a few hundred pc, have disky isophotes, which suggest that they may be inner disks, rings, or bars; their colors are redder than those of the underlying bulge, arguing against a recent origin for their stellar populations. Surface brightness profiles (of the total galaxy light, and the bulge component on its own) rise inward to the resolution limit of the data, with a continuous distribution of logarithmic slopes from the low values typical of dwarf ellipticals ($0.1 \\leq \\gamma \\leq 0.3$) to the high values ($\\gamma\\sim 1$) typical of intermediate luminosity ellipticals; the nuclear slope bi-modality reported by others is not present in our sample. ",
        "introduction": "\\label{Sec:Introduction}  The {\\it Hubble Space Telescope (HST)} enables the study of the inner regions of nearby bulges and ellipticals down to spatial scales of $\\sim$10 pc, roughly one order of magnitude closer to the center than is feasible with typical ground-based data.  These inner regions contain a small fraction of the ellipsoid mass, but they harbor the highest density regions of the galaxies and contain useful clues to their formation.  The availability of NIR array detectors in the nineties fostered significant progress in the understanding of many aspects of bulges, including structural parameters, colors, dust content and stellar populations, as well as the scaling of disk and bulge parameters, using ground-based imaging (e.g.\\ Andredakis, Peletier, \\& Balcells 1995, hereafter APB95; de Jong 1996; Seigar \\& James 1998; Knapen et al.\\ 1995; Khosroshahi et al.\\ 2000;  Graham 2001a, hereafter G01; M\\\"ollenhoff \\& Heidt 2001; Graham 2001b, 2002; Eskridge et al.\\ 2002; MacArthur, Courteau, \\& Holtzman 2003; Castro-Rodr\\'\\i guez \\& Garz\\'on 2003).  NIR data helped to establish that exponential profiles provide better fits to the surface brightness profiles of bulges than \\r14\\ models (Kent et al.\\ 1991 for the MW bulge; Andredakis \\& Sanders 1994; de Jong 1996), and soon thereafter it was demonstrated that profiles of bulges of all Hubble types admit a particularly simple fit using the S\\'ersic (1963; see Graham \\& Driver 2005) function  \\begin{equation} I(R) = I(0)\\,\\exp\\{-b_n\\,(R/\\reff)^{1/n}\\} \\label{Eqn:Sersic} \\end{equation}  \\noindent(APB95; G01; M\\\"ollenhoff \\& Heidt 2001; MacArthur et al.\\ 2003; see Caon et al.~2003 for the case of elliptical galaxies). In eqn.~\\ref{Eqn:Sersic}, \\reff\\ is the half-light radius of the bulge, and $b_n\\approx 1.9992n-0.3271$.  The S\\'ersic index $n$, which measures the curvature of the surface brightness profile, scales with bulge-to-disk luminosity ratio (B/D) and with bulge luminosity.   The S\\'ersic index also provides a concentration parameter (Trujillo et al.\\ 2001) which strongly correlates with the velocity dispersion and central supermassive black hole mass (Graham et al.\\ 2001a, 2001b), hence it is linked to global physical parameters of the spheroid.  Numerical simulations also suggest that bulges have a range of profile shapes.  Aguerri, Balcells, \\& Peletier (2001) show that the accretion of dense satellites onto disk-bulge-halo galaxies yields a growth of both the S\\'ersic index and B/D, hinting that $n$ may be linked to the accretion history and to the growth of bulges.  $\\Lambda$ cold dark matter cosmological simulations of galaxy formation yield bulge-disk structures where the bulge profile shape ranges from exponential to \\r14 (Scannapieco \\& Tissera 2003; Sommer-Larsen, G\\\"oth, \\& Portinari 2003).  The results given above, derived from ground-based data, bear the question of whether the inner regions to which the \\HST\\ gives access also follow the S\\'ersic function.  Our picture of elliptical galaxy nuclei had to be revised in several ways after the \\HST\\  imaging campaigns.   Giant ellipticals often show a rather sudden inward flattening of their surface brightness profiles, confirming the result from ground-based data that some ellipticals have \"cores\" (Kormendy 1985), while intermediate-luminosity ellipticals ($-18\\leq  M_B \\leq -20.5$) do not show cores; their profiles approach power laws throughout the inner regions (\"power-law\" galaxies); see Faber et al.\\ (1997, hereafter F97) and Rest et al.\\ (2001, hereafter R01).  Inner profile slopes decrease toward fainter luminosities, and, for dwarf ellipticals, approach the slopes seen in the nuclei of giant, core galaxies, although dwarfs do not show profile discontinuities, i.e., do not show 'cores'  (Graham \\& Guzm\\'an 2003; Ferrarese et al.\\ 2006, hereafter F06).  Many cores of ellipticals and S0s are dusty, and a fraction of them harbor central unresolved sources at \\HST\\  resolution (Lauer et al.\\ 1995, hereafter L95; Phillips et al.\\ 1996; Carollo et al.~1997; Ravindranath et al.\\ 2001; Stiavelli et al.\\ 2001). Inasmuch as bulges share global similarities with ellipticals when studied from the ground, we enquire whether bulges show \"cores\", whether bulges show nuclear sources.  Bulges of disk galaxies have been targeted less often than ellipticals by the \\HST. Peletier et al.\\ (1999, hereafter Paper\\,I) analyzed a sample of 19 field S0-Sbc galaxies using WFPC2 F450W, F814W and NICMOS F160W images, with the goal of obtaining bulge stellar population diagnostics.  The combination of blue and NIR colors allowed them to put tight limits on the ages of bulge populations.  Ages of S0 to Sb bulges were found to be comparable to those of ellipticals in the Coma cluster, with a small age spread $<2$ Gyr (Sbc bulges showed colors corresponding to younger ages).   Nuclei were found to be dusty, with $A_V =0.6-1.0$ mag.  Carollo and collaborators surveyed mid- to late-type bulges using WFPC2 and NICMOS (e.g.\\ Carollo 1999; Carollo \\& Stiavelli 1998; Carollo et al.\\ 1997, 1998, 2001, 2002; Seigar et al.\\ 2002).  These authors focus on bulge structure.  They provide fits using the \\r14, exponential and Nuker models, and propose a structural classification of bulges into `\\r14 classical' and 'exponential'. Carollo et al.\\ (2002) find nuclear resolved components (NC) in the centers of 60\\% of the exponential bulges.  In their view, '\\r14' and 'exponential' bulges respectively show 'high' and 'low' nuclear profile slopes, a structural difference which would trace different formation histories.  Whether bulges come in two families with distinct structural properties has implications for formation mechanisms of bulges.  Several models have been proposed (see Wyse, Gilmore, \\& Franx 1997; Bouwens, Cayon, \\& Silk 1999; Kormendy \\& Kennicutt 2004): early collapse (Renzini 1999; Zoccali et al.~2003); mergers prior to disk formation (Kauffmann, Charlot \\& White 1996); satellite accretion (Pfenniger 1993; Aguerri et al.\\ 2001); and disk instabilities (Pfenniger \\& Norman 1990; Zhang 1999). Bulges with \\r14\\ structure fit in the early collapse or merger scenarios, while exponential bulges are destroyed by mergers (Aguerri et al.\\  2001) and may instead be expected from disk instabilities (Combes et al.\\ 1990). Edge-on, peanut-shaped bulges are known to have bar dynamics, and are therefore also expected to form from disk instabilities (Kuijken \\& Merrifield 1995; Bureau \\& Freeman 1999). The existence of two classes of bulges is commonly understood as evidence that massive bulges come from mergers while less massive bulges grow as a result of disk instabilities (see, e.g., Athanassoula 2005).  In this paper we analyze the structural properties of bulges of early- to intermediate-type galaxies at \\HST\\ resolution using the S0-Sbc sample presented in Paper I.  We address profile shapes, nuclear sources, nuclear slopes, and central massive black hole mass estimates.  Given the ability of the S\\'ersic model to describe the profiles of spheroids at ground-based resolution, we use the S\\'ersic model as our starting point and enquire whether the increased spatial resolution of the \\HST\\  contributes to support or to modify the ground-based picture.   We perform a bulge-disk decomposition of the surface brightness profiles using combined \\HST+ground-based profiles that sample the galaxy light distribution from the nucleus to the disk-dominated region.  Ignoring this step would bring up two problems: the un-modeled disk contribution to the inner profile would bias the bulge nuclear parameters; and, we would not be able to derive basic bulge parameters such as the total luminosity and the effective radius as the \\HST\\  images do not cover the entire bulge at the distances of our target galaxies.  We avoid using the \\r14\\ or exponential models, rather we focus on S\\'ersic fits to the bulge profiles to test if the profile shape dichotomy appears when it is not forced.  Our first results on bulge profile shapes using \\HST\\  data were presented in Balcells et al.\\ (2003, hereafter Paper\\,II).  In that paper we show that \\r14\\ bulge profiles are exceedingly rare.  In this and a companion paper (Balcells, Graham, \\& Peletier 2007, hereafter Paper~IV) we perform a comprehensive analysis of those profiles.  We will show that inner surface brightness profiles show excesses, over the best-fit bulge S\\'ersic model, which can be successfully modeled by adding central unresolved sources and/or inner exponential components to the fitting function (\\S\\,\\ref{Sec:BulgeDiskFits}).  Section~\\ref{Sec:ParameterUncertainties} provides details on the estimation of  parameter errors through fits to simulated profiles.  The subsequent sections analyze the properties of the nuclear excess light.  Sect.\\,\\ref{Sec:NuclearExtendedComponents} shows that the galaxies with extended nuclear components closely match those with nuclear disky isophotes, which suggests that the excess light in the surface brightness profiles comes from flattened components such as disks, rings or inner bars.  Sect.\\,\\ref{Sec:PointSources} derives luminosities and masses for the unresolved nuclear sources and addresses the Compact Massive Object (CMOs) paradigm, i.e., that nuclear star clusters are the low-mass extension to central supermassive black holes.  In \\S\\,\\ref{Sec:SMBH} we relate the point sources to black hole mass estimates from the bulge velocity dispersions.  Finally, in \\S\\,\\ref{Sec:Gammas} we present and discuss the nuclear surface brightness profile slopes and compare them to those of ellipticals, bulges and dwarf ellipticals. In Paper~IV, we discuss global bulge and disk scaling relations as inferred from the profile decompositions. A Hubble constant of $H_{0} = 75$ km\\,s$^{-1}$\\,Mpc$^{-1}$ is used throughout. ",
        "conclusions": "\\label{Sec:Conclusions}  At \\HST\\ resolution, nuclear photometric components, in addition to the S\\'ersic bulge and the exponential outer disk, are exceedingly common ($\\sim$90\\%) in early- to intermediate-type disk galaxies. Spatially-resolved nuclear components are found in 58\\% of our sample.  These components are geometrically flat systems, and could be disks, bars or rings.  The ones detected have comparable central surface brightness to the underlying bulges, but fainter such systems may exist.  The isophotal signatures indicate total sizes of a few hundred pc, similar to those of inner bars in double-barred galaxies.  Often, such components are reddened by dust; the evidence from optical and NIR colors, presented in Paper~I, as well as their high densities, suggest that they are old rather than late additions to the bulges.     A majority of the galaxies ($\\sim$58\\%) harbor sources unresolved by \\HST/NICMOS2.  They are most likely star clusters, with luminosities corresponding to 10--20 globular clusters, and similar to other unresolved sources found in the nuclei of ellipticals, dwarf ellipticals and bulges.   When combined with similar nuclear components in dE galaxies, their photometric masses scale with spheroid mass as $\\MassPt/\\MassSun = 10^{7.73\\pm 0.16} (\\MassBul/10^{10}\\,\\MassSun)^{0.76\\pm 0.13}$. Our central star clusters fall above the faint-ward extrapolation of \\MassSMBH--\\MassBul\\ relations derived by HR04 or F06.  In order to extend a CMO-style relation to faint spheroid luminosities, a moderate non-linearity is needed, and we propose the relation $\\MassCMO/\\MassSun = 10^{7.51\\pm 0.06} (\\MassBul/10^{10}\\,\\MassSun)^{0.84 \\pm 0.06}$. But we see additional difficulties with the CMO picture in that all of our PS show masses above the cluster-black hole transitional mass of $10^{7} \\MassSun$ proposed by Wehner \\& Harris (2006).  Bulge surface brightness profiles rise inward to the limit of the \\HST/NICMOS resolution, $\\sim$10 pc for the current sample.  While the inner bulge profiles deviate from pure power-laws, \"break radii\" in a Nuker-law sense are not present.  Structurally, the bulges of early- to intermediate-type galaxies may be globally grouped with the \"power-law\", intermediate- and low-luminosity elliptical galaxies. Negative logarithmic nuclear profile slopes of the S\\'ersic bulge components $\\gamma$ cover a continuous range of $0<\\gamma<1$, overlapping with dwarf-ellipticals at the faint end, and with intermediate luminosity ellipticals at the bright end.  We find no evidence to support a bimodal distribution of $\\gamma$ reported by others."
    },
    "0404/hep-ph0404234_arXiv.txt": {
        "abstract": "We derive general expressions at the one-loop level for the coefficients of the covariant structure of the neutrino self-energy in the presence of a constant magnetic field. The neutrino energy spectrum and index of refraction are obtained for neutral and charged media in the strong-field limit ($M_{W}\\gg \\sqrt{B}\\gg m_{e},T,\\mu ,\\left| \\mathbf{p}\\right| $) using the lowest Landau level approximation. The results found within the lowest Landau level approximation are numerically validated, summing in all Landau levels, for strong $B\\gg T^{2}$ and weakly-strong $B \\gtrsim T^{2}$ fields. The neutrino energy in leading order of the Fermi coupling constant is expressed as the sum of three terms: a kinetic-energy term, a term of interaction between the magnetic field and an induced neutrino magnetic moment, and a rest-energy term. The leading radiative correction to the kinetic-energy term depends linearly on the magnetic field strength and is independent of the chemical potential. The other two terms are only present in a charged medium. For strong and weakly-strong fields, it is found that the field-dependent correction to the neutrino energy in a neutral medium is much larger than the thermal one. Possible applications to cosmology and astrophysics are considered. \\pacs{13.15.+g, 14.60.Pq, 95.30.Cq, 98.80.Cq} ",
        "introduction": "There are many astrophysical systems on which the physics of neutrinos in a magnetic field plays an important role. Let us recall that proto-neutron stars typically possess very strong magnetic fields. Large magnetic fields $% B=10^{12}-10^{14}$  G have been associated with the surface of supernovas \\cite{Ginzburg} and neutron stars \\cite{Fushiki}, and fields perhaps as large as $% 10^{16}$  G with magnetars \\cite{Duncan}. Even larger fields could exist in the star's interior. It is presumed from the scalar virial theorem \\cite{Lai} that the interior field in neutron stars could be as high as $10^{18}$  G. A magnetic field as this ($\\sim 10^{18}$  G) in the interior of a compact star will be larger in two orders than the chemical potential characterizing its quark matter density.   Unveiling the interconnection between the star magnetic field and its particle current flows could shed new light to the question of the star evolution. For example, it is well known that neutrinos drive supernova dynamics from beginning to end. Neutrino emission and interactions play a crucial role in core collapse supernovae \\cite{SuperN}. Their eventual emission from the proto-neutron star contains nearly all the energy released in the star explosion. Neutrino luminosity, emissivity and the specific heat of the densest parts of the star are governed by charged and neutral current interactions involving matter at high densities and in the presence of strong magnetic fields. Thus, a total understanding of the star cooling mechanism in a strongly magnetized medium is crucial for astrophysics.   On the other hand, the explanation of large-scale magnetic fields observed in a number of galaxies, and in clusters of galaxies \\cite{Kronberg} seems to require the existence of seed fields of primordial origin \\cite{Elect-W}. According to several mechanisms \\cite{Grasso}, strong primordial fields $% \\sim 10^{24}$  G could be generated at the electroweak transition. Even larger fields have been associated with superconducting magnetic strings, which would generate fields $\\sim 10^{30}$  G in their vicinity if created after inflation \\cite{Witten}.   Were primordial magnetic fields present in the early universe, they would have had non-trivial consequences for particle-physics cosmology. For instance, as it is well known, oscillations between neutrino flavors may change the relative abundance of neutrino species and may thereby affect primordial nucleosynthesis (for a recent review on neutrinos in cosmology see \\cite {Dolgov}). Therefore, if a strong magnetic field ($M_{W}\\gg \\sqrt{B}\\gg m_{e},T,\\mu ,\\left| \\mathbf{p}\\right| $, with $ M_{W}$ and $m_{e}$ the W-boson and electron masses respectively) modifies the neutrino energy spectrum of different flavors in different ways, a primordial magnetic field can consequently influence the oscillation process in the primeval plasma \\cite{Australia}.   The propagation of neutrinos in magnetized media has been previously investigated by several authors \\cite{McKeon}-\\cite{Nieves}. Weak-field calculations were done for magnetized vacuum in Refs. \\cite{McKeon}, \\cite{Feldman}, and at $ T\\neq 0$ and $\\mu =0$ ($\\mu $ is the electric chemical potential) in Refs. \\cite{Weak-F}. In the $\\mu =0$ case, as long as $B< T^{2}$, both the field- and temperature-dependent leading contributions to the neutrino energy resulted of $1/M_{W}^{4}$-order \\cite{McKeon}-\\cite{Weak-F}. In the charged plasma \\cite{Nieves}, the field-dependent terms were much larger, $\\sim 1/M_{W}^{2}$, but they vanished in the ($\\mu \\rightarrow 0$)-limit. The weak-field results of papers \\cite{McKeon}-\\cite{Nieves} led to think that the magnetic-field effects could be significant in astrophysics, because of the field- and $\\mu$-dependent terms of order $1/M_{W}^{2}$, but were irrelevant ($\\sim 1/M_{W}^{4}$-order) in the early universe due to its charge-symmetric character ($\\mu=0$). Hence, it has been assumed that in cosmology the main contribution to neutrino energy was the purely thermal term of order $T^{4}/M_{W}^{4}$ \\cite{Raffelt}. However, as we will show below, a strong magnetic field ($M_{W}\\gg \\sqrt{B}\\gg m_{e},T,\\mu ,\\left| \\mathbf{p}\\right| $) gives rise to a new contribution to the neutrino energy that is linear in the field, independent of the chemical potential, and that is of the same order ($1/M_{W}^{2}$) as the largest terms found in the weak-field charged-medium case. This new result can turn magnetic-field effects relevant for cosmology.   In recent papers \\cite{Elizalde}, \\cite{Efrain}, we investigated the effects of a strong magnetic field on neutrinos in magnetized vacuum (i.e. with $T=0$ and $\\mu =0$). There, to facilitate the calculations in the strong-field limit, we extended the Ritus' Ep-eigenfunction method of diagonalization of the Green functions of spin-1/2 charged particle in electromagnetic field \\cite{Ritus}-\\cite{Ritus-Book}, to the case of spin-1 charged particles \\cite {Elizalde}. This formulation, which is particularly advantageous for strong-field calculations, provides an alternative method to the Schwinger approach to address QFT problems in electromagnetic backgrounds \\cite {Schwinger}. The use of the Ritus' method resulted very convenient to study the neutrino-self-energy in magnetized media, since it allowed to diagonalize in momentum space both the electron and the W-boson Green's functions in the presence of a magnetic field. Ritus' formalism has also been recently applied to investigate non-perturbative QFT in electromagnetic backgrounds \\cite{Ng}.   From the above considerations it is clear that strong magnetic fields can play a significant role in a variety of astrophysical systems, and possibly also in the early universe. For these applications, the analysis has to be carried out in the presence of a medium. Thus, in the present paper we extend the results obtained in papers \\cite{Elizalde} and \\cite{Efrain} to include finite temperature and density, performing a detailed study of the effects of a strong magnetic field on neutrino propagation in neutral and charged media, and discussing possible applications to astrophysics and cosmology.   We stress that in calculating the neutrino self-energy in a magnetized medium, we should consider, as usual, its vacuum and statistical parts. In this case the vacuum part depends on the magnetic field, and for strong fields it can make important contributions even at high temperatures $T^{2}\\gtrsim eB$. The reason is that the vacuum and statistical terms have different analytical behaviors, due to the lack of the statistic ultraviolet cutoff in the vacuum part. This fact gives rise to a field-dependent vacuum contribution ($1/M_{W}^{2}$-order) which is larger than the thermal one ($1/M_{W}^{4}$-order), into the self-energy. Therefore, as shown in this paper, a strong magnetic field can become more relevant than temperature for neutrino propagation in neutral media.   The plan of the paper is as follows. In Sec.~\\ref{selfen-gen-struc} we consider the radiative correction to the neutrino dispersion relation in the presence of a constant magnetic field. We introduce the general covariant structure of the neutrino self-energy in the presence of an external field, and find the dispersion relation as a function of the coefficients of each independent term in the covariant structure. The general form of the found dispersion relation goes beyond any given approximation and medium characteristic and serves as a guidance for particular applications. The general expressions for the coefficients of the self-energy structures in the presence of the magnetic field are found in the one-loop approximation in Sec.~\\ref{1-loop-selfen}. The leading behavior in the $1/M_{W}^{2}$ expansion of those coefficients are then calculated in Sec.~\\ref{strong-b-selfen} in the strong-field limit (i.e. in the lowest Landau level (LLL) approximation) for neutral and charged magnetized media. These results are then used in Sec.~\\ref{index-ref} to find the corresponding neutrino dispersion relations and the indexes of refraction in neutral and charged strongly magnetized media. In Sec.~\\ref{numerical}, the LLL-approximation is numerically corroborated by summing in all Landau levels and finding the values of the coefficients for parameter ranges corresponding to strong $B \\gg T^{2}$ and weakly strong $B \\gtrsim T^{2}$ fields. Possible applications to cosmology and astrophysics are discussed in Sec.~\\ref{applicat}. Finally, in Sec.~\\ref{conclusions} we summarize the main outcomes of the paper and make some final remarks. ",
        "conclusions": "In this paper we carried out a thorough study of the propagation of neutrinos in strongly magnetized neutral and charged media. We started from the most general structure of the neutrino self-energy in a magnetic field, expressing it as the sum of four independent covariant terms with coefficients that are functions of the physical variables of the theory and whose values depend on the approximation considered. General expressions of the four coefficients at one-loop approximation were given in Eqs. (\\ref{25}-\\ref{28}).  The coefficients were then calculated in the strong-field limit using the LLL aproximation for the electrons. The LLL was assumed to be valid in the parameter range $M_{W}\\gg \\sqrt{B}\\gg m_{e},\\left| \\mathbf{p}\\right| $, $eB\\gtrsim $ $T^{2}$. To justify it one should keep in mind that under these conditions most electrons would not have enough energy to overcome the gap between the Landau levels. Hence, they will be mainly confined to their lower levels and the leading contribution would come from the LLL. This assumption was also corroborated for the above parameters' range by numerical calculations summing in all Landau levels.  The dispersion relation of the neutrinos was written as a function of the four coefficients of the self-energy structures, allowing in this way to straightforwardly obtain the neutrino's energy in the strong-field limit for each physical case.  In concordance with results previously obtained in charged media at weak fields \\cite{Nieves},% \\cite{Weak-F}, in the strong-field case an energy term associated with the interaction between the magnetic field and the effective magnetic moment was also found at leading order in $G_{F}$. This interaction energy disappears in the neutral medium, since in a charged-symmetric plasma the contribution to the effective magnetic moment coming from electrons and positrons cancels out.  A main outcome of our investigation was to show that in strongly magnetized systems a term of different nature emerges in both charged and neutral media. The new term, which is linear in the magnetic field and of first order in $G_{F}$, enters as a correction to the neutrino kinetic energy in the presence of a strong-magnetic field. This correction is present even in a strongly magnetized vacuum, since it is related to the vacuum part ($T=0,$ $\\mu =0$) of the neutrino self-energy at $B\\neq 0$.  A characteristic of the field-dependent corrections to the neutrino energy is that they produce an anisotropic index of refraction, since neutrinos moving along different directions have different field-dependent dispersion relations. We should underline that while the magnetic moment interaction term produces a maximum field effect for neutrinos propagating along the field lines, the field correction to the kinetic energy does not contribute to those propagation modes, but on the contrary, the maximum kinetic-energy effect takes place for neutrinos propagating perpendicularly to the field direction. We stress that the anisotropy does not differentiate between neutrinos and antineutrinos.  The charged medium results reported in the current work can be of interest for the astrophysics of neutrinos in stars with large magnetic fields. On the other hand, our finding for the neutral medium can have applications in cosmology, if the existence of high primordial magnetic fields is finally confirmed. Contrary to some authors' belief \\cite{Weak-F},\\cite{Dolgov} that, regardless of the field intensity, the neutrino dispersion relation in the early universe is well approximated by the dispersion relation in the zero field medium, our results indicate that strong, and even weakly strong, magnetic fields can give rise to a contribution to the neutrino energy that is several orders larger than the pure thermal contribution.  The field-dependent correction to the neutrino energy in a neutral medium with strong magnetic field can have an impact in neutrino flavor-oscillations in the primeval plasma \\cite{Australia} and therefore affect primordial nucleosynthesis. Hence, this new effect could be important to establish possible limits to the strength of the primordial magnetic field.  \\medskip  \\textbf"
    },
    "0404/astro-ph0404192_arXiv.txt": {
        "abstract": "I review our understanding of the structure and kinematics of the Large Magellanic Cloud (LMC), with a particular focus on recent results. This is an important topic, given the status of the LMC as a benchmark for studies of microlensing, tidal interactions, stellar populations, and the extragalactic distance scale. I address the observed morphology and kinematics of the LMC; the angles under which we view the LMC disk; its in-plane and vertical structure; the LMC self-lensing contribution to the total microlensing optical depth; the LMC orbit around the Milky Way; and the origin and interpretation of the Magellanic Stream. Our understanding of these topics is evolving rapidly, in particular due to the many large photometric and kinematic datasets that have become available in the last few years. It has now been established that: the LMC is considerably elongated in its disk plane; the LMC disk is thicker than previously believed; the LMC disk may have warps and twists; the LMC may have a pressure-supported halo; the inner regions of the LMC show unexpected complexities in their vertical structure; and precession and nutation of the LMC disk plane contribute measurably to the observed line-of-sight velocity field. However, many open questions remain and more work is needed before we can expect to converge on a fully coherent structural, dynamical and evolutionary picture that explains all observed features of the LMC. ",
        "introduction": "\\label{s:intro}  The Large Magellanic Cloud (LMC) is one of our closest neighbor galaxies at a distance of $\\sim 50 \\kpc$. The Sagitarrius dwarf is closer at $\\sim 24 \\kpc$, but its contrast with respect to the Milky Way foreground stars is so low that it was discovered only about a decade ago. The LMC is therefore the closest, big, easily observable galaxy from our vantage point in the Milky Way. As such, it has become a benchmark for studies on various topics. It is of fundamental importance for studies of stellar populations and the interstellar medium (ISM), it is being used to study the presence of dark objects in the Galactic Halo through microlensing (e.g., Alcock \\etal 2000a), and it plays a key role in determinations of the cosmological distance scale (e.g., Freedman \\etal 2001). For all these applications it is important to have an understanding of the structure and kinematics of the LMC. This is the topic of the present review. For information on other aspects of the LMC, the reader is referred to the book by Westerlund (1997). The book by van den Bergh (2000) discusses more generally how the properties of the LMC compare to those of other galaxies in the Local Group.  The Small Magellanic Cloud (SMC) at a distance of $\\sim 62 \\kpc$ is a little further from us than the LMC, and is about 5 times less massive. Its structure is more irregular than that of the LMC, and it is less well studied and understood. Recent studies of SMC structure and kinematics include the work by Hatzidimitriou \\etal (1997), Udalski \\etal (1998), Stanimirovic \\etal (1999, 2004), Kunkel, Demers \\& Irwin (2000), Cioni, Habing \\& Israel (2000b), Zaritsky \\etal (2000, 2002), Crowl \\etal (2001) and Maragoudaki \\etal (2001). However, our overall understanding of SMC structure and kinematics has not evolved much since the reviews by Westerlund and van den Bergh. The present review is therefore restricted to the LMC. ",
        "conclusions": "\\label{s:conc}  The structure and kinematics of the LMC continue to be active areas of research. As outlined in this review, much progress has been made recently. Improved datasets have played a key role in this, most notably the advent of large stellar datasets of magnitudes in many bands, lightcurves, and line-of-sight kinematics, and also the availability of sensitive HI observations over large areas. As a result we now have a fairly good understanding of the LMC morphology and kinematics. The proper motion of the LMC is reasonably well measured and the global properties of the LMC orbit around the Milky Way are understood. The angles that determine how we view the LMC are now known much more accurately than before and this has led to the realization that the LMC is quite elliptical in its disk plane. We are starting to delineate the vertical structure of the LMC and are finding complexities that were not previously expected.  Despite the excellent progress, many questions on LMC structure still remain open. Why is the bar offset from the center of the outer isophotes of the LMC? Why is the dynamical center of the HI offset from the center of the bar, from the center of the outer isophotes, and from the dynamical center of the carbon stars? Why do studies of the inner and outer regions of the LMC yield differences in line-of-nodes position angle of up to $30^{\\circ}$? Does the LMC have a pressure supported halo? Are there populations of stars at large distances from the LMC plane? What is the origin and dynamical nature of the non-planar structures detected in the inner regions of the LMC? Do different tracers outline the same non-planar structures?  It might be necessary to answer all of these open questions before we can convince ourselves that the optical depth for LMC self-lensing has been correctly estimated. This seems to be the most critical step in establishing whether or not the Milky Way halo contains hitherto unknown compact lensing objects (MACHOs). The open questions about LMC structure are important also in their own right. The tidal interaction between the Magellanic Clouds and the Milky Way provides one of our best laboratories for studying the processes of tidal disruption and hierarchical merging by which all galaxies are believed to grow. A better understanding of LMC structure may also provide new insight into the origin of the Magellanic Stream, which continues to be debated. And with improved proper motion measurements of the Magellanic Clouds, the Stream may become a unique tool to constrain the shape and radial density distribution of the Milky Way halo at radii inaccessible using other tracers."
    },
    "0404/astro-ph0404471_arXiv.txt": {
        "abstract": "We recently identified a substantial population of galaxies at $z>2$ with comparatively red rest-frame optical colors. These distant red galaxies (DRGs) are efficiently selected by the simple observed color criterion $J_s-K_s>2.3$. In this paper we present near-infrared spectroscopy with Keck/NIRSPEC of six DRGs with previously measured redshifts $2.4<z<3.2$, two of which were known to host an active nucleus. We detect continuum emission and emission lines of all observed galaxies.  Equivalent widths of H$\\alpha$ in the non-active galaxies are 20\\,\\AA --30\\,\\AA, smaller than measured for Lyman break galaxies (LBGs) and nearby luminous infrared galaxies, and comparable to  normal nearby galaxies. The modest equivalent widths imply that the galaxies either have a decreasing star formation rate, or that they are very dusty. Fitting both the photometry and the H$\\alpha$ lines, we find continuum extinction $A_V =1-2$\\,mag, ages $1 - 2.5$\\,Gyr, star formation rates $200-400\\,M_{\\odot}$\\,yr$^{-1}$, and stellar masses $1-5 \\times 10^{11}$\\,$M_{\\odot}$ for models with constant star formation rates. Models with a declining star formation lead to significantly lower extinction, star formation rates, and ages, but similar stellar masses. From [N\\,{\\sc ii}]\\,/\\,H$\\alpha$ ratios we infer that the metallicities are high, $1-1.5 \\times$  Solar. For four galaxies we can determine line widths from the optical emission lines. The widths are high, ranging from $130$\\,\\kms\\ to 240\\,\\kms, and by combining data for LBGs and DRGs we find significant correlations between linewidth and restframe $U-V$ color, and between linewidth and stellar mass. The latter correlation has a similar slope and offset as the ``baryonic Tully-Fisher relation'' for nearby galaxies.  From the linewidths and effective radii we infer dynamical masses and mass-to-light ($M/L$) ratios. The median dynamical mass is $\\sim 2 \\times 10^{11}\\,M_{\\odot}$, supporting the high stellar masses inferred from the photometry. We find that the median $M/L_V \\approx 0.8\\,(M/L)_{\\odot}$, a factor of $\\sim 5$ higher than measured for LBGs.  We infer from our small sample that DRGs are dustier, more metal rich, more massive, and have higher ages than $z\\approx 3$ LBGs of the same rest-frame $V$-band luminosity. Although their volume density is still uncertain, their high $M/L$ ratios imply that they contribute significantly  to the stellar mass density at $z\\approx 2.5$.  As their stellar masses are comparable to those of early-type galaxies they may have already assembled most of their final mass. ",
        "introduction": "It is still not known how and when massive galaxies were formed. Early studies envisioned an epoch of rapid collapse at high redshift, followed by sustained star formation in a disk and a smooth and regular dimming of the stellar light in the spheroidal component (e.g., {Eggen}, {Lynden-Bell}, \\&  {Sandage} 1962). More recently, hierarchical galaxy formation models in Cold Dark Matter (CDM) cosmologies have postulated that massive galaxies have complex assembly histories, and were built up gradually through mergers and periods of star formation ({White} \\& {Frenk} 1991). In these models properties such as mass, star formation rate, and morphology are transient, depending largely on the merger history and the time elapsed since the most recent merger (e.g., {Kauffmann}, {White}, \\&  {Guiderdoni} 1993; {Kauffmann} {et~al.} 1999; {Baugh} {et~al.} 1998; {Meza} {et~al.} 2003).  One of the most direct tests of hierarchical galaxy formation models is the predicted decline of the abundance of massive galaxies with redshift. The best studied galaxies at high redshift are $z\\approx 3$ Lyman break galaxies (LBGs) ({Steidel} \\& {Hamilton} 1992; {Steidel}, {Pettini}, \\&  {Hamilton} 1995; {Steidel} {et~al.} 1996).  Measurements have been made of their clustering properties ({Giavalisco} {et~al.} 1998), star formation histories ({Papovich}, {Dickinson}, \\&  {Ferguson} 2001; {Shapley} {et~al.} 2001), contribution to the cosmic star formation rate ({Madau} {et~al.} 1996; {Steidel} {et~al.} 1999; {Adelberger} \\& {Steidel} 2000), rest-frame optical emission lines ({Pettini} {et~al.} 2001), and interaction with the IGM ({Adelberger} {et~al.} 2003). Initially LBGs were thought to be very massive ({Steidel} {et~al.} 1996), but these estimates have been revised following the realization that their rest-frame UV kinematics are dominated by winds rather than gravitational motions ({Franx} {et~al.} 1997; {Pettini} {et~al.} 1998). The masses of luminous $z\\approx 3$ LBGs appear to be $\\sim 10^{10} M_{\\odot}$ ({Pettini} {et~al.} 2001; {Shapley} {et~al.} 2001), a factor of $\\sim 10$ lower than the most massive galaxies today. Such relatively low masses are qualitatively consistent with hierarchical models. LBGs in these models are ``seeds'' marking the highest density peaks in the early Universe, and form the low mass building blocks of massive galaxies in groups and clusters (e.g., {Baugh} {et~al.} 1998).  Despite these advances and the qualitative consistency of theory and observations many uncertainties remain. Perhaps most fundamentally, it is still not clear whether LBGs are the (only) progenitors of today's massive galaxies.  The highly successful Lyman break technique selects objects with strong UV emission, corresponding to galaxies with high star formation rates and a limited amount of obscuration of the stellar continuum.  Galaxies whose light is dominated by evolved stellar populations, and those that are heavily obscured by dust, may therefore be underrepresented in LBG samples (see, e.g., Francis, Woodgate, \\& Danks 1997; Stiavelli et al.\\ 2001; Hall et al.\\ 2001; {Blain} {et~al.} 2002; {Franx} {et~al.} 2003).  Recent advances in near-infrared (NIR) capabilities on large telescopes have made it possible to select high redshift galaxies in the rest-frame optical rather than the rest-frame UV. The rest-frame optical is much less sensitive to dust extinction and is expected to be a better tracer of stellar mass.  The FIRES project ({Franx} {et~al.} 2003) is the deepest ground-based NIR survey to date, with 101.5 hrs of VLT time invested in a single pointing on the Hubble Deep Field South (HDF-S) ({Labb{\\' e}} {et~al.} 2003), and a further 77 hrs on a mosaic of four pointings centered on the foreground cluster \\msclus\\ (van Dokkum et al.\\ 2003; F\\\"orster Schreiber et al.\\ 2004).  We have selected high redshift galaxies in the FIRES fields by their observed NIR colors.  The simple criterion $J_s-K_s>2.3$ efficiently isolates galaxies with prominent Balmer- or 4000\\,\\AA-breaks at $z>2$ (see {Franx} {et~al.} 2003).  This rest-frame optical break selection is complementary to the rest-frame UV Lyman break selection.  We find large numbers of red $z>2$ objects in both fields ({Franx} {et~al.} 2003; {van Dokkum} {et~al.} 2003). Surface densities are $\\sim 3$\\,arcmin$^{-2}$ to $K_s=22.5$ and $\\sim 1$\\,arcmin$^{-2}$ to $K_s=21$, and the space density is 30--50\\,\\% of that of LBGs. Their much redder colors suggest they have higher mass-to-light ($M/L$) ratios, and they may contribute equally to the stellar mass density. Most of the galaxies are too faint in the rest-frame UV to be selected as LBGs. Although the samples are still too small for robust measurements, there are indications that the population is highly clustered ({Daddi} {et~al.} 2003; van Dokkum et al.\\ 2003; {R{\\\" o}ttgering} {et~al.} 2003). The available evidence suggests they could be the most massive galaxies at high redshift, and progenitors of today's early-type galaxies.  Although these results are intriguing, large uncertainties remain. The density and clustering measurements are based on very small areas, comprising less than five percent of the area surveyed for LBGs by Steidel and collaborators. Furthermore, owing to their faintness in the observer's optical, spectroscopic redshifts have been secured for only a handful of objects (van Dokkum et al.\\ 2003). Finally, current estimates of ages, $M/L$ ratios, star formation rates, and extinction are solely based on modeling of broad band spectral energy distributions (SEDs), and as is well known this type of analyis suffers from significant degeneracies in the fitted parameters (see, e.g., {Papovich} {et~al.} 2001; {Shapley} {et~al.} 2001).  Confirmation of the high stellar masses and improved constraints on the stellar populations require spectroscopy in the rest-frame optical (the observer's NIR). Emission lines such as \\oiii\\ $\\lambda 4959,5007$ and the Balmer H$\\alpha$ and H$\\beta$ lines have been studied extensively at low redshift, allowing direct comparisons to nearby galaxies. The H$\\alpha$ line is particularly valuable, as its luminosity is proportional to the star formation rate ({Kennicutt} 1998), and its equivalent width is sensitive to the ratio of current and past star formation activity. When more lines are available metallicity and reddening can be constrained as well.  Finally,  the widths of rest-frame optical lines better reflect the velocity dispersion of the H\\,{\\sc ii} regions than the widths of rest-UV lines, which are very sensitive to outflows and supernova-driven winds (see, e.g., {Pettini} {et~al.} 1998).  In this paper, we present NIR spectroscopy of a small sample of $J_s-K_s$ selected galaxies. Line luminosities, equivalent widths, and linewidths are determined, and the derived constraints on the stellar populations and masses are combined with results from fits to the broad band SEDs. Results are compared to nearby galaxies, and also to LBGs: {Pettini} {et~al.} (1998, 2001) and {Erb} {et~al.} (2003) have studied the rest-frame optical spectra of LBGs in great detail, providing an excellent benchmark for such comparisons.  For convenience we use the term \\orgsf, or \\orgs, for galaxies having $J_s-K_s>2.3$ and redshifts $z\\gtrsim 2$. This term is more general than ``optical-break galaxies'', as it allows for the possibility that in some galaxies the red colors are mainly caused by dust rather than a strong continuum break. {Im} {et~al.} (2002) use the designation Hyper Extremely Red Objects, or HEROs, for galaxies with $J-K\\gtrsim 2$. However, the corresponding rest-frame optical limits are not really ``hyper extreme'', as they would include all but the bluest nearby galaxies: at $z=2.7$, our $J_s-K_s$ limit corresponds to $U-V\\gtrsim 0.1$ in the rest-frame. We use $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$, and $H_0=70$\\,\\kms\\,Mpc$^{-1}$ ({Riess} {et~al.} 1998; {Spergel} {et~al.} 2003). All magnitudes are on the Vega system; AB conversion constants for the ISAAC filters are given in {Labb{\\' e}} {et~al.} (2003). ",
        "conclusions": "Near-infrared spectroscopy with Keck/NIRSPEC has provided new constraints on the nature of red galaxies at redshifts $z>2$.  The results are not unique, as they depend on the assumed functional form of star formation history, and on the (unknown) extinction towards the line emitting gas.  The low H$\\alpha$ equivalent widths imply either significant extinction or a declining star formation rate. For constant star formation models we find total extinction toward H{\\sc ii} regions of $A_V = 2-3$\\,mag, extinction-corrected star formation rates of 200--400 $M_{\\odot}$\\,yr$^{-1}$, ages $1-2.5$\\,Gyr, and stellar masses $1-5\\times 10^{11}\\,M_{\\odot}$. In declining models the star formation rates and extinction are lower but the stellar masses are similar. Measurements of H$\\alpha$ and H$\\beta$ for the same objects will help constrain the extinction toward the line-emitting gas.  The $J_s-K_s$ selected objects extend the known range of properties of early galaxies to include higher masses, mass-to-light ratios, and ages allowing a re-evaluation of the existence of galaxy scaling relations in the early universe. By combining our small sample of \\orgs\\ with a published sample of LBGs we find significant correlations between line width and color and line width and stellar mass. The latter correlation is intruiging, as it has the same form as the ``baryonic Tully-Fisher relation'' in the local universe. There are, however, many uncertainties.  In particular, it is unclear whether the line widths at $z>2$ measure regular rotation in disks, and if they do, whether they sample the rotation curves out to the same radii as in local galaxies. Also, the correlations are currently only marginally significant, and larger samples are needed to confirm them. Finally, the comparison to $z\\approx 3$ LBGs may not be appropriate: initial results from studies at $z\\approx 2$ indicate that their measured properties are markedly different from the $z\\approx 3$ LBGs, with the lower redshift galaxies being more massive and more metal rich (Erb et al.\\ 2003 and private communication). A plausible reason is that the $z\\approx 3$ LBGs were selected to be bright in the rest-frame ultraviolet, which may select lower mass and lower metallicity objects (Erb et al.\\ 2003). It will be interesting to see if the $z\\approx 2$ LBGs deviate from the relations seen in Fig.\\ \\ref{sigcorr.plot}.  It is still unclear how \\orgs\\ fit in the general picture of galaxy formation. They are redder, more metal rich, and more massive than $z\\approx 3$ Lyman break galaxies of the same luminosity, which suggests they will evolve into more massive galaxies. A plausible scenario is that they are descendants of LBGs at $z\\sim 5$ or beyond, in which continued star formation led to a gradual build-up of dust, metals, and stellar mass.  It is tempting to place \\orgs\\ in an evolutionary sequence linking them to Extremely Red Objects (EROs) and massive $K-$selected galaxies at $1<z<2$ (e.g., {Elston}, {Rieke}, \\& {Rieke} 1988; {McCarthy} {et~al.} 2001; {Cimatti} {et~al.} 2002; Glazebrook et al.\\ 2004) and early-type galaxies in the local universe. The apparently smooth sequence of massive galaxies in the $M/L$ versus redshift plane is suggestive of such evolution. However, it is not clear whether galaxies follow ``parallel tracks'' in such plots (i.e., \\org\\ $\\rightarrow$ ERO $\\rightarrow$ early-type, and LBG $\\rightarrow$ spiral) or evolve in complex ways due to mergers or later addition of metal-poor gas (e.g., Trager et al.\\ 2000; van Dokkum \\& Franx 2001; Bell et al.\\ 2004).  The subsequent evolution of \\orgs\\ should also be placed in context with the evolution of the cosmic stellar mass density.  The red galaxies likely contribute $\\gtrsim 50$\\,\\% of the stellar mass density at $z\\approx 2.5$ ({Franx} {et~al.} 2003; {Rudnick} {et~al.} 2003), a similar contribution as that of elliptical galaxies to the local stellas mass density (e.g., {Fukugita}, {Hogan}, \\&  {Peebles} 1998). However, studies of the Hubble Deep Fields suggest that the stellar mass density of the Universe has increased by a factor of $\\sim 10$ since $z\\approx 3$ ({Dickinson} {et~al.} 2003; {Rudnick} {et~al.} 2003), and it is unknown whether \\orgs\\ participated in this rapid evolution.  The high dust content that we infer for \\orgs\\  has important consequences for tests of galaxy formation models that use the abundance of $K$-selected galaxies at high redshift (e.g., Kauffmann \\& Charlot 1998; Cimatti et al.\\ 2003). Our results show that any observed decline in the abundance will at least in part be due to dust: the most luminous $K-$selected galaxies at $0<z<1$ are dust-free elliptical galaxies, but the $K$ magnitudes of red galaxies at $z>2$ are probably attenuated by 1--2 magnitudes. Therefore, tests of galaxy formation models using the ``K20'' technique require a careful treatment of dust, either by correcting the observed abundance for extinction or by incorporating dust in the models.    There are several ways to establish the properties of red $z>2$ galaxies more firmly, and to enlarge the sample for improved statistics. Photometric surveys over larger areas and multiple lines of sight are needed to better quantify the surface density, clustering, and to measure the luminosity function. Furthermore, the high dust content and star formation rates imply that a significant fraction of \\orgs\\ may shine brightly in the submm, as evidenced by the SCUBA detection of the galaxy with the highest dust content in our sample. The star formation may also be detectable with Chandra in stacked exposures, as has been demonstrated recently for LBGs (Reddy \\& Steidel 2004).  Chandra can also provide estimates of the prevalence of AGN in \\orgs; this fraction could be relatively high if black hole accretion rates correlate with stellar mass at early epochs. Spitzer can provide better constraints on the stellar masses, as it samples the rest-frame $K$ band.  Finally, the results in this paper can clearly be improved by obtaining more NIR spectroscopy of red $z>2$ galaxies. Contrary to the situation for LBGs, the sample of \\orgs\\ with measured redshift is rather limited due to their faintness in the rest-frame UV and the resulting difficulty of measuring redshifts in the observer's optical (see van Dokkum et al.\\ 2003; Wuyts et al.\\ in preperation). Furthermore, samples of galaxies with optically-measured redshifts are obviously biased toward galaxies with high star formation rates, low dust content, and/or active nuclei; it may not be a coincidence that the galaxy with the highest dust content is the only one in our sample without Ly$\\alpha$ emission. ``Blind'' NIR spectroscopy is notoriously difficult, but may be the best way of determining the masses, star formation rates, and dust content of \\orgs\\ in an unbiased fashion. It is expected that such studies will be possible for large samples when multi-object NIR spectrographs become available on 8--10m class telescopes."
    },
    "0404/astro-ph0404251_arXiv.txt": {
        "abstract": "{We discuss the physical basis of the statistical mechanics of self-gravitating systems.  We show the correspondance between statistical mechanics methods based on the evaluation of the density of states and partition function and thermodynamical methods based on the maximization of a thermodynamical potential (entropy or free energy). We address the question of the thermodynamic limit of self-gravitating systems, the justification of the mean-field approximation, the validity of the saddle point approximation near the transition point, the lifetime of metastable states and the fluctuations in isothermal spheres. In particular, we emphasize the tremendously long lifetime of metastable states of self-gravitating systems which increases exponentially with the number of particles $N$ except in the vicinity of the critical point. More specifically, using an adaptation of the Kramers formula justified by a kinetic theory, we show that the lifetime of a metastable state scales as $e^{N\\Delta s}$ in microcanonical ensemble and $e^{N\\Delta j}$ in canonical ensemble, where $\\Delta s$ and $\\Delta j$ are the barriers of entropy and free energy $j=s-\\beta \\epsilon$ (per particle) respectively. The physical caloric curve must take these metastable states (local entropy maxima) into account. As a result, it becomes multi-valued and leads to microcanonical phase transitions and ``dinosaur's necks'' (Chavanis 2002b, Chavanis \\& Rieutord 2003). The consideration of metastable states answers the critics raised by D.H.E. Gross [cond-mat/0307535/0403582]. ",
        "introduction": "\\label{sec_introduction}   The statistical mechanics of self-gravitating systems has a long history starting with the seminal papers of Antonov (1962) and Lynden-Bell \\& Wood (1968). A statistical mechanics approach is particularly relevant to describe the late stages of ``small'' groups of stars ($N\\sim 10^{6}$), such as globular clusters, which evolve under the influence of stellar encounters (``collisional'' relaxation). Apart from astrophysical applications, the statistical mechanics of stellar systems is of great interest in physics because it differs in many respects from that of more familiar systems with short-range interactions (Padmanabhan 1990).  In particular, for systems with long-range interactions, the thermodynamical ensembles are not equivalent, negative specific heats are allowed in the microcanonical ensemble (but not in the canonical ensemble) and metastable equilibrium states can have tremendously long lifetimes making them of considerable interest.  Two types of approaches have been developed to determine the statistical equilibrium state of a self-gravitating system. In the {\\it thermodynamical approach}, one determines the { most probable} distribution of particles by maximizing the Boltzmann entropy at fixed mass and energy in the microcanonical ensemble or by minimizing the free energy $F=E-TS$ at fixed mass and temperature in the canonical ensemble (Lynden-Bell \\& Wood 1968, Katz 1978, Chavanis 2002a). This approach is the simplest and the most illuminating. In addition, it is directly related to kinetic theories (based on the Landau or on the Fokker-Planck equation) for which the Boltzmann entropy (or the Boltzmann free energy) plays the role of a Lyapunov functional and satisfies a H-theorem. Alternatively, in the {\\it statistical mechanics approach}, one starts from the density of states or partition function, transforms it into a functional integral and uses a saddle point approximation valid in a properly defined thermodynamic limit (Horwitz \\& Katz 1978, de Vega \\& Sanchez 2002, Katz 2003).  In the first part of this paper, we discuss the connexion between these two procedures. We remain at a heuristic level, stressing more the physical ideas than the mathematical formalism. In Sec. \\ref{sec_fermions}, we introduce the entropy by a combinatorial analysis. In order to regularize the problem at short distances, we consider either the case of self-gravitating fermions or the case of self-gravitating particles with a soften potential. We also discuss the thermodynamic limit of the classical and quantum self-gravitating gas. In Sec. \\ref{sec_connexion}, we show the relation between the density of states $g(E)$ and the entropy functional $S[f]$ and between the partition function $Z(\\beta)$ and the free energy functional $J[f]=S[f]-\\beta E[f]$. In the thermodynamic limit, the saddle point approximation amounts to maximizing the entropy at fixed mass and energy (microcanonical ensemble) or to minimizing the free energy at fixed mass and temperature (canonical ensemble). In Sec. \\ref{sec_corr}, we discuss the notion of canonical and microcanonical phase transitions in self-gravitating systems.  We perform the (standard) horizontal and (less standard) vertical Maxwell constructions and discuss the validity of the saddle point approximation near the transition point for finite $N$ systems. These results (e.g., microcanonical first order phase transitions) are relatively new in statistical mechanics and still subject to controversy (Gross 2003,2004). Therefore, we provide a relatively detailed discussion of these issues.  In the second part of the paper, we emphasize the importance of metastable states in astrophysics and show how they can be taken into account in the statistical approach. In Sec. \\ref{sec_persistence}, we use the Kramers formula to estimate the lifetime of a metastable state. We show that the lifetime of a metastable state scales as $e^{N\\Delta s}$ in microcanonical ensemble and $e^{N\\Delta j}$ in canonical ensemble, where $\\Delta s$ and $\\Delta j$ are the barriers of entropy and free energy $j=s-\\beta \\epsilon$ (per particle) respectively. Therefore, the typical lifetime of a metastable state scales as $e^{N}$ except in the vicinity of the critical point $E_{c}$ (Antonov energy) or $T_{c}$ (Emden-Jeans temperature). We explicitly compute the barriers of entropy and free energy close to the critical point for classical self-gravitating particles (stars). The very long lifetime of metastable states, scaling as $e^{N}$, was pointed out by Chavanis \\& Rieutord (2003) and the difficulty of a stellar system to overcome the entropic barrier and collapse was qualitatively discussed in Chavanis \\& Sommeria (1998). We here improve these arguments by developing a theory of fluctuations in isothermal spheres, following the approach of Katz \\& Okamoto (2000). We also determine how finite $N$ effects affect the collapse temperature and the collapse energy. Finally, in Sec. \\ref{sec_kramers}, we derive a Fokker-Planck equation for the evolution of the distribution of energies $P(E,t)$ in the canonical ensemble and make contact with the standard Kramers problem. We determine the typical lifetime of a metastable state by calculating the escape time accross a barrier of free energy. ",
        "conclusions": "\\label{sec_conclusion}  In this paper, we have completed previous investigations concerning the statistical mechanics of self-gravitating systems in microcanonical and canonical ensembles. The microcanonical ensemble is the proper description of isolated Hamiltonian systems such as globular clusters (Binney \\& Tremaine 1987). The canonical ensemble is relevant for systems in contact with a heat bath of non-gravitational origin. It is also the proper description of stochastically forced systems such as self-gravitating Brownian particles (Chavanis, Rosier \\& Sire 2002).  We have justified the mean-field approximation, in a proper thermodynamic limit $N\\rightarrow +\\infty$ with $\\eta=\\beta GMm/R$ and $\\epsilon=ER/GM^{2}$ fixed, from the equilibrium BBGKY hierarchy. In this thermodynamic limit, the equilibrium state is determined by a maximization problem: the maximization of entropy at fixed mass and energy in the microcanonical ensemble and the minimization of free energy at fixed mass and temperature in the canonical ensemble. This determines the {\\it most probable} macroscopic distribution of particles at equilibrium. This can also be seen as a saddle point approximation in the functional integral formulation of the density of states and partition function. We have shown that the saddle point approximation is less and less accurate close to the transition point since the condition $N|E-E_{t}|\\gg 1$ (in microcanonical ensemble) or $N|T-T_{t}|\\gg 1$ (in canonical ensemble) must be satisfied.  We have also argued that the lifetime of metastable states (local entropy maxima) scales as ${\\rm exp}(N)$ due to the long-range nature of the interaction. Therefore, the importance of these metastable states is considerable and they cannot be simply ignored. Metastable states are in fact {\\it stable} and they correspond to observed structures in the universe such as globular clusters. The preceding estimate must, however, be revised close to the critical point. By solving a Fokker-Planck equation, we have shown the the lifetime of metastable states is given by the Kramers formula involving the barrier of entropy or free energy. These barriers have been calculated exactly close to the Antonov energy $E_{c}$ (in microcanonical ensemble) and close to the Jeans-Emden temperature $T_{c}$ (in canonical ensemble). We have obtained the estimates $t_{life}\\sim {\\rm exp}\\lbrace 1.726\\ N (\\Lambda_{c}-\\Lambda)^{3/2}\\rbrace$ (in microcanonical ensemble) and $t_{life}\\sim {\\rm exp}\\lbrace 0.339\\ N (\\eta_{c}-\\eta)^{3/2}\\rbrace$ (in canonical ensemble) so that the lifetime decreases as we approach $E_c$ or $T_{c}$. This implies that the collapse will take place slightly above $E_c$ or $T_{c}$ at an energy $\\Lambda_l= \\Lambda_{c}(1-2.077\\ N^{-2/3})$ or temperature $\\eta_{l}=\\eta_{c}(1-0.816\\ N^{-2/3})$. Similar conclusions have been reached by Katz \\& Okamoto (2000). Yet, these predictions do not seem to be consistent with the Monte Carlo simulations of de Vega \\& Sanchez (2002), although they find that the collapse indeed takes place slightly before the critical point. Independent simulations are under preparation to check that point.  Finally, a part of our discussion was devoted to answer the critics raised by Gross (2003,2004) in recent comments. This author argues that the microcanonical entropy $S_{micro}(E)$ and the microcanonical temperature $\\beta_{micro}(E)$ must be single valued. This is true in a strict sense, but the problem is richer than that because of the existence of long-lived metastable states. Therefore, the {\\it physical} caloric curve/series of equilibria $\\beta(E)$ is multi-valued and leads to ``dinosaur's necks'' and special ``microcanonical phase transitions'' (Chavanis 2002). This is specific to systems with long-range interactions in view of the long lifetime of metastable states (local entropy maxima). These results have stimulated a general classification of phase transitions by Bouchet \\& Barr\\'e (2004). Microcanonical phase transitions (as in Fig. \\ref{le5}) have not been fully appreciated by Gross and his collaborators because their studies (e.g., Votyakov et al. 2002) consider a {\\it large} small-scale cut-off for which the caloric curve looks like Fig. \\ref{fel} and is univalued. If these authors reduce their small-scale cut-offs, they will see ``dinosaurs'' appear!"
    },
    "0404/astro-ph0404584_arXiv.txt": {
        "abstract": "Since the beginning of precise Doppler surveys, which have had stunning success in detecting extrasolar planetary companions, one surprising enigma has emerged: the relative paucity of spectroscopic binaries where the secondary mass lies in between the stellar and planetary mass regime. This gap in the mass function for close-in ($a<3-4$~AU) companions to solar-type stars is generally referred to as the ``Brown Dwarf Desert''. Here we report the detection of a companion to HD~137510 (G0IV), with a minimum mass of $26~{\\rm M}_{\\rm Jupiter}$, moving in an eccentric orbit ($e=0.4$) with a period of $798$ days and an orbital semimajor axis of $1.85$~AU. The detection is based on precise differential radial velocity (RV) data obtained by the McDonald Observatory and Th\\\"uringer Landessternwarte Tautenburg planet search programs. ",
        "introduction": "It was the common expectation that, with the increase in the precision of RV measurements, the detected secondary masses of spectroscopic binaries will gradually decrease, first into the brown dwarf mass range (${\\rm M}<80~{\\rm M}_{\\rm Jup}$) and eventually into the planetary mass regime of roughly ${\\rm M}<13~{\\rm M}_{\\rm Jup}$ (below the deuterium burning limit). Quite the opposite turned out to happen. The pioneering long-term Doppler survey of Campbell, Walker, \\& Yang (1988), and later also in the southern hemisphere by Murdoch, Hearnshaw, \\& Clark~(1993), did not reveal any brown dwarf companions to nearby solar-type stars, despite their sufficient sensitivity. There appeared to be an abrupt drop in the frequency of companions at the hydrogen burning limit of $\\approx 0.075~{\\rm M}_{\\odot}$. The first very low-mass companion ($m \\sin i = 11~{\\rm M}_{\\rm Jup}$) detected by the RV technique is the companion to HD~114762 (Latham et al.~1989). The discoverers classified the object as a brown dwarf companion, but according to Cochran, Hatzes, \\& Hancock~(1991) this companion might be in fact a late M dwarf due to a possible low value of the inclination angle $i$. In the years after the discovery of the first unambiguous extrasolar planet orbiting a main-sequence star by Mayor \\& Queloz (1995), a total of more than a $100$ extrasolar planets have been accumulated by several precise Doppler searches all around the globe. Still, with a few exceptions (HD~127506: Mayor, Queloz, \\& Udry~1998; HD~29587 \\& HD~140913: Mazeh, Latham, \\& Stefanik~1996; HD~168443 (second companion): Marcy \\& Butler~2001; HD~184860: Vogt et al.~2002), the part of the mass function for spectroscopic secondaries, in between the minimum hydrogen burning mass and the deuterium burning limit of ${\\rm M}\\approx13~{\\rm M}_{\\rm Jup}$, remained surprisingly sparsely populated. A notable exception is HD~10697, where Vogt et al.~(2000) detected a companion with $m \\sin i = 6.35~{\\rm M}_{\\rm Jup}$, while Zucker \\& Mazeh (2000) determined a true mass for the secondary of $40~{\\rm M}_{\\rm Jup}$ after combining the spectroscopic solution with $Hipparcos$ astrometry. Udry et al.~(2002) found 4 companions with minimum masses close to the planet/brown dwarf border, which could be either ``superplanets'' or real brown dwarfs. However, most of the candidate brown dwarf secondaries considered in the Halbwachs et al.~(2000) study turned out to be stellar mass companions, again after combination with $Hipparcos$ astrometry.  This suggests the existence of two distinctive binary formation processes for certain mass-ratios. The gap in the mass function between the stellar and planetary mass range for close binaries is generally dubbed the ``Brown Dwarf Desert'' (this desert does not appear to exist at wide separations, e.g. Gizis et al. 2001; Neuh\\\"auser \\& Guenther 2004). The gap is particularily obvious as a sudden drop at the deuterium burning limit in the substellar mass function for close companions. While planets are relatively common as companions to solar-type stars, brown dwarfs are rare. This fact is especially surprising since brown dwarf companions are much easier to detect, even for lower precision surveys, due to the larger RV amplitudes. Moreover, these objects cover a more than 4 times larger mass range than planetary companions, still their detected numbers per mass bin is extremely low compared to lower and higher mass companions.  In this work we present the discovery of another possible ``oasis'' in this desert, a brown dwarf orbiting the G-type star HD~137510. ",
        "conclusions": "We have presented observational evidence for the substellar nature of the companion to HD~137510. {\\it The question, why this star constitutes one of the very rare oases in the brown dwarf desert, remains unanswered.} However, due to the well established parameters of the primary and the orbit, this new system might help to understand binary formation processes for these mass-ratios at small orbital separations."
    },
    "0404/astro-ph0404317_arXiv.txt": {
        "abstract": "{ We present a prospective analysis of a combined cosmic shear and cosmic microwave background data set, focusing on a Canada France Hawaii Telescope Legacy Survey (CFHTLS) type lensing survey and the current WMAP-1 year and CBI data. We investigate the parameter degeneracies and error estimates of a seven parameters model, for the lensing alone as well as for the combined experiments. The analysis is performed using a Monte Carlo Markov Chain calculation, allowing for a more realistic estimate of errors and degeneracies than a Fisher matrix approach. After a detailed discussion of the relevant statistical techniques, the set of the most relevant 2 and 3-dimensional lensing contours are given. It is shown that the combined cosmic shear and CMB is particularly efficient to break some parameter degeneracies. The principal components directions are computed and it is found that the most orthogonal contours between the two experiments are for the parameter pairs $(\\Omega_m,\\sigma_8)$, $(h,n_s)$ and $(n_s,\\alpha_s)$, where $n_s$ and $\\alpha_s$ are the slope of the primordial mass power spectrum and the running spectral index respectively. It is shown that an improvement of a factor $2$ is expected on the running spectral index from the combined data sets. Forecasts for error improvements from a wide field space telescope lensing survey are also given. ",
        "introduction": "The Canada-France-Hawaii Telescope Legacy Survey \\footnote{http://www.cfht.hawaii.edu/Science/CFHTLS/} (CFHTLS) is a long term wide field imaging project that started in early 2003 and should be completed by 2008. The French and Canadian astronomical communities will spend about 500 CFHT nights to carry out imaging surveys with the new Megaprime/Megacam instrument recently mounted at the CFHT prime focus. About 160 nights will focus on the \"CFHTLS-Wide\" survey that will cover 170 deg$^2$, spread over 3 uncorrelated patches of $7^o \\times 7^o$ each,  in $u^*,g',r',i',z'$ bands, with typical exposure times of about one hour per filter.  The \"CFHTLS-Wide\" survey design and observing strategy are similar to the {\\sc Virmos}-Descart cosmic shear survey \\footnote{http://terapix.iap.fr/cplt/oldSite/Descart/} but it will have a sky coverage 20 times larger.  It is widely seen as a typical second generation cosmic shear survey.  The exploration of weak gravitational distortion produced by the large scale structures of the universe over field of views as large as \"CFHTLS-Wide\" has an enormous potential for cosmology.  Past experiences based on first generations cosmic shear surveys (see for example reviews in \\cite{VWM03,REFRE03}) have demonstrated they can constrain the dark matter properties ($\\sigma_8$, $\\Omega_m$ and the shape of the dark matter power spectrum) from a careful investigation of the ellipticity induced by weak gravitational shear on distant galaxies. For example, the  most recent cosmic shear results from the {\\sc Virmos}-Descart survey (\\cite{VWMH04}) lead to the conservative limits $\\sigma_8 =0.85 \\pm0.15$ (99\\% C.L.) and $\\Omega_m=0.3\\pm 0.15$ (99\\% C.L.), which means an expected accuracy of $\\approx$1-3\\% can be expected with the \"CFHTLS-Wide\" for the same set of cosmological parameters.  The CFHTLS-Wide will also explore a broader wavenumber range (10$^5$-10$^2$) than {\\sc Virmos}-Descart and will extend to  linear scale, that will considerably ease cosmological interpretation of weak lensing data. The second generation surveys will therefore permit to investigate more thoroughly different cosmological models, taking into account a broad range of cosmological parameters. For instance, the CFHTLS as a probe dark energy evolution was stressed out by \\cite{VWBENAB03}.  The full scientific outcome of the cosmic shear data from the \"CFHTLS-Wide\" will only be complete with a joint analysis with other data sets, like Type Ia Supernovae, galaxy redshift surveys, Lyman-alpha forest, or CMB observations.  \\cite{CONTALDHOEK03} have used the Red Cluster Sequence (RCS)  cosmic shear survey together with the CMB data. It was shown that the $\\Omega_m,\\sigma_8$ degeneracies for lensing and CMB are nearly orthogonal, which makes this set of parameters particularly relevant for such combined analysis (\\cite{VW02}). The search for orthogonal parameter degeneracies between different observations is one of the most important aspects of parameters measurements.  $\\,$\\cite{ISHAK03} recently argued that joint CMB-cosmic shear surveys provide an optimal data set to explore the amplitude of the running spectral index and probe inflation models. They used a Fisher-Matrix analysis on  WMAP+ACBAR+CBI and a \"reference survey\". Their simulated survey covers about 400 deg$^2$ with a depth corresponding to a galaxy number density of lensed sources of about 60 arc-min$^{-2}$, and they restricted their analysis to 3000$>l>$20. They found that several parameters can be significantly improved (like $\\sigma_8$, $\\Omega_mh^2$, $\\Omega_{\\Lambda}$)  and in particular that both the spectral index $n_s$ and the running spectral index $\\alpha_s$ errors are reduced by a factor of 2. Their encouraging results show that joint CMB and weak lensing data may provide interesting insights on inflation models. Here, we investigate the 2-dimensional structure of the parameter degeneracies between lensing and CMB data sets, and look for the expectation with a \"CFHTLS-Wide\"-like survey design.  To explore the smaller scales probed by the \"CFHTLS-Wide\", which will provide cosmic shear information down to 20 arc-seconds, it is preferable to avoid prior assumptions regarding the Gaussian nature of the underlying distribution, and to discard a Fisher matrix analysis. We used in this work the so-called Markov Chain Monte Carlo (MCMC) method. The MCMC computing time linearly scales with the number of parameters and eases the exploration of a large sample of parameters and a broad range of values for each. \\cite{CONTALDHOEK03} already used this approach with the RCS survey to map the $\\Omega_m,\\sigma_8$ parameter space, but marginalised over a small set of cosmological parameters. The goal of this present work  is  to  map the parameter space that describes cosmological models in order to extract series of parameter combinations that would minimise intersections of CMB and cosmic shear degeneracy tracks. Compared to the Fisher-Matrix approach which produces ellipses only, MCMC provides more details of the parameter space and eventually a more realistic estimate of error improvements of the joint analyses.  The paper is organised as follows: Section 2 introduces the gravitational lensing and defines the cosmic shear fiducial data used and the parameter space investigated. Section 3 gives the details of our MCMC calculations, limitations and convergence criteria. Section 4 shows the MCMC results from the cosmic shear alone, assuming a lensing survey similar to the CFHTLS. In Section 5 we present the results of the parameter degeneracies analysis on the combined cosmic shear and cosmic microwave background observations. The assumptions made and the results obtained are discussed in section 6 and we conclude in section 7. ",
        "conclusions": "The CMB/cosmic shear complementarity opens good prospectives for the determinations of cosmological parameters by combining CMB and cosmic shear data sets. In fact, even for CFHTLS, whose contours are, in general, noticeably larger than the WMAP+CBI ones (Fig. \\ref{joint2d}), we predict non neglectable gains. Figure \\ref{snapcmb} shows what can be expected with future space telescope data. \\begin{table} \\caption{Cosmic shear : Wide field space telescope illustration specifications.} \\label{snaptab} \\begin{center} \\begin{tabular}{ll} \\hline \\hline Size of the survey: &A=$1000\\,{\\rm deg} ^2$ \\\\ Density of galaxies: & $n_g=50\\,{\\rm arcmin}^{-2}$ \\\\ Intrinsic ellipticity dispersion:& $\\sigma_\\epsilon=0.3$ \\\\ Scales probed:& $0.6\\,'<\\theta<5\\deg$\\\\ \\ & $40<l<18000$\\\\ \\hline \\hline \\end{tabular} \\end{center} \\end{table} \\begin{figure} \\centering \\includegraphics [height=6cm] {snapcmbbm.ps} \\caption{68\\% and 95\\% C.L. for the 4 most orthogonal cases found in section 5. Blue is WMAP+CBI data and red shows predictions for the wide field space telescope - cosmic shear parameters of Table (\\ref{snaptab}) (99 \\% C.L. is also shown in this case).} \\label{snapcmb} \\end{figure} These results were produced with a cosmic shear Fisher matrix analysis, using Eq. (\\ref{fisher}) and the fiducial model of Table \\ref{fiduc} (except for the redshift of the sources which was moved to $z_s=1.1$). For this illustration, the data covariance matrix of Eq. (\\ref{bigcov}) was computed for the configuration shown in Table \\ref{snaptab}, which is close to the SuperNova Accelerator Probe / Joint Dark Energy Mission (SNAP/JDEM) \"Wide+\" case of \\cite{REFSNAP3}. The CMB ellipses are plotted from the parameters covariance matrix found with our WMAP+CBI chains.  In summary, we found the best constrained parameters by 2-point cosmic shear correlation functions to be $\\sigma_8\\,\\Omega_m^{0.5}$ and $\\Gamma\\,n_s^{0.6}$ (with $\\Gamma=\\Omega_m\\,h$). We have shown that 2-dimensional degeneracies defined by these parameters plus another one defined by $n_s$ and $\\alpha_s$ are orthogonal to CMB degeneracies. Due to this CMB/cosmic shear complementarity, current weak lensing surveys, such  as the CFHTLS, already have the potential to improve the precision on several cosmological parameters. In particular, a better knowledge of $\\alpha_s$ will have an impact on inflationary scenarios. The crucial information provided by the cosmic shear comes from the small scales it probes. Thus, it provides an additional possibility, along with galaxy redshift surveys and Lyman-$\\alpha$ forest, to combine with CMB data."
    },
    "0404/astro-ph0404499_arXiv.txt": {
        "abstract": "The 511 keV emission from the Galactic Bulge observed by INTEGRAL/SPI could be the product of light (1-100 MeV) annihilating dark matter particles. In order to distinguish between annihilating light dark matter scenarios and more conventional astrophysical sources for the bulge emission, we propose to test the light dark matter hypothesis on the Sagittarius Dwarf Galaxy, a close-by galaxy dominated by dark matter. The predicted flux being in the SPI sensitivity range, the detection of a substantial 511 keV emission line from this galaxy would provide a strong evidence for the light dark matter hypothesis.  During the two Galactic Center Deep Exposures performed in 2003, the Sagittarius Dwarf Galaxy was in the field of view of several INTEGRAL observations. In this paper we present preliminary results of the analysis of these data. ",
        "introduction": "\\begin{figure*} \\centering \\includegraphics[width=0.8\\linewidth]{SDG.eps} \\caption{Surface density map of the Sagittarius Dwarf Galaxy. On the right, Integral exposure map.\\label{fig:expomap}} \\end{figure*}  The \\unmarkedfootnote{ corresponding author: b.cordier@cea.fr \\\\ $^{3}$ Institut d'Astrophysique de Paris, 98 bis Boulevard Arago, 75014 Paris, France \\\\ $^{4}$ F\\'ed\\'eration de Recherche Astroparticule et Cosmologie, % Coll\\`ege de France, 11 Place Marcellin Berthelot, 75231 Paris, France } 511 keV emission from the Galactic bulge \\citep{Jean03} observed by INTEGRAL/SPI \\citep{Vedrenne03,Schanne02,Attie03,Roques03} could be the signature of light (1-100 MeV) dark matter particles \\citep{Casse04}. Such particles could annihilate throughout the Galactic bulge into positrons (and electrons) which, after losing energy, annihilate into 511 keV gamma-rays \\citep{Boehm04}. Note that the SPI observation implies an injection rate of about $10^{43}$ positrons per second within the Galactic bulge. It seems difficult to account for such an injection rate and to reproduce the morphology observed by INTEGRAL/SPI by standard astrophysical scenarios such as Type Ia supernovae \\citep{Schanne04}. However astrophysical scenarios cannot be ruled out and more evidence is needed to further motivate the light dark matter annihilation hypothesis. In this context, \\cite{Hooper04} have proposed to test the light dark matter hypothesis on the Sagittarius Dwarf Galaxy (SDG), a closeby galaxy dominated by dark matter. A large dark matter annihilation rate is predicted from this galaxy and as it contains comparatively few stars, the detection of a substantial 511 keV emission line from SDG would provide a strong evidence for the light dark matter hypothesis. Indeed SGR is a good place to search for the effect of pure dark matter since star formation has virtually stopped due to the lack of gas and thus supernovae and hypernovae should be essentially absent. Assuming that:  \\begin{itemize} \\item all the 511 keV photons observed in the bulge of our Galaxy are the signature of dark matter annihilation \\item the spatial distribution of the SDG diffuse halo gas is less than 2 degrees \\item the shape of the dark matter halo profile for the SDG is similar to that of Draco, another dwarf galaxy \\item the hydrogen density in the SDG halo is of the order of ~ $10^{-4}~cm^{-3}$ \\item the annihilation occurs directly on dust grains \\end{itemize}  \\cite{Hooper04} found a predicted flux at 511 keV from SDG: (1-7)$\\times10^{-4}~ph~cm^{-2}~s^{-1}$ . This work presents attempts to search for light dark matter annihilation in the SDG with SPI/INTEGRAL. ",
        "conclusions": "After two Galactic Center Deep Exposures, we have not detected a 511 keV emission from SDG, but this non-detection cannot exclude the light dark matter hypotheses. The SPI/INTEGRAL effective observation time is today not sufficient to reach the adequate sensitivity. In the following months and years, by accumulating the data, the SPI sensitivity in the direction of the SDG will reach the predicted fluxes. In this study we considered the 511 keV source as a point like source (less than 2 degrees), in the future we will search for different source extent. Moreover two parameters are very important for the prediction of the 511 keV emission as a signature of the light dark matter: the hydrogen density within the SDG and the dust fraction. Future observations from FUSE will provide a more accurate value of the gas density. With more SPI data and more accurate parameters, the observation of the SDG in the 511 keV energy range could permit to put a relevant constrain on the mass of the light dark matter."
    },
    "0404/astro-ph0404450_arXiv.txt": {
        "abstract": "{High spatial resolution observations in the 1 to 3.5\\micron\\/ region of the Galactic Center source known historically as \\object{IRS~13} are presented. They include ground-based adaptive optics images in the H, Kp (2.12/0.4\\micron) and L bands, HST-NICMOS data in filters between 1.1 and 2.2\\micron, and integral field spectroscopic data from BEAR, an Imaging FTS, in the \\ion{He}{i} 2.06\\micron\\/ and the Br$\\gamma$ line regions. Analysis of all these data provides a completely new picture of the main component, \\object{IRS~13E}, which appears as a cluster of seven individual stars within a projected diameter of $\\sim$\\,0.5\\arcsec\\ (0.02 pc). The brightest sources, 13E1, 13E2, 13E3 which is detected as a binary, and 13E4, are all massive stars of different type. The star 13E1 is a luminous, blue object, with no detected emission line. 13E2 and 13E4 are two hot, high-mass emission line stars, 13E2 being at the WR stage and 13E4 a massive O-type star. In contrast, 13E3A and B are extremely red objects, proposed as other examples of dusty WR stars, like \\object{IRS~21} (Tanner \\etal \\cite{tanner}). All these sources have a common westward proper motion (Ott \\etal \\cite{ott2}) indicating they are bounded.  Two other sources, detected after deconvolution of the AO images in the H and Kp bands, are also identified. One, that we call 13E5, is a red source similar to 13E3A and B, while the other one, 13E6, is probably a main sequence O~star in front of the cluster.  Considering this exceptional concentration of comoving massive hot stars, IRS~13E is proposed as the remaining core of a massive star cluster, which could harbor an intermediate-mass black hole (IMBH) (Portegies Zwart \\& McMillan \\cite{zwart2}) of $\\sim$\\,1300\\,M$_{\\odot}$. This detection plays in favor of a scenario, first suggested by Gerhard (\\cite{gerhard}), in which the helium stars and the other hot stars in the central parsec originate from the stripping of a massive cluster formed several tens of pc from the center. This cluster would have spiraled towards \\SgrA, and IRS~13E would be its remnant. Furthermore, IRS~13E might be the second black hole needed according to a model by Hansen \\& Milosavljevi\\'c (\\cite{hansen}) to drag massive main-sequence stars, in the required timescale, very close to the massive black hole.  The detection of a discrete X-ray emission (Baganoff \\etal \\cite{baganoff}) at the IRS~13 position (within the positional accuracy) is examined in this context. ",
        "introduction": "In the early images of the central region of the GC recorded in the near infrared, at the best seeing-limited resolution, several bright point sources dominate the $\\sim$\\,20\\arcsec$\\times$20\\arcsec\\/ field centered on \\SgrA. With the radical improvement of angular resolution through multiple short exposures with shift-and-add (SHARP camera, Eckart \\etal \\cite{eckart}) or speckle techniques (Ghez \\etal \\cite{ghez}), with the images from the NICMOS cameras on board HST (Stolovy \\etal \\cite{stolovy}), to the advent of adaptive optics correctors behind large telescopes as NAOS (Ott \\etal \\cite{ott3}), a more complex vision of this crowded field has emerged. At a spatial resolution in the best case of 0.05\\arcsec\\/ (Ghez \\etal \\cite{ghez}) new, faint stellar sources appear and the early-identified sources are often resolved into several components. Therefore, the observed spectra, generally made at a resolution limited at best by a slit not less than 1\\arcsec\\/ wide, are actually composites of emission from stellar objects of different spectral type as well as from emission lines from the surrounding interstellar medium (ISM). The stellar type identified from these spectra can be wrong if attributed to a single source, which can erroneously look unusual. This care led to the paper on the revised identifications of the central cluster of \\ion{He}{i} stars by Paumard \\etal (\\cite{paumard}) made from slitless integral field spectroscopy, with BEAR an imaging Fourier Transform Spectrometer (Maillard \\cite{maillard}). This instrument makes it possible to obtain near-infrared spectroscopy at the seeing-limited resolution of Mauna Kea (i.e. $\\simeq$\\,0.5\\arcsec), which represents a significant improvement.  However, this resolution is not sufficient for such a crowded field. Therefore, the data were compared to adaptive optics (AO) images in the K band of the same field, with a spatial resolution of $\\sim$\\,0.15\\arcsec, in order to check whether the emission sources at the angular resolution of seeing were single stars or not.  Among the sources studied in this paper the object historically named IRS~13, is a typical example. Located approximately 3.6\\arcsec\\/ south-west of \\SgrA, at the edge of the circular structure of the \\object{SgrA~West} HII region called the Minicavity, it appears in early works at near-infrared wavelengths, in J, H, K (Rieke \\etal \\cite{rieke}) and L (Allen \\& Sanders \\cite{allen}), as a bright spot. However, from a discussion on the sources of energy at the Galactic Center, Rieke \\etal (\\cite{rieke}) note \\lq\\lq there are separate luminosity sources in the core of the 10\\micron\\/ sources 1, 9 and 13\\rq\\rq\\/, because, particularly for IRS~13, the energy distribution from 1 to 5\\micron\\/ presents a sudden, steep increase beyond 3\\micron. A first work at subarcsecond resolution, by lunar occultation in the K band (Simon \\etal \\cite{simon}), indicates that IRS~13 resolves into a pair of sources, separated by $\\sim$\\,1.2\\arcsec, which were henceforth designated as IRS~13E and 13W.  In the photometric survey of Ott \\etal (\\cite{ott}) from SHARP imaging data (Eckart \\etal \\cite{eckart}), at a limit of resolution of $\\sim$\\,0.15\\arcsec\\/ obtained by deconvolution, two equally bright components (K$_{mag}$\\,=\\,10.26) are reported, IRS~13E1 and 13E2, with a separation of 0.2\\arcsec.  Paumard \\etal \\cite{paumard}) published the first AO map of IRS~13E in the K band, extracted from a larger image of the central region obtained with the CFHT-AO system (Lai \\etal \\cite{lai}). We noted that a fainter third source, we called \\object{IRS~13E3}, was also present, forming a kind of equilateral triangle with 13E1 and 13E2. In parallel, Cl\\'enet \\etal (\\cite{clenet}) published a photometric analysis of the same data.  For IRS~13E they reported three components in the K band, noted 13E1, 13E2 and 13N. From the given offset coordinates the source called 13N in this paper is not exactly coincident with the source we had previously called 13E3 (Paumard \\etal \\cite{paumard}).  In the meantime, spectroscopic works on the stellar population of the central parsec were conducted. Several spectra of IRS~13 (Blum \\etal \\cite{blum95}, Libonate \\etal \\cite{libo}, Tamblyn \\etal \\cite{tamblyn}), and specifically of IRS~13E ( Krabbe \\etal \\cite{krabbe}, Genzel \\etal \\cite{genzel96}, Najarro \\etal \\cite{najarro}), and of \\object{IRS~13W} (Krabbe \\etal \\cite{krabbe}) have been published. They cover mostly the K band, i.e. all or part of the 1.95 -- 2.45\\micron\\/ range, and one the 1.57 -- 1.75\\micron\\/ region of the H band (Libonate \\etal \\cite{libo}).  From these spectra IRS~13W is unambiguously a cool star with the strong vib-rot CO signatures at 2.3\\micron.  Emission lines were detected on IRS~13E, typical of a luminous, hot star: strong \\ion{He}{i} 2.058, 2.112\\micron, Br$\\gamma$ line and other Brackett lines up to Br12, plus weaker lines of [\\ion{Fe}{ii}], [\\ion{Fe}{iii}], and a weak emission at 2.189\\micron\\/ attributed to \\ion{He}{ii}.  The 7--10 line is the most intense of this ion in the K band, and is supposed unblended, while the 8--14 line cannot be distinguished from Br$\\gamma$ of which it is separated by less than 2\\,cm$^{-1}$. The \\ion{He}{ii} line if detected is a precious type indicator in the classification of WR stars (Tamblyn \\etal \\cite{tamblyn}, Figer \\etal \\cite{figer97}).  From these spectral characteristics Libonate \\etal (\\cite{libo}) concluded \\lq\\lq the IRS~13 spectrum bears a strong resemblance to the low-resolution K-band spectra of P~Cygni (a LBV) and the AF source\\rq\\rq\\/. Krabbe \\etal (\\cite{krabbe}) proposed IRS~13E as a WN9 star.  From a K-band spectrum of the Minicavity Lutz \\etal (\\cite{lutz}) attributed to [\\ion{Fe}{iii}] several emission lines (2.145, 2.218, 2.242, 2.348\\micron). Images at 2.218\\micron\\/ show that this line is present all over the Minicavity, with a particular enhancement at the position of IRS~13.  From NICMOS data in the F164N filter (Stolovy \\etal \\cite{stolovy}) the lower ionized Fe species, [\\ion{Fe}{ii}] is detected in emission in the central ISM, particularly strong at the edge of the Minicavity, but remarkably absent at the position of IRS~13E. As mentioned in this paper the ionization condition for [\\ion{Fe}{iii}] requires 16.2\\,eV while only 7.2\\,eV are required for [\\ion{Fe}{ii}]. Hard ionization radiation is likely originating from IRS~13E. Hence, it can concluded that these iron lines are not of stellar origin but are present in the surrounding ionized gas.  From the extraction of the \\ion{He}{i} 2.058\\micron\\/ line profile at high spectral resolution from the BEAR data, Paumard \\etal (\\cite{paumard}) concluded that one of the three sources identified as forming IRS~13E, instead of a LBV-type star should be already at the WR stage. The main argument was the width of the observed emission line profile (FWHM $\\simeq$\\,974\\kms), typical of WR stars, making the source belonging to a class of broad-line stars including 8 other stars in the central cluster of helium stars.  It was also measured that the broad-line stars were in average weaker by $\\simeq$\\,2.4 mag in the K band than the rest of helium sources which were characterized by a narrow \\ion{He}{i} emission line. With a K$_{mag}$ for IRS~13E3 consistent with the magnitude measured for the other broad-line stars, and 13E1 and 13E2 much brighter, it was proposed that IRS~13E3 should be the WR-type helium emitter. of because and It was not possible to precise the 13E1 and 13E2 stellar type.  In the centimetric domain, Zhao \\& Goss (\\cite{zhao}) presented the detection of IRS~13 at 7 and 13~mm with the VLA, at a resolution of 0.06\\arcsec, which they reported to be the brightest radio continuum source after \\SgrA\\/ at the Galactic Center.  They resolved the source into two components, one with no significant proper motion while the other one is moving south at a rate 6.2\\,$\\pm$\\,1.1\\,mas~yr$^{-1}$. They called the two compact \\ion{H}{ii} regions IRS~13E and IRS~13W, which was improper, since these denominations had already been given to infrared sources at IRS~13 as reminded above, with which they are not coincident. However, this detection is another element which makes this source special.  The detection of a bright, discrete X-ray emission source (Baganoff \\etal \\cite{baganoff}, Muno \\etal \\cite{muno}) at the IRS~13 position within the positional accuracy -- source \\object{CXOGC J174539.7-290029} -- among the brightest sources within the central parsec besides \\SgrA, is a last element contributing to make this source an object of interest.  This coincidence already triggered the interest of Coker \\etal (\\cite{coker}) who presented the first Chandra X-ray spectrum of IRS~13 and deduced that it was consistent with a highly absorbed X-ray binary system. They concluded that IRS~13E2 should be a compact post-LBV binary whose colliding winds were the source of the X-ray emission.  As a conclusion, the origin of the brightness of IRS~13, one of the brightest objects at all wavelengths in the vicinity of \\SgrA, has been a matter of debate.  We examined all the high-angular resolution images in the near infrared currently available.  This analysis made it possible to build a completely new picture of IRS~13 which is presented in this paper. We will show that the peculiar spectral energy distribution (SED) previously reported is well explained by the nature of the individual sources which compose IRS~13. Unexpectedly, this analysis provides new insight to the unsolved question of the formation of the central massive cluster. It has also given the opportunity through the study of the stars in the IRS~13 field to characterize a sample of the star population in the central parsec.  A preliminary version of the paper was presented at the Galactic Center Workshop 2002 (Maillard \\etal \\cite{maillard2}). ",
        "conclusions": "The presence of a compact star cluster of six hot, massive stars at the position of IRS~13E from high-resolution near-infrared observations is demonstrated. The spectral types of the various members range from O to WR, including dusty WRs. Proper motion measurements indicate that the brightest stars are co-moving suggesting that the members of the cluster are bound by a central IMBH with a mass $\\geq$\\,1300\\,M$_{\\odot}$.  Such a secondary black-hole in the vicinity of \\SgrA\\/ could be the element needed to explain the population of massive young stars observed today (Hansen \\& Milosavljevi\\'c \\cite{hansen}).  To precise the spectral type of the components, better constrain the mass of the IMBH, spectroscopy in the 1-5\\micron\\/ range of all the individual sources within IRS~13E, at angular resolution as good as 0.1\\arcsec, is required.  In the L and M bands it should confirm the expected featureless spectrum, except dust signatures, of the IRS~13E3 and 13E5 objects. These studies will need near-infrared 3-D spectrometers behind an AO system on a 8-m telescope, like SINFONI (Mengel \\etal \\cite{mengel} and AMBER behind VLTI (Petrov \\etal \\cite{petrov}). Deeper AO imaging, as already obtained with NAOS/CONICA north of the IRS~13E center (Eckart \\etal \\cite{eckart2}), could make it possible to detect more members of the cluster.  Proper motions of the fainter members would help to confirm which of the individual sources are kinematically bound together. Theoretical work is needed to confirm whether IRS~13E can be the remnant of a massive cluster. More generally, to address the problem of recent star formation in the vicinity of \\SgrA\\/ a full census of the hot star population, illustrated in this paper on a small field, remains to be completed. It can be done by deep AO imaging to the condition to acquire data down to 1\\micron. Up to now, such data, which appeared essential as shown in the current analysis, have only been possible with NICMOS on HST. Search for optical counterparts of the star-like X-ray sources detected by Chandra in the central parsec is another objective."
    },
    "0404/astro-ph0404385_arXiv.txt": {
        "abstract": "{We present an X-ray spectral analysis of 21 low redshift quasars observed with {\\sl XMM-Newton} EPIC. All the sources are Palomar Green quasars with redshifts between 0.05 and 0.4 and have low Galactic absorption along the line-of-sight.  A large majority of quasars in the sample (19/21) exhibit a significant soft excess below $\\sim$  1--1.5\\,keV, whilst two objects (PG\\,1114+445 and I\\,Zw1) show a deficit of soft X-ray flux due to the presence of a strong warm absorber. Indeed, contrary to previous studies with {\\sl ASCA} and {\\sl ROSAT}, we find that the presence of absorption features near  0.6--1.0\\,keV is common in our sample. At least half of the objects appear to harbor a warm absorber, as found previously in Seyfert 1 galaxies. We find significant detections of Fe\\,K$\\alpha$ emission lines in at least twelve objects, whilst there is evidence for some broadening of the line profile, compared to the EPIC-pn resolution, in five of these quasars. The determination of the nature of this broadening (e.g., Keplerian motion, a blend of lines, relativistic effects) is not possible with the present data and requires either higher S/N or higher resolution spectra. In seven objects the line is located between 6.7--7\\,keV, corresponding to highly ionized iron, whereas in the other five objects the line energy is consistent with 6.4 keV, i.e. corresponding to near neutral iron. The ionized lines tend to be found in the quasars with the steepest X-ray spectra. We also find a correlation between the continuum power law index $\\Gamma$ and the optical H$\\beta$ width, in both the soft and hard X-ray bands, whereby the steepest X-ray spectra are found in objects with narrow H$\\beta$ widths, which confirms previous {\\sl ROSAT} and  {\\sl ASCA} results. The soft and hard band X-ray photon indices are also strongly correlated, i.e. the steepest soft X-ray spectra correspond the steepest hard X-ray spectra. We propose that a high accretion rate and a smaller black hole mass is likely to be the physical driver responsible for these trends, with the steep spectrum objects likely to have smaller black hole masses accreting near the Eddington rate. ",
        "introduction": "In Active Galactic Nuclei (AGN), from Seyfert galaxies to quasars, the analysis of spectral X-ray features can help us to understand the central regions of these powerful objects. The first is the so-called soft excess seen below 2--3\\,keV (Arnaud et al. \\cite{Arnaud85}; Turner \\& Pounds \\cite{Turner89}). This spectral characteristic is thought to be the high energy part of the optical-UV ``big blue bump''  extending down to 1$\\mu$m, which contains a large fraction of the bolometric luminosity. Soft X-ray excesses were detected for most AGN with {\\sl ROSAT} (e.g., Brinkmann \\cite{Br92}, Buehler et al. \\cite{Bu95}). Using {\\sl XMM-Newton} observations, Pounds \\& Reeves (\\cite{Pounds2002}), recently confirmed  the presence of a soft X-ray excess in all objects of their sample (6 Seyfert galaxies), with  amplitude and width increasing with luminosity. Current interpretations of the soft excess range from direct thermal emission from the accretion disk to reprocessing of harder radiation absorbed in the thin disk (Pounds \\& Reeves \\cite{Pounds2002}). Other important components are emission and/or absorption structures, mainly in the soft X-ray range, due to a warm absorbing medium supposed to be located between the  Broad Line Region and the Narrow Line Region (e.g., Reynolds \\& Fabian \\cite{RF95}; Porquet et al. \\cite{Porquet99}). These absorption and/or emission features are commonly seen in low-luminosity AGN such as Seyfert galaxies, but are rather rare in quasars. The Fe\\,K$_{\\alpha}$ line complex observed in the 6--7\\,keV range is also an important spectral diagnostic tool to probe dense matter from the inner disk  (e.g., \\object{MCG-06-30-15}, Tanaka et al. \\cite{T95}, Fabian et al. \\cite{F2002}; \\object{Q0056-363}, Porquet \\& Reeves \\cite{PR2003}) to the Broad Line Region and the molecular torus (e.g.  \\object{NGC 5548}, Yaqoob \\cite{Ya2001}; \\object{Mrk 205}, Reeves et al. \\cite{R2001}; \\object{Mrk 509}, Pounds et al. \\cite{Pounds2001}). Recently, Page et al. (\\cite{Pa2004}) have shown that the majority of 53 AGN observed with {\\sl XMM-Newton}, including both Seyfert 1 and quasars, exhibit narrow Fe\\,K line near 6.4\\,keV, due to the emission from neutral material, such as the BLR or the putative molecular torus. \\\\ \\indent In this paper,  we present the spectral analysis of a sample of 21 low-redshift  quasars (z$\\sim$0.05--0.4), obtained with {\\it XMM-Newton}. One great advantage of such study with the EPIC cameras (MOS and pn) is the combination of high sensitivity with moderate spectral resolution over a broad energy band, i.e. from 0.3\\,keV up to 12\\,keV for the pn camera. All sources are Palomar Green objects (Green et al. \\cite{Green86}) with low absorption column density  along the line-of-sight. Note that all fit parameters given in this paper are in the quasars rest frame, with values of H$_{\\rm 0}$=75\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, and q$_{\\rm 0}$=0.5 assumed throughout. \\\\ \\indent In section~\\ref{sec:xmm}, we present the quasar sample used in this work as well as the data reduction techniques. In section~\\ref{sec:mdot}, we estimate the accretion rate with respect to the Eddington rate for each quasar. In section~\\ref{sec:xmmspectra}, we analyze the spectral shapes over different energy bands. The possible presence of a warm  absorber medium and of a Fe\\,K line is checked in sections~\\ref{sec:WA} and \\ref{sec:FeK}, respectively. Then, in sect.~\\ref{sec:all}, we investigate possible potential origins for the soft excess and which model could account for the whole spectral shape. Finally, in sect.~\\ref{sec:corr}, we report the global properties of the quasar sample and check for correlations between the X-ray properties and other quasar properties. In appendix~A, we report individual notes about each object of our sample. ",
        "conclusions": "We have presented results from {\\sl XMM-Newton} EPIC observations of 21 low redshift quasars with redshifts less than 0.4.  All  objects are Palomar Green quasars and have low Galactic absorption along the line-of-sight. Five objects are Narrow Line Seyfert1/quasars, two are RLQ, and the other 14 objects are RQQ. We now summarize our main results: \\\\  About 90$\\%$ of the present quasars (19/21) exhibit a significant soft excess at energies below about 1--1.5\\,keV. This result is in contrast to previous findings with {\\sl ASCA}, e.g. George et al. (2000) found that only five of their 14 PG quasars (with z$\\leq$0.25) require a soft excess below 1\\,keV. This is likely due to the smaller band-pass of {\\sl ASCA} compared to {\\sl XMM-Newton} at lower energies. For objects with a strong soft excess and a large Fe\\,K line, the soft excess cannot be solely due to reprocessing of the primary X-ray component by an ionized disk. In all objects in the present sample (except possibly PG\\,1244+026), the soft X-ray excess is much better explained by Comptonization of soft photons in a hot plasma, than by the thermal emission from a thin standard accretion disk, as the inferred disc temperatures appear to be too hot. However, to definitely discriminate between these two models, a larger bandpass is required, from the EUV (for the soft component) to hard X-rays up to few 100\\,keV (for the hard component).\\\\  PG\\,1114+445 shows a huge absorption trough between 0.6 and 1.5\\,keV, due to the presence of a warm absorber. Fitting the spectrum with two absorption edges, we find very large optical depths of $\\tau\\sim$2.23 and  $\\tau\\sim$0.53 at E=0.72\\,keV and E=0.93\\,keV, respectively. Further nine other objects show the presence of one or two absorption edges near those of \\ion{O}{vii} (0.74\\,keV) and \\ion{O}{viii} (0.87\\,keV). This means that at least 50$\\%$ of our present quasars show absorption, which is compatible with the fraction found in Seyfert\\,1 galaxies at lower redshift (e.g., Reynolds \\cite{Re97} with {\\sl ASCA}). This is the first time that such a high fraction of warm absorbers is found in quasars and may imply a similar covering fraction of the absorbing material compared to the Seyfert 1s. The lack of detection of warm absorbers in quasars with {\\sl ROSAT} and {\\sl ASCA} may have been due to the limited bandpass for both instruments, i.e. the lack of soft response of {\\sl ASCA} and the limited spectral resolution of the {\\sl ROSAT} PSPC. The present data do not enable us to determine the exact nature (ionization degree, velocity, etc.)  and the complexity (e.g., absorption/emission lines, absorption edges) of the absorbing medium. Higher spectral resolution data are required (e.g. RGS), which will be investigated in a subsequent paper. \\\\  We find evidence for Fe K shell emission lines in at least twelve objects. Indeed, this is probably a lower limit on the number of objects with iron lines, since the line detection depends on the exposure times which are very short in some objects. Nonetheless, iron lines are detected in five objects at $>$99$\\%$ confidence (I\\,Zw1, PG\\,0804+761, PG\\,1114+445, PG\\,1116+215 and PG\\,1402+261), whilst seven objects have a more marginal detection at 90--99$\\%$ confidence (PG\\,0947+396, PG\\,1048+342, PG\\,1115+407, PG\\,1244+026, PG\\,1309+355, Mkn\\,1383, and PG\\,1512+370). In five objects the line energy is consistent within the errors with 6.4 keV, a ``neutral'' to moderately ionized iron (i.e. $\\leq$ \\ion{Fe}{XVII}), while for the other seven objects the line energy is consistent with  highly ionized iron:  I\\,Zw1, PG\\,0804+761, PG\\,1115+407, PG\\,1116+215, PG\\,1244+026, PG\\,1402+261, and Mrk\\,1383. For  five objects the fit is statistically improved when leaving the width of the line as a free parameter. All these line profiles can be well fitted with either a broad Gaussian line, or with a relativistic line profile (non-rotating Schwarzschild BH or rapidly rotating Kerr BH). No significant differences have been found comparing the $\\chi^{2}$ obtained for the relativistic line profile and for the broad Gaussian line. The width of the Gaussian line could be due to a blend of several (neutral/ionized) lines or simply to Keplerian motion. With the present S/N data we are not able to infer the incidence of genuine relativistic effects in quasars. Higher S/N data or larger spectral resolution are needed for such investigation. The line profiles can be also be interpreted as ionized disk reflection, which is able to explain the spectral shape over the 2--12\\,keV energy range. At least five significant detections (at $\\geq$ 99$\\%$ confidence) of broad iron lines in quasars have been found with {\\sl XMM-Newton} to date: I\\,Zw1, PG\\,0804+761, PG\\,1116+215, and PG\\,1402+261 (this work), and Q\\,0056+363 (Porquet et al. \\cite{PR2003}). The first three objects exhibit line energies corresponding to very ionized iron (He-like, H-like). As shown previously by Porquet \\& Reeves (\\cite{PR2003}), broad and intense lines at 6.4\\,keV in quasars appear to be very rare, with the quasar Q0056-363 being the most luminous quasar found to date exhibiting such characteristics.  The RQQ PG\\,1402+261 shows a very large positive deviation near 7--8\\,keV compared to a power law model. Assuming that this feature is a relativistic Fe\\,K line, we find a huge EW of about 2\\,keV, thus the hard X-ray spectrum of this object is likely to be dominated by reflected emission from the disc. Other interpretations such as partial covering absorber models will be investigated in detail in a \\ forthcoming paper (Reeves et al. 2004, in preparation). Four of the five NLG of the sample show the presence of a very highly ionized  Fe\\,K line: I\\,Zw1, PG\\,1115+407, PG\\,1244+026, and PG\\,1402+261. Our results are also consistent with the correlations found between the Fe\\,K line energy and the 2--10\\,keV  X-ray power law slope by Dewangan (\\cite{D2002}). This means that the steep X-ray spectrum objects (such as NLG) tend to have Fe\\,K lines formed in a highly ionized medium, while objects with flatter X-ray spectra tend to be associated with near neutral or weakly ionized iron line emission. \\\\  A strong correlation is found between $\\Gamma$ (both in the soft 0.3--2\\,keV and hard 2--10\\,keV energy bands) and optical H$\\beta$ width, whereby the steepest X-ray spectra tend to be found in those objects with narrow H$\\beta$ widths. This is consistent with previous results (e.g., Boller et al. \\cite{Bo96}, Wang et al. \\cite{WBB96}, Brandt et al. \\cite{B97}, Vaughan et al. \\cite{V99}, Reeves \\& Turner \\cite{RT2000}, and Dewangan \\cite{D2002}). The soft and hard X-ray photon indices are also linked by a very strong correlation, i.e. the steepest soft X-ray spectra lead to the steepest hard X-ray spectra. The strongest correlations are found between $\\Gamma$ (soft and hard), H$\\beta$ width  and the black hole mass and the accretion rate. Therefore, we conclude that a high accretion rate and a smaller black hole mass is likely to be the physical driver responsible for these trends, i.e. the steepest X-ray spectra are often found in objects accreting at high fraction of the Eddington rate, with smaller black hole masses. \\\\"
    },
    "0404/astro-ph0404516_arXiv.txt": {
        "abstract": "We have used the Green Bank Telescope (GBT) and Berkeley-Illinois-Maryland Association (BIMA) array to search for redshifted millimetre absorption in a sample of damped Lyman-alpha absorption systems (DLAs). This brings the number of published systems searched from 18 to 30. In 17 cases we reach $3\\sigma$ limits of $\\tau\\leq0.1$, which is a significant improvement over the previous searches and more than sufficient to detect the 4 known redshifted millimetre absorbers ($\\tau\\gapp1$). While the CO rotational (millimetre) column density limits obtained are weaker than the electronic (optical) limits, they may provide useful limits below the atmospheric cut-off for the Lyman and Werner \\MOLH-bands in the UV ($z_{\\rm abs}\\lapp1.8$). Using a model for the DLA metallicity evolution combined with assumed HCO$^+$/\\MOLH ~and CO/\\MOLH ~conversion ratios, we use the molecular column density limits to calculate plausible \\MOLH ~molecular fraction limits. Finally, we use these results to discuss the feasibility of detecting rotational CO transitions in DLAs with the next generation of large radio telescopes. ",
        "introduction": "\\label{sec:intro}  Molecular absorption lines trace, and provide detailed physical and chemical information about, the cold dense component of the interstellar medium (ISM).  Although many detailed studies exist for molecular clouds within our own Galaxy, only relatively recently has detailed information emerged for molecular abundances at high redshift, through absorption studies of redshifted UV molecular hydrogen lines (e.g. \\citealt{lps03,rbql03}) and millimetre-band rotational lines from molecular tracers (e.g. \\citealt{wc96a}).   Observations of a range of different molecular transitions in gas clouds at high redshift would provide a wealth of information on star formation activity in external galaxies, potentially viewed at epochs when chemical abundances and environments were markedly different to today.  Such information is invaluable for a detailed understanding of galactic formation and evolution. Furthermore, the narrowness of molecular lines reveals information about small-scale structure in the ISM. Comparison of relative line strengths yields information about the excitation mechanism, and in particular, probes the temperature of the Cosmic Microwave Background (CMB), and hence the expected ($1+z$)--dependence (e.g. \\citealt{wc96a}). Finally, comparisons of the relative observed frequencies of millimetre molecular lines (with each other, and/or with atomic transitions arising in the same cloud) can be used to check on any possible variation in certain combinations of the fundamental constants \\citep{dwbf98,mwf+00}. This last point was the prime motivation for our search for new redshifted millimetre absorbers reported in this paper, although we use the upper limits obtained to yield molecular fraction limits in low redshift DLAs, which is not possible using the UV \\MOLH ~lines, which fall below the atmospheric cut-off of $z_{\\rm abs}\\lapp1.8$.  Currently, only 4 redshifted millimetre absorption systems are known (see \\citealt{wc99b} and references therein), of which the highest redshift is 0.886 (PKS 1830--211). As a means of approaching a search for new high redshift radio absorbers systematically, we produced a catalogue of DLAs \\citep{cwbc01}\\footnote{A version of this catalogue is kept updated on-line and is available from http://www.phys.unsw.edu.au/$\\sim$sjc/dla}, where large column densities ($N_{\\rm HI}\\geq2\\times10^{20}$ cm$^{-2}$) are known to exist, and shortlisted those which are illuminated by radio-loud quasars (i.e. those with a measured radio flux density $>0.1$ Jy). This yielded 60 DLAs and sub-DLAs occulting radio-loud quasars. Of these, 37 have been searched for 21-cm absorption (see \\citealt{kc02,cmp+03}). Selecting those of 12-mm and 3-mm flux densities $\\gapp0.1$ Jy, gives 18 systems which have previously been searched for millimetre absorption (\\citealt{cmwp02}), this number now being increased in total to 30 with the observations we present in this paper. ",
        "conclusions": "\\subsection{Search results}   In Table \\ref{sum} we show the best previously published optical depths together with our new results, for which, despite the improved limits, there are no absorption features of $\\geq3\\sigma$ over $\\geq3$ \\kms, the full resolution of the BIMA array observations. For all of the limits we estimate the $3\\sigma$ upper limits on the total column density of each molecule from \\begin{equation} N_{\\rm mm}=\\frac{8\\pi}{c^3}\\frac{\\nu^{3}}{g_{J+1}A_{J+1\\rightarrow J}}\\frac{Qe^{E_J/kT_x}}{1-e^{-h\\nu/kT_x}} \\left.\\int\\right.\\tau dv, \\end{equation} where $\\nu$ is the rest frequency of the $J\\rightarrow J+1$ transition, $g_{J+1}$ and $A_{J+1\\rightarrow J}$ are the statistical weight and the Einstein A-coefficient\\footnote{These are taken from \\citet{cklh95,cms96} or derived from the dipole moment (e.g. \\citealt{rw00}).} of the transition, respectively, $Q = \\sum^{\\infty}_{J=0}g_{J}~e^{-E_J/kT_x}$ is the partition function\\footnote{The energy of each level, $E_J$, is obtained from the JPL Spectral Line Catalog \\citep{ppc+98}.}, for the excitation temperature, $T_x$, and $\\int\\tau dv$ is the $3\\sigma$ upper limit of the velocity integrated optical depth of the line\\footnote{$\\int\\tau dv\\approx1.06\\,\\tau_{\\rm mm}\\times{\\rm FWHM}$ for a Gaussian profile, where $\\tau_{\\rm mm}$ is the $3\\sigma$ peak optical depth limit (Table \\ref{sum}).}.  Since the derived optical depth limits depend upon the r.m.s. noise and thus the velocity resolution (see \\citealt{cwn+02}), in our previous articles we presented all of the optical depth limits normalised to a spectral resolution of 1 \\kms, the finest resolution typical of most current wide-band spectrometers.  Previously, we also quoted column density upper limits per unit \\kms~ line-width. However, in order to use our limits to constrain molecular fractions at low redshift (Section 3.2), we shall now adopt a line-width (FWHM) of 10 \\kms, which is close to those of the 4 known systems (e.g. \\citealt{wc96a}). This is done for a resolution of the same value and so the column density limits, which are calculated for an excitation temperature\\footnote{Due to the increase of CMB temperature with redshift, $T_{\\rm CMB}=2.73 (1+z_{\\rm abs})$, $T_{x}$ increases to 20 K at $z_{\\rm abs}=3.75$, the highest redshift of our sample. The main effect of this is to increase the HCO$^+$ $0\\rightarrow1$ column density estimates by a factor of $\\leq4$ in comparison to a constant value of 10 K.} of 10 K at $z_{\\rm abs}=0$, represent a one channel $3\\sigma$ detection of a $10$ \\kms ~wide line. \\begin{table*} \\centering \\begin{minipage}{162mm} \\caption{Summary of published searches for molecular tracer absorption in DLAs and sub-DLAs. $\\nu_{\\rm obs}$ is the approximate observed frequency (GHz), $V$ is the visual magnitude of the background quasar and $S$ is the flux density (Jy) at $\\nu_{\\rm obs}$ which, unless flagged, is the value obtained during the actual observations (the limits are $1\\sigma$ and blanks in this field indicate that no value is given in the literature). $N_{\\rm HI}$ (\\scm) is the DLA column density from the Lyman-alpha line and $\\tau_{\\rm 21~cm}$ is the normalised peak optical depth of the redshifted 21-cm line (see \\protect\\citealt{cmp+03} for details). The optical depth of the relevant millimetre line is calculated from $\\tau=-\\ln(1-3\\sigma_{{\\rm rms}}/S)$, where $\\sigma_{{\\rm rms}}$ is the r.m.s. noise level. Since $\\sigma_{{\\rm rms}}$ is dependent on the spectral resolution, we take the various published values and recalculate $\\sigma_{{\\rm rms}}$ at a resolution of 10 \\kms ~($\\tau_{\\rm mm}$) and quote only the best existing limit. For all optical depths, $3\\sigma$ upper limits are quoted and ``--'' designates where $3\\sigma>S_{{\\rm cont}}$, thus not giving a meaningful value for this limit. The penultimate column gives the corresponding column density (\\scm) per 10 \\kms ~channel (see main text).}% \\begin{tabular}{@{}l c c r c c c c c c c @{}} \\hline DLA & $z_{\\rm abs}$ & Transition & $\\nu_{\\rm obs}$ & $V$ & $S$ & $N_{\\rm HI}$ & $\\tau_{\\rm 21~cm}$ &$\\tau_{\\rm mm}$ & $N_{\\rm mm}$ & Ref.\\\\ \\hline 0201+113 & 3.38639 & HCO$^+$ $0\\rightarrow1$ & 20.3 & 19.5 &0.47 & $2\\times10^{21}$ &  $\\leq0.09$ &$<0.007$ & $<1\\times10^{12}$ & 10\\\\ 0201+365 & 2.4614 & HCO$^+$ $0\\rightarrow1$ & 25.8 & 17.9 &0.09$^V$  & $3\\times10^{20}$  & & $<0.1$ & $<1\\times10^{13}$ & 10\\\\ 0235+164 & 0.52400 & CO $0\\rightarrow1$& 75.6 &15.5 & 2.5& $4\\times10^{21}$ & $0.64$ &$<0.02$&$<1\\times10^{15}$ & 2\\\\ ... &0.52398 & CO $1\\rightarrow2$ & 151.3& ... & 1.27& ... & ... & $<0.03$& $<1\\times10^{15}$& 6\\\\ ...\t&  ...& HCO$^+$ $3\\rightarrow4$ & 234.1 & ... & 0.75&... &...  & $<0.1$  &$<7\\times10^{12}$ & 6\\\\ ...& 0.523869& CS $2\\rightarrow3$ & 96.4 &... & $1.7$ &...&  ...&$<0.1$  &$<4\\times10^{13}$ & 8\\\\ 0248+430 & 0.3939& CS $2\\rightarrow3$ & 105.4 &17.7 &$<0.2$ &$4\\times10^{21}$ &  0.2& $<0.4$&$<1\\times10^{14}$ & 8\\\\ ...& 0.394 &  CO $0\\rightarrow1$& 82.7 & ... & 0.21    &...&  ...&  $<0.4$ &$<3\\times10^{16}$  & 10\\\\ 0335--122 & 3.178 &  HCO$^+$ $0\\rightarrow1$ & 21.4 & 20.2 & 0.13$^{A,V}$& $6\\times10^{20}$  &$<0.008$ &$<0.03$&  $<6\\times10^{12}$& 10\\\\ 0336--017 & 3.0619 & HCO$^+$ $0\\rightarrow1$ &22.0 &18.8 & 0.15$^A$& $2\\times10^{21}$ &$<0.007$ &$<0.03$& $<5\\times10^{12}$& 10\\\\ 0405--331 &2.570 & HCO$^+$ $0\\rightarrow1$ & 25.0 & 19.0 & 0.53&$4\\times10^{20}$ & &$<0.02$&$<3\\times10^{12}$ & 10\\\\ 0458--020 & 2.0397  & HCO$^+$  $2\\rightarrow3$  & 88.0 &18.4 & & $5\\times10^{21}$ &0.3  & $<0.3$& $<1\\times10^{13}$& 7\\\\ ...\t&   2.0399 &... & 88.0 & ...&$1.3$& ... &... & $<0.07$ &$<3\\times10^{12}$ & 8\\\\ ...\t& 2.03937 &CO $0\\rightarrow1$ & 37.9 &  ...& $\\approx0.8$ & ...&  ... & $<0.01$ &$<1\\times10^{15}$ & 4\\\\ ...\t&  2.0398 & CO $2\\rightarrow3$ & 113.8 &... &0.53 & ...&  ... &  $<0.1$& $<6\\times10^{15}$ & 10\\\\ ...\t&  2.0397 & CO $2\\rightarrow3$ & 113.8 &... &0.53$^a$ & ...&  ... &$<0.6$&  $<4\\times10^{16}$& 5\\\\ ...\t&   2.0399 & CO $3\\rightarrow4$& 151.7 &... & $0.4$& ...& ... &$<0.4$ & $<4\\times10^{16}$ & 8\\\\ ...& 2.04 & H$_2$CO $1_{10}\\rightarrow1_{11}$ & 1.6 & ...& & ...&...&$<0.01$ &  & 3\\\\ 0528--2505 & 2.1408 & CO $2\\rightarrow3$ &110.1 &19.0 & $\\approx0.2^b$ &$4\\times10^{20}$ & $<0.3$ & $<1.0$& $<6\\times10^{16}$& 5\\\\ ... & ... & HCO$^+$ $0\\rightarrow1$ &23.4 &  ...& 0.3$^{A,V}$ & ... & ... & $<0.02$&$<3\\times10^{12}$ & 10\\\\ 0537--286 & 2.974 & HCO$^+$ $0\\rightarrow1$ & 22.4 &19.0 & 0.58& $2\\times10^{20}$ &$<0.007$ &$<0.02$ &$<3\\times10^{12}$ & 9\\\\ 0738+313 &  0.2212 & CO $0\\rightarrow1$ & 94.4 &16.1 &$0.48$ & $2\\times10^{21}$ & $\\approx0.07$& $<0.04$& $<3\\times10^{15}$  & 10\\\\ 0827+243 & 0.5247 & CS $2\\rightarrow3$ & 96.4 &17.3  &$2.7$ & $2\\times10^{20}$ &0.007  &$<0.09$ & $<3\\times10^{13}$& 8\\\\ 08279+5255  & 2.97364& HCO$^+$ $0\\rightarrow1$ & 44.9 &15.2 & &$1\\times10^{20}$ &  & \\multicolumn{2}{c}{{\\it No 7 mm flux available}}   & 8\\\\ ...\t&  ... & CO $2\\rightarrow3$ & 87.0 &... & $<0.1$ &... & ... & & & 8\\\\ 0834--201 & 1.715 & HCO$^+$  $2\\rightarrow3$ &98.6 &  18.5& $1.7$ & $3\\times10^{20}$ & & $<0.1$&$<4\\times10^{12}$  & 8\\\\ ...\t&  ... &HCO$^+$  $3\\rightarrow4$& 131.4 &... & & ...&  ...& $<0.2$&$<1\\times10^{13}$ & 7\\\\ ...\t&  ... &  CO $3\\rightarrow4$& 169.8 &... & $0.9$& ...&  ...&$<0.7$ &$<8\\times10^{16}$  & 8\\\\ 0913+003 & 2.774 & HCO$^+$ $0\\rightarrow1$ & 23.8 & -- & $\\approx0.17$$^{A,V}$ & $2\\times10^{20}$ &  &$<0.04$ & $<6\\times10^{12}$ &10\\\\ 1017+1055 & 2.380& CS $2\\rightarrow3$ & 43.5 &17.2 & & $8\\times10^{19}$ &  & \\multicolumn{2}{c}{{\\it No 7 mm flux available}}  & 8\\\\ ...\t&  ... & CO $2\\rightarrow3$ & 102.3 &... &$<0.2$ &  ...&  ... & --& --&   8\\\\ 1215+333 & 1.9984 & CO $2\\rightarrow3$ & 115.3 &18.1 & & $1\\times10^{21}$ &   &   \\multicolumn{2}{c}{{\\it No 3 mm flux available}}  & 5\\\\ 1229--021 & 0.3950 & CO $0\\rightarrow1$  & 82.6 &16.8  &$0.2$ & $1\\times10^{21}$ &0.05  & $<0.7$& $<5\\times10^{16}$& 8\\\\ ...\t&  ... & CO $1\\rightarrow2$& 165.3 &... &$<0.1$ &  ...&  ...& --& --& 8\\\\ ...\t& 0.39498  & CO $1\\rightarrow2$& 165.3 &... &$0.11$ &...&  ...&-- &-- &6\\\\ 1251--407 & 3.752 &   HCO$^+$ $0\\rightarrow1$ & 18.8 &23.7 & $\\approx0.1$$^{A,V}$ & $2\\times10^{20}$ & &$<0.09$& $<2\\times10^{13}$& 10\\\\ 1328+307 & 0.69215 & HCO$^+$  $1\\rightarrow2$ & 105.4 &17.3  &0.50 & $2\\times10^{21}$ &0.02  & $<0.2$& $<6\\times10^{12}$& 6\\\\ ...\t&  ... & CS $2\\rightarrow3$ & 86.9 &... & $1.0$ &... & .... & $<0.3$& $<8\\times10^{13}$ &  8\\\\ ...\t&  ... & CO $1\\rightarrow2$ & 136.2 &... & 0.39& ...& ... &$<0.3$ &$<1\\times10^{16}$ & 6\\\\ ...\t&  ... & CO $2\\rightarrow3$ & 204.4 &... & 0.27& ...& ... &$<0.5$ & $<3\\times10^{16}$& 6\\\\ 1331+170 & 1.7764 & CO $0\\rightarrow1$ & 41.5 &16.7 & 0.5& $3\\times10^{21}$ & 0.02 &$<0.4$ & $<5\\times10^{16}$ & 1\\\\ ...& 1.7755 & ...&41.5 &  ... & ...& ...& ... &$<0.4$ &$<5\\times10^{16}$ & 1\\\\ 1354--107 & 2.501 & HCO$^+$ $0\\rightarrow1$ & 25.5 & 19.2 &$\\approx0.07$$^{A,V}$  &  $3\\times10^{20}$ & $<0.015$ & $<0.08$&$<1\\times10^{13}$ & 10\\\\ ... &  2.996 & ...& 22.5 & ... &...  & ... & ... & $<0.1$& $<2\\times10^{13}$& 10\\\\ 1402+044 &  2.713 & HCO$^+$ $0\\rightarrow1$ & 24.0 & 19.8 & 0.21$^*$ & $8\\times10^{19}$ &  &$<0.008$& $<1\\times10^{12}$ & 10\\\\ ... &... & HCN  $0\\rightarrow1$ & 23.9 & ... & 0.6 & ... & ...& $<0.02$& $<8\\times10^{12}$ & 10\\\\ 1418--064 & 3.449 &  HCO$^+$ $0\\rightarrow1$ & 20.1 & 18.5 & 0.14 & $3\\times10^{20}$ &   &$<0.02$ & $<4\\times10^{12}$ & 10\\\\ 1451--375 &0.2761 & HCO$^+$ $1\\rightarrow2$ & 139.8 & 16.7 &$0.6$ & $1\\times10^{20}$ & $<0.007$ & $<0.2$ &$<7\\times10^{12}$  & 8\\\\ ...\t&  ... &CO $0\\rightarrow1$  &90.3 & ... & $1.2$ & ...& ... & $<0.1$&$<8\\times10^{15}$& 8\\\\ \\hline \\end{tabular} \\label{sum} {References: (1) \\citet{tsi+84}, (2) \\citet{tnb+87}, (3) \\citet{bwl+89}, (4) \\citet{tn91}, (5) \\citet{wc94}, (6) \\citet{wc95}, (7) \\citet{wc96b}, (8) \\citet{cwn+02}, (9) \\citet{cmwp02}, (10) This paper.\\\\ Flux densities: $^A$ATCA June \\& August 2002 \\citep{cmwp02,rcm+03}, $^V$VLA May 2003. Where the flux densities could not be obtained from the data/article -- $^a$the BIMA array value, $^b$interpolated between 11 GHz and K-band, $^c$interpolated between radio and 0.5--10 keV, $^d$extrapolated from 0.4 and 5 GHz. $^*$This observation of 1402+044 is from preliminary (February 2003) GBT observations using the wide band spectrometer.} \\end{minipage} \\end{table*} \\begin{table*} \\addtocounter{table}{-1} \\centering \\begin{minipage}{167mm} \\caption{{\\it Continued}} \\begin{tabular}{@{}l c c r c c c c c c c @{}} \\hline DLA & $z_{\\rm abs}$ & Transition & $\\nu_{\\rm obs}$ & $V$ & $S$ & $N_{\\rm HI}$ & $\\tau_{\\rm 21~cm}$ &$\\tau_{\\rm mm}$ & $N_{\\rm mm}$ & Ref.\\\\ \\hline 1614+051 &  2.52 & HCO$^+$ $0\\rightarrow1$ & 25.3 & 19.5 & $\\sim0.3$$^c$ & $3\\times10^{20}$ & & $<0.03$&$<4\\times10^{12}$ & 10\\\\ 2131--045 & 3.27 & HCO$^+$ $0\\rightarrow1$ & 20.9 & 20.0 &  $<0.06^A$ &$1\\times10^{20}$ & & --&--  & 10\\\\ 2128--123 & 0.4298 & CS $2\\rightarrow3$ & 102.8 & 15.5 & 0.75 & $2\\times10^{19}$& $<0.003$ & $<0.04$ &$<1\\times10^{13}$ & 7\\\\ ...\t&  ... & CS $3\\rightarrow4$ & 137.1 & ... & 0.70 &  ...&   ...& $<0.05$ &$<2\\times10^{13}$ & 7\\\\ ...\t&  ... & CO $1\\rightarrow2$ & 161.2 & ... & 0.5 & ...&   ...& $<0.2$ &$<6\\times10^{15}$ & 7\\\\ ...\t&  ... & CO $2\\rightarrow3$ & 241.9 & ... & 0.4 & ...&   ...& $<0.2$ &$<1\\times10^{16}$ & 7\\\\ 2136+141 & 2.1346 & HCO$^+$ $2\\rightarrow3$&85.4 & 18.9  & 0.59& $6\\times10^{19}$&   &$<0.1$ &$<4\\times10^{12}$ & 6\\\\ ...\t&  ... &  CO $2\\rightarrow3$ & 110.3 &... &0.27 & ...&   ...&$<0.2$  &$<9\\times10^{15}$ &7\\\\ ...\t&  ... &  CO $3\\rightarrow4$ &147.1 & ... & 0.21& ...&   ...&$<0.2$  & $<2\\times10^{16}$& 7\\\\ ...\t&  ...&  CO $5\\rightarrow6$ & 220.6 &... & 0.16 &...&  ... &$<0.6$ & $<8\\times10^{17}$ &  7\\\\ 2342+342 & 2.908 &  HCO$^+$ $0\\rightarrow1$ & 22.8 & 18.4 & $\\sim0.1$$^d$ &$2\\times10^{21}$  &$<0.03$ & $<0.04$& $<7\\times10^{12}$ & 10\\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*}  \\subsection{Deriving molecular fraction limits}  Despite our improved limits there has yet to be a detection of a rotational molecular transition in a DLA. In many cases the limits far exceed those required to detect the 4 known redshifted millimetre absorbers (see below, Figs.  \\ref{co} and \\ref{hco}). These systems have molecular fractions of $f\\equiv\\frac{2N_{{\\rm H}_{2}}}{2N_{{\\rm H}_{2}}+N_{{\\rm HI}}}\\approx0.3-1.0$ (e.g. \\citealt{cw98b}), c.f. $f\\sim10^{-7}-10^{-2}$ for the 10 DLAs in which molecular hydrogen has been detected at UV and optical wavelengths (\\citealt{rbql03} and references therein). In this section we investigate the possibility of detecting molecular tracers in systems with such low molecular fraction.  \\begin{figure} \\vspace{5.3cm} \\special{psfile=MD1376rv-fig1.ps hoffset=-12 voffset=-160 hscale=43 vscale=43 angle=0} \\caption{The normalised CO column density versus the absorption redshift. The unfilled markers designate upper limits -- squares for the millimetre searches and circles for the optical searches. The filled stars represent the 4 known millimetre absorption systems. The lookback time is for $\\Omega_{\\rm m} = 0.30, \\Omega_{\\Lambda}=0.70, H_0 = 70\\,{\\rm km\\,s}^{-1}\\,{\\rm Mpc}^{-1}$.} \\label{co} \\end{figure} In Fig. \\ref{co} we show the CO rotational (Table \\ref{sum}) and electronic \\citep{gbwb97,lsb99,psl02} column density limits along with the values for the 4 known systems. While this illustrates the sensitivity to detecting absorption in the known systems, it clearly shows that the radio searches are considerably less sensitive than the optical CO searches\\footnote{Although the deep integrations of 0235+164 and 0458--020 with the NRO 45-m telescope \\citep{tnb+87,tn91} do approach the optical limits.}. They do, however, complement the optical results in giving limits at redshifts of $z_{\\rm abs}\\lapp1.8$ (below the atmospheric cut-off for direct optical \\MOLH ~detections). In Fig. \\ref{hco} we show the HCO$^+$ column density limits. \\begin{figure} \\vspace{5.3cm} \\special{psfile=MD1376rv-fig2.ps hoffset=-12 voffset=-160 hscale=43 vscale=43 angle=0} \\caption{The normalised HCO$^+$ column density versus the absorption redshift. Again the stars represent the 4 known systems and the inverted triangles are the DLA upper limits (Table \\ref{sum}).} \\label{hco} \\end{figure} Again we see that our limits are more than sufficient to detect HCO$^+$ in the 4 known systems \\citep{wc95,wc96,wc96b,wc97}, with two cases having been searched to at least an order of magnitude better than that required to detect the known systems at low redshift.  Although we cannot match the sensitivity to molecular absorption provided by the optical results, we can nevertheless attempt to derive constraints on the molecular fraction for absorption systems at low redshift. The major obstacle in this is determining which $N_{{\\rm CO}}$--$N_{{\\rm H}_{2}}$ conversion ratio to apply: Unlike HCO$^+$ (discussed later), the ratio may vary from $N_{{\\rm H}_{2}}\\sim10^6 N_{{\\rm CO}}$ for diffuse gas \\citep{ll00a} to $N_{{\\rm H}_{2}}\\sim10^4 N_{{\\rm CO}}$ for dense, dark Galactic clouds. This latter value is applied to the 4 known millimetre absorbers to yield molecular fractions from rotational lines (see \\citealt{wc98})\\footnote{Since all of these sources absorb at low redshift, $z_{\\rm abs}\\leq0.87$, and have $V\\gapp20$ there is currently little chance of detecting the \\MOLH ~line directly.}. However, DLAs have lower metallicities than Galactic systems and are less visually obscured than the 4 known systems, thus casting doubt on whether applying this conversion ratio is justified.  If we extrapolate the metallicity ([M/H]\\footnote{Defined as the heavy element abundance with respect to hydrogen, relative to that of the solar neighbourhood: ${\\rm [M/H]} \\equiv \\log_{10}[N({\\rm M})/N({\\rm H})] - \\log_{10}[N({\\rm M})/N({\\rm H})]_\\odot$.}) evolution of the general DLA population\\footnote{Due to many of the sources in Table \\ref{sum} having low redshifts and/or high visual magnitudes, there are only metallicity measurements for half of these sources, all of which belong to the sample of \\citet{pgw+03}. Therefore, by definition these follow the same metallicity evolution of the general DLA population rather than that of the \\MOLH-bearing DLAs \\citep{cwmc03}.} according to $[{\\rm M/H}]\\approx-0.26 z_{\\rm abs}$ \\citep{kf02,pgw+03}, \\begin{figure} \\vspace{5.3cm} \\special{psfile=MD1376rv-fig3.ps hoffset=-12 voffset=-162 hscale=43 vscale=43 angle=0} \\caption{The metallicity evolution of DLAs. The filled circles represent the DLAs which exhibit H$_2$ absorption with the least-squares fit shown \\protect\\citep{cwmc03}, the small unfilled diamonds represent the metallicity measurements of 60 DLAs between $1.9 < z_{\\rm abs} < 3.4$ (i.e. the general population), again with the fit shown \\citep{pgw+03}, and the unfilled squares are the DLAs searched for millimetre absorption (Table \\ref{sum}). } \\label{metal} \\end{figure} we obtain $[{\\rm M/H}]\\approx-0.75$ at $z_{\\rm abs}\\sim0$ (Fig. \\ref{metal}). At zero redshift this is still significantly lower than solar values suggesting that, in addition to lower enrichment at high redshift, there may be appreciable dust depletion of the metals, as in dense Galactic clouds (e.g. \\citealt{ss96}), although this may be more relevant to the \\MOLH-bearing DLAs where $f\\gapp10^{-2}$ at $z_{\\rm abs}\\lapp1.8$ \\citep{cwmc03}. Moreover, \\citet{lps03} suggest that the large depletion factors in DLAs indicate a significant dust content.  Therefore, a fair compromise between the diffuse/dark cases may be to assume the conversion ratio for dark clouds but scaled by the abundance of metals from which the tracer molecules form, i.e. $N_{{\\rm H}_{2}}\\approx10^{4.75} N_{{\\rm CO}}$ at $z_{\\rm abs}\\sim0$. Since the metallicity decreases with increasing redshift, we may expect the conversion of tracer to molecular hydrogen column density to scale accordingly. We therefore apply the ratio of $N_{{\\rm H}_{2}}\\approx10^{0.26 z_{\\rm abs} + 4.75} N_{{\\rm CO}}$, based upon the metallicity evolution of the general DLA population\\footnote{Since the molecular hydrogen fraction shows a strong anti-correlation with redshift for the \\MOLH-bearing DLAs (\\citealt{cwmc03} and Fig. \\ref{frac}), we may expect less depletion and thus relatively higher metallicities. Fig. \\ref{metal} does appear to suggest, however, that the metallicity is dominated by poorer chemical enrichment at higher redshifts.}. This gives $N_{{\\rm H}_{2}}\\sim10^5 N_{{\\rm CO}}$ at $z\\sim1$ rising to $N_{{\\rm H}_{2}}\\sim10^6 N_{{\\rm CO}}$ at $z\\sim5$.  From optical CO limits, \\citet{bcf87} and \\citet{cbf88} have previously noted that the carbon to hydrogen column density ratio is a tenth of the Galactic value at $z_{\\rm abs}=1.7$ and 2.3.  In addition to CO, we can use the HCO$^+$ limits, which could be a better choice of tracer, since a near constant ratio of $N_{{\\rm HCO^+}}=2-3\\times10^{-9} N_{{\\rm H}_{2}}$ is found over various regimes in the Galaxy \\citep{ll00}. Applying the conversion thus suggests that $N_{{\\rm H}_{2}}\\sim10^9 \\rightarrow 10^{10} N_{{\\rm HCO^+}}$ at $z\\sim1\\rightarrow 5$, i.e. $N_{{\\rm H}_{2}}\\approx10^{0.26 z_{\\rm abs} + 9.34} N_{{\\rm HCO}^+}$. In Fig. \\ref{frac} we show the derived molecular fractions \\begin{figure} \\vspace{9.0cm} \\special{psfile=MD1376rv-fig4.ps hoffset=-12 voffset=-57 hscale=43 vscale=43 angle=0} \\caption{Molecular fractions measured in redshifted absorbers. Again the unfilled symbols represent the upper limits from the millimetre searches, $N_{{\\rm H}_{2}}=10^{4.75} N_{{\\rm CO}}$ and $T_x = 10$ K at $z_{\\rm abs}=0$, with the squares representing CO and the inverted triangles HCO$^+$. The circles represent the H$_2$-bearing DLAs and the stars the 4 known millimetre absorption systems ($N_{{\\rm H}_{2}}=10^{4} N_{{\\rm CO}}$ and  $T_x = 10$ K, $\\forall z_{\\rm abs}$). A log version of the plot is also shown by means of a detail of the H$_2$-bearing systems.  This clearly illustrates the evolution of $f$ described by \\citet{cwmc03}.} \\label{frac} \\end{figure} of the CO and HCO$^+$ limits together with those of the H$_2$-bearing DLAs\\footnote{Versions of Fig. \\ref{frac} for $N_{{\\rm H}_{2}}=10^6 N_{{\\rm CO}}$ and $10^4 N_{{\\rm CO}}$ can be viewed at \\bf synergy url here please}.  From this we see that the high conversion ratio and excitation above the $0\\rightarrow1$ transition make the HCO$^+$ molecule insensitive with current telescopes to $f\\lapp1$: Although in many cases our GBT observations yield significantly better limits than previously (Table \\ref{sum}), these are all at high redshift and therefore convert less favourably to molecular fractions due to the steep metallicity evolution. Thus, we obtain near constant limits to the fraction across the entire redshift range searched for HCO$^+$.  Note that the optical CO limits (Fig. \\ref{co}) occupy the space between the HCO$^+$ limits and the \\MOLH-bearing DLAs ($1.7\\lapp z_{\\rm abs}\\lapp4.4$, $-3 \\lapp \\log_{10} f \\lapp-0.5$) when the conversion is applied.  We do see however, based on our redshift dependent conversion factor, that the millimetre CO searches generally give interesting (low redshift) limits, with two of the limits within the range of the H$_2$-bearing DLAs. This suggests that, were these DLAs \\MOLH-bearing\\footnote{\\MOLH ~is detected in $\\lapp20\\%$ of recent DLA surveys \\citep{lps03}, and so we perhaps expect a $\\approx1/5$ chance of having observed an \\MOLH-bearing system.}, making the $N_{{\\rm CO}}$--$N_{{\\rm H}_{2}}$ conversion reasonable, the CO rotational transition could have been detected in these DLAs: 0235+164 ($z_{\\rm abs}=0.524$) and 0738+313 ($z_{\\rm abs}=0.221$). Unfortunately, there are no metallicity measurements available for these, although \\citep{pgw+03} use the value of ${\\rm[M/H]}=-0.22\\pm0.15$ for 0235+164 which comes from model dependent X-ray measurements \\citep{trp+03}. Several other models give lower values than this, thus suggesting that this DLA may belong to the general population (${\\rm [M/H]}\\approx-0.9$), although the value quoted by \\citet{pgw+03} could place it in either population (Fig. \\ref{metal}). The general population fit is, however, consistent with the range of $-1.4 < {\\rm [M/H]} < -0.4$ at $z_{\\rm abs}\\approx0.4$ estimated by \\citet{trp+03} and the fact that we did not detect CO. Regarding the molecular fraction, the data are sensitive to $\\log_{10} f\\sim-1.5$ for the general population metallicity and $\\log_{10} f\\sim-2.1$ for the value used by \\citet{pgw+03}. This latter fraction shifts this DLA further into the \\MOLH-bearing DLA regime.  For more diffuse clouds of low dust content our limits become weaker with $f\\geq0.4$, c.f. $\\geq0.03$ previously, being obtained for $N_{{\\rm H}_{2}}\\sim10^6 N_{{\\rm CO}}$ and $T_x = 10$ K at $z_{\\rm abs}=0$, although in the diffuse gas case the column density limits would improve slightly due to lower excitation temperatures (at best a factor of $\\approx0.5$ lower than given in Table \\ref{sum})$^{13}$. That is, the clouds being diffuse could also explain why no millimetre absorption of CO was detected.  \\subsection{Detecting millimetre lines in DLAs with the next generation of large radio telescopes}  Applying the conversion of tracer column densities described above, we can ascertain the likelihood of detecting these molecular tracers in DLAs with the next generation of large radio telescopes. For $N_{{\\rm H}_{2}}=10^{4.75} N_{{\\rm CO}}$ and $T_x = 10$ K at $z_{\\rm abs}=0$: \\begin{enumerate} \\item The Atacama Large Millimeter Array (ALMA)\\footnote{http://www.eso.org/projects/alma/}: The 3-mm band of this telescope will give CO $0\\rightarrow1$ coverage to $z\\leq0.34$. At 10 \\kms ~resolution, one hour of integration will give an r.m.s. of $\\approx0.3$ mJy. For a 0.3 Jy continuum flux (fairly typical in Table \\ref{sum}), this corresponds to a $3\\sigma$ column density limit of $N_{{\\rm CO}}\\sim10^{14}$ \\scm ~per channel. Applying the conversion at these redshifts gives $N_{{\\rm H}_{2}}\\sim10^{19}$ \\scm ~thus giving $f\\sim10^{-3}-0.1$ for DLAs. If we extrapolate the molecular fractions in  \\MOLH-bearing systems back to $z\\lapp0.3$ (Fig. \\ref{frac}), we see that such limits could well be sufficient to detect rotational CO in DLAs at low redshift. \\item The Square Kilometre Array (SKA)\\footnote{http://www.skatelescope.org/}: With a tuning range of 0.15 to 20 GHz this telescope can observe CO at $z\\geq4.8$. With the same parameters as above, an r.m.s. of 5 $\\mu$Jy is reached giving $N_{{\\rm CO}}\\sim10^{12}$ \\scm ~per channel\\footnote{Again, this is for a flux density of 0.3 Jy. Due to the antenna temperature dominating the system at these low frequencies, an increased flux does not improve the optical depth to the same extent as for the millimetre observations. Therefore, the column density estimate is fairly robust.}, thus rivalling current optical limits.  At $z\\sim5$, $N_{{\\rm H}_{2}}\\sim10^6 N_{{\\rm CO}}$ giving $f\\sim10^{-4} - 10^{-2}$ for DLAs, i.e. the molecular fraction range of 7 of the 10 \\MOLH-bearing DLAs (e.g. \\citealt{rbql03}). Unfortunately at $z\\gapp5$, extrapolating the molecular fraction--redshift anti-correlation, we expect $f\\lapp10^{-12}$, thus making it unlikely that redshifted millimetre transitions could be detected\\footnote{Note, however, that from Galactic studies, \\citet{lis02} suggests that at very low metallicities \\MOLH ~will still via form slow gas-phase processes, giving minimum molecular fractions of $f\\sim10^{-8}-10^{-7}$. This suggests that current DLA molecular fraction measurements may be close to the minimum possible value (see \\citealt{lps03,rbql03}).} with the SKA, despite its superior sensitivity\\footnote{Although \\citet{cdk04} discuss the possibility of detecting low redshift OH 18-cm lines in DLAs with this telescope.}.  \\end{enumerate}"
    },
    "0404/astro-ph0404046_arXiv.txt": {
        "abstract": "3C 111  is an X-ray bright broad-line radio galaxy which is classified as a Fanaroff-Riley type II  source with a double-lobe/single jet morphology, and reported superluminal motion. It is a well-known X-ray source, and was observed by every major X-ray observatory since \\textit{HEAO-1}. In this paper we present the results of the \\textit{RXTE} and \\textit{INTEGRAL} data analysis and compare them with the results of the previous observations. ",
        "introduction": "3C~111 ($z=0.0485$) is an X-ray bright broad-line radio galaxy (BLRG) which is classified as a Fanaroff-Riley type II  source with a double-lobe/single jet morphology (Linfield \\& Perley 1984, see Figure 1) and reported superluminal motion (Preuss et al. 1988).  It is  a well-known X-ray source, and was observed by every major X-ray observatory since \\textit{HEAO-1}.  \\begin{figure} \\psfig{figure=3C111_jet.ps,width=\\columnwidth } \\caption{ Intensity map of 3C111 at 1.4 GHz. The bright central component is coincident with the nucleus of the  galaxy in the optical band. Figure taken from Linfield \\& Perley 1984.} \\end{figure}   During the last decade numerous attempts were made to compare the X-ray properties of radio quiet  Seyfert~1 galaxies and  bright broad-line radio galaxies (BLRG) (e.g. Zdziarski et al. 1995, Wo\\'zniak et al. 1998, Eracleous et al. 2000). It was shown that the typical spectral shape of radio quiet  Seyfert~1 is well described  in the 2--500 keV energy range by an intrinsic  power law with $\\Gamma\\sim 1.9$ with an exponential cutoff energy of the order of several hundred keV and a Compton reflection component due to  reflection of the power-law photons from cold matter covering a solid angle of $\\sim 2 \\pi$ (e.g. Nandra \\& Pounds 1994,  Zdziarski et al. 1995). As for the BLRG, their average spectrum, with a photon index $\\Gamma\\sim 1.7$ power law,  seems to have no, or little indication of the reflected component. If associated with the jet their X-ray emission is likely to be related to non-thermal Compton scattering.  With the present data a thermal origin of the  emission  still cannot be ruled out.  Main attention in all the X-ray observations of 3C~111 has been focused on characterizing its spectral shape. In the 2 -- 20 keV energy range all measured spectra are consistent with a power-law spectrum modified by the effects of neutral absorption (e.g. Weaver et al. 1995, Nandra \\& Pounds 1994, Reynolds et al.~1998).  For the hard X-ray spectrum we have less reliable results. Observations by \\textit{HEXTE}, \\textit{OSSE} and \\textit{BeppoSAX} are not sufficient enough to confirm or to rule out the break in the hard X-ray tail of the source spectrum (Wo\\'zniak et al. 1998, Eracleous et al. 2000, Grandi et al. 2002).  No significant short time variability was reported (see Eracleous et al. 2000 for a discussion on \\textit{RXTE} 1997 data). As for the long term variability, the 2-10 keV flux of 3C~111 varied by a factor of $\\sim$5 during its whole observation history  (e.g. Reynolds et al.~1998).  In this paper we present the results of the \\textit{INTEGRAL} and \\textit{RXTE} observations of 3C~111, and compare them with the results of previous missions. We retrieved the {\\it RXTE} data from the High Energy Astrophysics Science Archive Research Center (HEASARC).The {\\it INTEGRAL} and part of the  \\textit{RXTE} data presented here are unpublished so far. ",
        "conclusions": "\\begin{figure}[h] \\psfig{figure=3C111_hist.ps, width=\\columnwidth} \\vspace{-4cm} \\caption{ History of the 3C~111 flux in 2-10 keV energy range, and its power-law index variation.} \\end{figure}  In Figure 6 the  history of the 3C~111 flux in 2-10 keV energy range, and its power-law index variation is summarized. \\textit{RXTE} was lucky to observe the source during one of its faintest states (January 1999) and its brightest state (March, 2001). However, no significant spectral changes are observed, except the fact that the brightest state is the only one where the introduction of the Fe emission line significantly improves the $\\chi^2$ statistics. But this effect can be easily explained by the higher data quality."
    },
    "0404/astro-ph0404270_arXiv.txt": {
        "abstract": "Following a brief review of the principles of the strong equivalence principle (SEP) and tests for its violation in the strong and weak gravitational field regimes, we present preliminary results of new tests using two long-period binary pulsars: J0407+1607 and J2016+1947. PSR J0407+1607 is in a 669-day orbit around a $\\gapp 0.2$ M$_{\\odot}$ companion, while J2016+1947 is in a 635-day orbit around a $\\gapp 0.3$ M$_{\\odot}$ companion. The small eccentricities of both orbits ($e \\sim 10^{-3}$) mean that these systems reduce previous limits on SEP violation by more than a factor of 4. ",
        "introduction": "The principle of equivalence between gravitational force and acceleration is a common feature to all viable theories of gravity. The Strong Equivalence Principle (SEP), however, is unique to Einstein's general theory of relativity (GR).  Unlike the weak equivalence principle (which dates back to Galileo's demonstration that all matter free falls in the same way) and the Einstein equivalence principle from special relativity (which states that the result of a non-gravitational experiment is independent of rest-frame velocity and location), the SEP states that free fall of a body is completely independent of its gravitational self energy.  Before examining how the SEP can be tested, let us first review the gravitational self energy, $\\epsilon$, which is a useful quantity for distinguishing between strong or weak gravitational fields.  Expressed in terms of the rest-mass energy of a body of mass $M$ and size $R$, $\\epsilon = -GM/Rc^2$. For most bodies, $\\epsilon$ is vanishingly small. For example $\\epsilon_{\\rm human} \\sim -10^{-26}$, $\\epsilon_{\\rm Earth} \\sim -5 \\times 10^{-10}$ and even $\\epsilon_{\\odot} = -2 \\times 10^{-6}$. Only for compact objects does $\\epsilon$ become significant and we enter the ``strong-field'' regime.  For a white dwarf $\\epsilon_{\\rm WD} \\sim -10^{-4}$, for a neutron star $\\epsilon_{\\rm NS} \\sim -0.3$ and for a non-rotating black hole, $\\epsilon_{\\rm BH}=-0.5$. ",
        "conclusions": ""
    },
    "0404/astro-ph0404100_arXiv.txt": {
        "abstract": "It has been noted that the Crab and Vela pulsar proper motions lie along the symmetry axes of their wind nebulae. In an effort to promote this observation to a serious test of kick physics, we are using CXO images and other data to estimate the angle between the proper motion and PWN (i.e. spin) axis for a number of pulsars. Here we give a progress report on this work and the constraints that these data provide on kick models. Present data suggest that a kick duration of $\\tau_K \\sim 3$s is sufficient to explain the alignment of most pulsars. This rules out E-M and hydrodynamic kick models, but is fairly consistent with proposed anisotropic $\\nu$ emission. However, some objects, especially PSR J0538+2817 show such good alignment that even $\\nu$ models are challenged. ",
        "introduction": "Typical pulsar velocities of $\\sim 500$km/s represent a lot of momentum, and the nature of the kick that gives neutron stars such speeds has long been one of the major problems in compact object physics. The distribution of kick speeds has important implications for the observed pulsar population, especially those in binaries; thus measurement of pulsar proper motion distributions and application to binary modeling sums has been a major activity (see Podsiadlowski; Burgay; Dewi The {\\it etc.}, these proceedings). However, ${\\vec v}$ is a vector quantity and comparison of its orientation with respect to that other relic of neutron star birth ${\\vec \\Omega}$, promises to provide additional insight into the kick physics.  The letter of Spruit \\& Phinney (1998) was influential in promoting thinking about the spin-kick connection. These authors, in fact, hypothesized that neutron star initial angular momentum was small due to strong core-envelope coupling in pre-collapse stars.  They suggested that an off-center kick, at impact parameter $d=R {\\rm sin}\\psi_{kick}$ imposed while the bloated proto-NS has radius $\\sim 3 \\times 10^6$cm produces a spin of $$ \\Omega_{rms} \\approx 42 {\\rm s^{-1}} \\left ( {{{\\rm sin}\\psi_{kick}} \\over 0.5} \\right ) \\left ( R_{10}/3 \\right ) \\langle v^2 \\rangle_7^{1/2} $$ when the resultant kick velocity was $100\\langle v^2 \\rangle_7^{1/2}$km/s. This gives a modest initial spin period $P_0 \\sim 150/v_7$ms. For a single impulse, the resulting ${\\vec \\Omega}$ is always orthogonal to the space velocity. Of course for long duration kicks $\\tau_K \\gg P_0$ the transverse component of the kick rotationally averages to 0, leading to an aligned spin.   More recent treatments of core coupling through collapse (eg. Heger et al. 2004) do not support the idea of very slow initial spin, instead suggesting $P_0 \\approx 3-10$ms. Such a pre-existing spin will make rotational averaging of the transverse kick component even more effective, increasing the tendency to an aligned proper motion. Lai, Cordes \\& Chernoff (2001) have discussed several physical mechanisms for producing a kick at core collapse. The most important are the Harrison-Tademaru electromagnetic kick (requiring $P_0$ of a few ms), hydrodynamically-driven anisotropy induced kicks ($\\tau_K \\sim \\tau_{dyn}|_{R\\sim 100 {\\rm km}} \\sim 100$ms) and magnetic field-induced neutrino anisotropy kicks ($\\tau_K \\sim \\tau_\\nu \\sim 3$s). They discussed rotational averaging of these kicks, concluding that the spin-kick orientation could be a significant constraint on these models.  For pulsars born in close binaries with aligned angular momenta, the Blaauw mechanism guarantees a component of the proper motion perpendicular to the (pre-SN) spin axis. Similarly the binary-like structure of a maximally rotating core with a strong $m=1$ perturbation can, when the lower mass proto-NS disrupts, induce a kick to the main core, as recently discussed by Colpi \\& Wasserman (2002). This can be thought of as an `intra-core Blaauw mechanism' and similarly gives rise to a kick component orthogonal to the initial spin. So there are viable models for both aligned and orthogonal momenta.  The spectacular CXO images of the Crab and Vela PWNe show clear symmetry axes. It was promptly noted that the proper motion vectors (from HST for the Crab and the LBA for Vela; Dodson et al 2003) were roughly aligned. A more careful assessment (Ng \\& Romani 2004) however shows a statistically significant misalignment; the chance probability of getting two such 2-D projected alignments is $\\sim 3$\\%. Thus, the alignment can provide a significant probe of core collapse physics, but more and better measurements are clearly needed. ",
        "conclusions": ""
    },
    "0404/astro-ph0404336_arXiv.txt": {
        "abstract": "{The spatial dependence of luminosity and mass functions of evolved open clusters is discussed in this work using J and H 2MASS photometry, which allows a wide spatial coverage and proper background determination. The target objects are the overlooked open cluster NGC\\,2180 ($\\ell=203.85^\\circ$, $b=-7.01^\\circ$) and the intermediate-age open cluster NGC\\,3680 ($\\ell=286.76^\\circ$\\ and $b=16.92^\\circ$), which has been reported as being in an advanced state of dissolution. We conclude that, although in an advanced dynamical state (mass segregated), NGC\\,3680 does not present strong signs of dissolution, having luminosity and mass functions very similar to those of the $\\age\\approx3.2$\\,Gyr open cluster M\\,67. On the other hand, NGC\\,2180 presents flat, eroded luminosity functions throughout its structure, indicating that in addition to mass segregation, Galactic tidal stripping has been effective in depleting this cluster of stars. Accordingly, NGC\\,2180 may be the missing link between evolved open clusters and remnants. For NGC\\,2180 we derive an age $\\age\\approx710$\\,Myr and an observed stellar mass of $\\mobs\\sim47\\ms$. Most of the colour-magnitude diagram features,  the main sequence in particular, are equally well reproduced by isochrones with metallicity $\\zz=-0.4$ and 0.0. The solution for $\\zz=-0.4$ results in $\\mM=9.40\\pm0.20$, $\\ebv=0.18$ and a distance to the Sun $\\ds=0.76\\pm0.06$\\,kpc, while that for $\\zz=0.0$ gives $\\mM=10.10\\pm0.20$, $\\ebv=0.0$ and $\\ds=1.05\\pm0.08$\\,kpc. For NGC\\,3680 we derive an age $\\age\\approx1.6$\\,Gyr, $\\ebv=0.0$ and $\\ds=1.00\\pm0.09$\\,kpc, confirming previous estimates. The observed stellar mass $\\mobs\\approx130\\,\\ms$ agrees with previous values. We study both clusters in the context of dynamical states estimated from diagnostic-diagrams involving photometric and structural parameters. Both clusters are dynamically evolved systems. In particular, NGC\\,2180 is closer to open cluster remnants than NGC\\,3680. ",
        "introduction": "\\label{intro}  Open clusters are formed along the gas and dust-rich Galactic plane and contain from tens to a few thousands of stars distributed in an approximately spherical structure of up to a few parsecs in radius. This loose condition makes them potentially short-lived stellar systems, since disruptions may occur by the cumulative effect of passages near interstellar clouds and/or by shocks with the Galactic disk. Cumulative orbital perturbations may lead to more internal orbits, enhancing such effects Bergond et al. (\\cite{Bergond2001}). Consequently, most of the open clusters in the Galaxy evaporate completely in less than 1\\,Gyr. Indeed, the open cluster catalogue of Lyng\\aa\\ (\\cite{Lyngaa1987}) indicated about 70 objects older than 1\\,Gyr ($\\approx6\\%$ of the total number).  The dynamical evolution of an open cluster depends both on internal and external factors. Internal factors are: {\\it (i)} after successive 2-body encounters with more massive stars, less-massive stars may acquire velocities larger than the cluster's escape velocity, and {\\it (ii)} the normal stellar evolution via mass-loss. The external factors are: {\\it (i)} large-scale encounters with giant molecular clouds (Wielen \\cite{Wielen1991}), and {\\it (ii)} tidal stripping by the Galactic gravitational field. A typical open cluster at the solar radius will cross the Galactic plane 10--20 times before being disrupted and leaving an open cluster remnant (de la Fuente Marcos \\cite{delaF1998}). Bergond et al. (\\cite{Bergond2001}) estimate the destruction time-scale for open clusters in the solar neighbourhood at about 600\\,Myr. Consequently, it is expected that only those open clusters which are born with the largest masses or those located at large Galactic radii will survive longer than a few Gyr (Friel \\cite{Friel1995}).  The numerical simulations of de la Fuente Marcos (\\cite{delaF1996}) have shown that the final cluster remnant composition depends on the initial mass function, fraction of primordial binaries and galactocentric distance. The resulting cluster remnants are rich in binaries and do not appear to contain collapsed objects. Remnants of poorly populated clusters are expected to contain early-type stars, while those of more massive clusters contain late-type stars (de la Fuente Marcos \\cite{delaF1996}), owing to different evolutionary time-scales. In the central region of the more evolved clusters, mass segregation should deplete the low main-sequence (MS) stars thus creating a core rich in compact and giant stars (Takahashi \\& Portegies Zwart \\cite{TakaP2000}).  Mass segregation in a star cluster scales with the relaxation time, defined as $t_{relax}=\\frac{N}{8\\ln N}t_{cross}$, where $t_{cross}=R/v$ is the crossing time (Binney \\& Tremaine \\cite{BinTre1987}). For a typical cluster radius of $R\\sim5$\\,pc and velocity dispersion $v\\sim1$\\,km\\,s$^{-1}$, $t_{relax}\\sim13$\\,Myr for a cluster with $N=10^2$ stars, and $t_{relax}\\sim90$\\,Myr for $N=10^3$ stars.  Recently, several studies called attention to the possibility of detecting open cluster remnants in the Galaxy, e.g. Bica et al. (\\cite{Bica2001}), Carraro (\\cite{Carraro2002}), Pavani et al. (\\cite{Pavani2002}, \\cite{Pavani2003}). A fundamental question to dynamical evolution studies is whether any open cluster can be observed right at the disrupting phase, when the remaining low-mass stars in the cluster's halo get dispersed into the background and the corresponding mass function becomes eroded.  Depletion of low-MS stars in the central parts of a cluster is a sign of advanced dynamical evolution. This has been detected e.g. in NGC\\,3680 (Anthony-Twarog et al. \\cite{Twa1991}) and M\\,67 (Bonatto \\& Bica \\cite{BB2003}), in which the presence of a corona rich in low-mass stars has been confirmed with 2MASS photometry.  The Two Micron All Sky Survey (hereafter 2MASS, Skrutskie et al. \\cite{2mass1997}) has proven to be a powerful tool in the analyses of the structure and stellar content of open clusters (e.g. Bonatto \\& Bica \\cite{BB2003}, Bica et al. \\cite{BBD2003}). Indeed, the uniform and essentially complete sky coverage provided by 2MASS allows one to properly take into account background regions with suitable star count statistics, which is fundamental in order to correctly identify and characterize the stellar content of clusters, since their ages and distances can be determined by fitting isochrones to their colour-magnitude diagrams (CMDs), with a precision depending on the depth of the photometry and field contamination.  In the present study we address the actual dynamical state of NGC\\,3680 analysing a large spatial area in the direction of the cluster, which 2MASS can provide. In particular, we search for the presence of a low-mass star-rich corona. In addition, we discuss NGC\\,2180, an overlooked open cluster. This cluster appears to be in a more advanced dynamical evolutionary stage than NGC\\,3680, and thus might be a missing link between evolved open clusters with a corona and final remnants.  In Section~\\ref{targets} we provide available information on NGC\\,2180 and the intermediate-age open cluster NGC\\,3680. In Sect.~\\ref{2massPh} we obtain the 2MASS photometry and introduce the $\\jj\\times\\jh$ CMDs. In Sect.~\\ref{StructAnal} we discuss the radial density distribution of stars and derive structural parameters for the clusters. In Sect.~\\ref{Fund_par} we fit isochrones to the near-infrared CMDs and derive cluster parameters. In Sect.~\\ref{LumFunc} we derive the luminosity and mass functions (hereafter LF and MF) and estimate the stellar masses of each cluster. In Sect.~\\ref{Comp} we compare NGC\\,2180 and NGC\\,3680 with well-known dynamically evolved open clusters and open cluster remnants. Concluding remarks are given in Sect.~\\ref{Conclu}. ",
        "conclusions": "\\label{Conclu}  As a consequence of the internal dynamical evolution and the relentless Galactic tidal pull, most open clusters are expected to evaporate completely in less than 1\\,Gyr, leaving behind a remnant with characteristics which depend on the cluster's initial conditions. Only the more massive open clusters may survive to old ages. In this context, the observation of an actually dissolving open cluster becomes very interesting to check existing theories on dynamical evolution of stellar systems, N-body codes in particular, as well as to test stellar evolution theories.  In the present work we analyse the physical structure, stellar content and dynamical state of the overlooked open cluster NGC\\,2180. We also examine in detail the intermediate-age open cluster NGC\\,3680, formerly considered to be in the last stages of its dynamical evolution. The present analyses make use mostly of J and H 2MASS All Sky data release photometry.  NGC\\,2180 presents a non-uniform radial distribution of stars (Sect.~\\ref{StructAnal}), with significant  $1-\\sigma$ Poissonian error bars due to the small number of member stars. Its radial density profile has a central concentration of stars, as well as a corona (Fig.~\\ref{fig4}). From a  King law fit we estimate $\\rc=0.7\\pm0.3$\\,pc and a linear limiting diameter of $9.5\\pm1.2$\\,pc. Its $\\mj\\times\\jh$ CMD (Fig.~\\ref{fig5}) is depleted of stars near the turnoff, and can be fitted with 710\\,Myr isochrones of solar and sub-solar metallicity. The $\\zz=0.0$ solution results in $\\mM=10.10\\pm0.20$, $\\ebv=0.0$ and a distance to the Sun $\\ds=1.05\\pm0.08$\\,kpc, while the $\\zz=-0.4$ solution gives $\\mM=9.40\\pm0.20$, $\\ebv=0.18$ and $\\ds=0.76\\pm0.06$\\,kpc. Thus, we adopt as distance to the Sun $\\ds=0.91\\pm0.15$\\,kpc, which puts NGC\\,2180 at a galactocentric distance of $8.8\\pm0.1$\\,kpc. Mass segregation and advanced Galactic tidal stripping on NGC\\,2180 are reflected on the spatial properties of its LFs (Fig.~\\ref{fig6}). Low-mass stars are severely depleted from the MS in each region, from the center to the cluster's limiting radius. In addition, the MS of the corona, although depleted as well, is slightly more populated of low-mass stars than the MS of more internal regions. The observed stellar mass in NGC\\,2180 is $\\sim47\\pm7\\,\\ms$.  NGC\\,3680 presents a uniform radial distribution of stars, with a well-defined core and a corona, with $\\rc=0.7\\pm0.1$\\,pc and a linear limiting diameter of $12.8\\pm1.3$\\,pc (Sect.~\\ref{StructAnal}). Its $\\mj\\times\\jh$ CMD (Fig.~\\ref{fig5}) presents a nearly complete MS, including the turnoff and giants. We derive an age of $\\age\\approx1.6$\\,Gyr, $\\mM=10.00\\pm0.20$, $\\ebv=0.00$ and $\\ds=1.00\\pm0.09$\\,kpc, in reasonable agreement with previous works. The LF of the central region of NGC\\,3680 (Fig.~\\ref{fig6}) is depleted of low-mass stars which, contrary to what is observed in NGC\\,2180, are still present in the external region and corona. Thus, Galactic tidal stripping has not yet been effective in severely depleting NGC\\,3680 of stars. For NGC\\,3680, a MF fit $\\phi(m)\\propto m^{-(1+\\chi)}$ resulted in a slope $\\chi=2.06\\pm1.08$, and in an observed stellar mass (MS and giants) of $\\approx130\\pm24\\,\\ms$. Extrapolating the MF fit down to the theoretical low-mass end $m_{low}=0.08\\,\\ms$, the total stellar mass in NGC\\,3680 turns out to be $\\sim(2.4\\pm1.2)\\times10^3\\,\\ms$, which agrees with previous estimates, within uncertainties.  Assuming a more representative IMF (Kroupa \\cite{Kroupa2001}), which flattens for masses below $\\sim0.5\\,\\ms$, the total stellar mass in NGC\\,3680 turns out to be $546\\pm206\\,\\ms$.  Finally, comparing NGC\\,2180 and NGC\\,3680 with clusters in other dynamical states, {as well as open cluster remnants}, we found that the less-massive nature of NGC\\,2180 put it closer to cluster remnants than NGC\\,3680.  The above arguments lead us to conclude that NGC\\,2180 is in a more advanced dynamical state than NGC\\,3680, on its way to become a fossil cluster. Thus, NGC\\,2180 may be the missing link between evolved open clusters and remnants."
    },
    "0404/astro-ph0404322_arXiv.txt": {
        "abstract": "This is the second in a series of papers presenting a new calculation of the mass of the Galaxy based on radial velocities and distances for a sample of faint $16<B<21.3$ field blue horizontal-branch (BHB) stars. We present accurate $BV$ CCD photometry and spectra for 142 candidate A--type stars selected from $ub_jr$ photometry of UK Schmidt telescope plates in six high--Galactic--latitude fields. Classification of these candidates produces a sample of 60 BHB stars at distances of 11--52 kpc from the Sun (mean 28 kpc), with heliocentric line-of-sight velocities accurate to 15 km\\,s$^{-1}$, and distance errors $<10\\%$. We provide a summary table listing coordinates and velocities of these stars. The measured dispersion of the radial component of the Galactocentric velocity for this sample is 108 $\\pm$ 10 km s$^{-1}$, in agreement with a recent study of the distant halo by Sirko and coworkers. Measurements of the Ca II K line indicate that nearly all the stars are metal-poor with a mean [Fe/H] = -1.8 with dispersion 0.5.  Subsequent papers will describe a second survey of BHBs to heliocentric distances $70<R<125$ kpc and present a new estimate of the mass of the Galaxy. ",
        "introduction": "In contrast with the detailed knowledge we have of the stellar content of the halo of the Galaxy, the dark matter content is far less well characterised, with both the total mass and the size of the halo poorly determined quantities.  The accurate measurement of the mass profile would provide important clues to the nature of the dark matter.  For instance, an accurate estimate of the mass within a radius of 50 kpc is required to establish the fraction of the mass in compact objects identified from microlensing experiments (e.g Alcock et al. 2000).  Quantifying the distribution of dark matter in the Galaxy is essential for understanding the assembly of the various baryonic components through comparison with simulations.  This paper is the second in a series of four, presenting a new measurement of the mass profile of the Galaxy out to the largest radii, $r>100$kpc,\\footnote{In this paper we use the coordinate $r$ to denote Galactocentric distances and the coordinate $R$ to denote heliocentric distances} using surveys for remote halo blue horizontal branch (BHB) stars.  The main shortcoming of previous analyses of the mass of the Galaxy is the small size of the sample of objects at large radii, used as dynamical tracers.  Wilkinson and Evans (1999, hereafter WE99) calculate the total mass of the Milky Way to be $M_{tot}=1.9^{+3.6}_{-1.7} \\times 10^{12} M_{\\odot}$, using the full set of 27 known satellite galaxies and globular clusters at Galactocentric radii $r>$ 20 kpc (six of which possess measured proper motions).  This sample must be nearly complete, so a new population of distant halo objects must be found in order to increase the number of dynamical tracers.  Field BHB stars are ideal for this purpose.  These A--type stars are luminous, $M_V=0.9$ (\\S5), which ensures that they can be detected to large distances, and have a small spread in absolute magnitude, so their distances may be determined accurately.  In a recent paper Sakamoto et al. (2003) added new kinematic data to the dataset used by WE99, and obtained a considerably more precise estimate of the total mass of the Galaxy of $M_{tot}=2.5^{+0.5}_{-1.0} \\times 10^{12} M_{\\odot}$.  Primarily, the new kinematic data comprise radial velocities of 412 BHB stars, of which 211 have proper motions, at heliocentric distances $R<10$kpc, with a median distance of $\\sim$ 4.5kpc. Sakamoto et al. adopted the WE99 mass model, where the total mass enclosed within Galactocentric radius $r$ is $M(r)=M_{tot}/(1+a^2/r^2)^{1/2}$.  Here $a$ is the scale length, for which the best-fit value of 225kpc was obtained.  For $r=20$kpc, i.e. a radius containing all the BHB stars used in the analysis, this gives $M(r)=0.09M_{tot}$.  This indicates that most of the improved precision derives from kinematic data within a radius containing only one tenth of the total mass, and therefore relies on extrapolation of a model that is tightly constrained only at small radii. In a review of mass estimates of the Milky Way halo, Zaritsky (1999) emphasises the dangers of such extrapolation (see also Bellazzini, 2004, for a useful discussion). This motivates a new survey for remote BHB stars, at distances that are a substantial fraction of the halo scale length $a$, in order to obtain a direct measure of the enclosed mass at large radii.  Although BHB stars are abundant in the Galaxy halo, selection of a clean sample is not straightforward.  Samples of field A--type stars in the halo are easily identifiable in UBR (or equivalent e.g.  $ugr$) multicolour datasets (e.g.  Yanny et al. 2000).  Samples selected by broadband colour include not only BHB stars but also stars of main sequence surface gravity that are some 2 magnitudes less luminous, predominantly field blue stragglers (hereafter A/BS), as well as a small proportion of quasars.  The reason that no large sample of remote $r>30$kpc BHB stars has yet been compiled, is that the methods developed for separating the two populations of A stars (e.g.  Kinman, Suntzeff, and Kraft, 1994) require signal-to-noise ($S/N$) ratios that are unfeasibly high for such faint stars ($B>18$).  In the first paper in this series (Clewley et al. 2002, hereafter Paper I), we developed two methods that overcome the difficulties.  The methods require relatively modest telescope resources, yet produce samples with high completeness and low contamination.  The methods are applicable specifically to the classification of halo stars with strong Balmer lines, defined by EW $H\\gamma>13$\\AA, i.e.  A stars in the approximate colour range $0.0<(B-V)_0<0.2$.  Both methods employ a Sersic function fit to the H$\\gamma$ and H$\\delta$ absorption lines.  The first method, the {\\em $D_{0.15}$--Colour} method, plots the average of the width of the two Balmer lines against $(B-V)_0$ colour.  The A/BS stars, having higher surface gravity, separate from the BHB stars because of their broader Balmer lines.  The EW of the CaK line is used to filter out a small number of interlopers in the BHB sample.  We used Monte Carlo methods to establish that with $(B-V)_0$ colours accurate to $0.03$ magnitudes, and spectra of $S/N=15{\\rm\\AA}^{-1}$, samples of BHB stars selected by this method would be about $87\\%$ complete, with a contamination of $7\\%$ by A/BS stars.  Contamination at this level can safely be accounted for in the dynamical analysis.  Spectra of this $S/N$ are also suitable for measurement of the radial velocities.  The second method, called the {\\em Scale width--Shape} method, plots two parameters of the Sersic fit, the scale width $b$, against the power--law exponent $c$.  The method is almost as efficient as the {\\em $D_{0.15}$--Colour} method, with $82\\%$ completeness and $12\\%$ contamination, for the same spectroscopic $S/N$ of $15{\\rm\\AA}^{-1}$. Again, the EW of the CaK line is used to filter out a small number of interlopers.  The advantage of the {\\em Scale width--Shape} method is that colours are not needed.  For samples of stars with existing accurate $(B-V)_0$ colours, the first method is preferred, as the contamination is lower, and the completeness higher.  The gain is small, however, and accurate photometry is time consuming.  Where accurate colours are not already available, if telescope resources are limited, the best practical solution will be simply to obtain spectra, and use the {\\em Scale width--Shape} method.  In this second paper we describe the compilation of a sample of BHB stars in the magnitude range $16<B<19.5$, corresponding to heliocentric distances $11<R<52$ kpc, using the two classification methods described above.  This first survey (hereafter the UKST survey) starts from photographic plates taken with the United Kingdom Schmidt Telescope (UKST).  In the third paper we describe a second survey, for fainter BHB stars, $20<g*<21.1$, corresponding to heliocentric distances $70<R<125$ kpc.  This second survey starts from the Sloan Digital Sky Survey (SDSS) Early Data Release (EDR) (Stoughton et al. 2002).  The fourth paper will present the dynamical analysis of these two data sets, combined with the 27 satellites and globular clusters analysed by WE99.  The layout of the remainder of the paper is as follows.  In section 2 we describe the selection of the UKST fields, the reduction of the plate scans, and the production of the multicolour photometric datasets.  We provide catalogues of a total of 461 objects with the colours of A--type stars, which are candidate halo BHB stars.  Section 3 covers CCD $B-V$ photometry of 280 candidate BHB stars, and section 4 describes the spectroscopic observations of 156 candidates.  We describe the measurement of the metallicity and the radial velocity from the spectra.  Section 5 reviews the measurement of distances of BHB stars.  In Section 6 we use these results and the methods developed in Paper I to classify these stars, and we provide a table of distances and radial velocities of the 60 confirmed BHB stars. Finally, in Section 7 we provide a summary of the main conclusions of the paper. ",
        "conclusions": "In this paper we have presented the results of a survey of remote halo A-type stars, beginning with catalogues of $ub_jr$ photometry from UKST plates in two northern and four southern high-Galactic-latitude fields.  Accurate CCD photometry and spectroscopy of candidate A-type stars produced a sample of 142 stars with data of suitable quality for classification into the classes BHB and A/BS. The final sample, Table 14, comprises 60 stars classified BHB, at distances of $11-52$kpc from the Sun (mean distance 28 kpc), with heliocentric radial velocities accurate to 15 km s$^{-1}$, and distance errors $<10\\%$. These stars are suitable for a dynamical study of the mass distribution in the Milky Way at large radii. The stars are all metal poor, with mean [Fe/H]=-1.8 with dispersion 0.5. The measured dispersion of the radial component of the Galactocentric velocity for this sample is 108 $\\pm$ 10 km s$^{-1}$ which is consistent with Sommer-Larsen et al. 1997 and Sirko et al. 2004b. The remaining 82 stars are classified A/BS. These stars have very similar metallicities and velocity dispersion to the BHB stars."
    },
    "0404/astro-ph0404608_arXiv.txt": {
        "abstract": "s{This rapporteur review summarizes results presented in Parallel Session GBT2 (Gamma Ray Burst Theory 2) on the Collapsar and Supranova Models held 25 July 2003 at the 10th Marcel Grossmann Meeting on General Relativity, Rio de Janeiro, Brazil.  A central issue in GRB studies is the process whereby energy is released from the GRB engine. One scenario is the collapsar model, where the evolved stellar core promptly collapses to a black hole surrounded by a massive, intermittently-accreting torus of nuclear density material. A second scenario is the supranova model, where the first step of a two-step collapse process leaves behind a rapidly rotating neutron star stabilized by rotation, which later collapses to a black hole while making the GRB. In the supranova model, a powerful pulsar wind lasting days to weeks after the supernova makes distinctive signatures from the heating of the supernova remnant (SNR) shell, which should be discovered with Swift. This model also predicts nonstandard reddened excesses from the cooling SNR due to the range of delays between the two collapse events, contrary to observations of SN emissions from low redshift GRBs which favor the collapsar model. The observational basis for both models is critically reviewed. Problems of the collapsar/internal shock scenario powered by the Blandford-Znajek process, and a two-collapse supranova scenario powered by the B-Z process or the Penrose mechanism, or even Ruffini's pair electromagnetic pulse, are briefly discussed. In view of the standard $\\gamma$-ray energy upper limit, an external shock model involving jetted relativistic ejecta powered by an explosive event provides the most consistent explanation for GRB observations.  An outline of a GRB model is proposed where the second collapse takes place within minutes to hours after the primary SN event, during which time loss of centrifugal support along the rotation axis of the neutron star provides a relatively baryon-clean polar environment along which a newly formed, rapidly spinning black hole drives collimated baryon-dilute outflows.} ",
        "introduction": "The contributions\\cite{bot04,laz04} entitled ``The Diagnostic Power of X-ray Emission Lines in GRBs,\" by Dr.\\ Markus B\\\"ottcher (Ohio University, USA), and ``Gamma-Ray Burst Progenitors confront Observations,\" by Dr.\\ Davide Lazzati (IoA, Cambridge, UK), frame the current debate. B\\\"ottcher's article\\cite{bot04} summarizes X-ray observations, and the case is made for a supernova (SN) taking place weeks to months before the GRB, so that the circumburst medium (CBM) contains dense ($\\gtrsim 10^{10}$ cm$^{-3}$), distant ($\\approx 10^{15}$ -- $10^{17}$ cm) SNR shell material whose illumination could account for the appearance of X-ray features in GRB prompt and afterglow radiation. These features give important clues about the GRB environment, and would be compelling if their significance were only greater. The marginal (3 -- 5 $\\sigma$) detections near threshold make for concern about the reality of the features.  Swift should clear this issue up, with its $24^\\prime\\times 24^\\prime$ FoV X-ray telescope reaching $2\\times 10^{-14}$ ergs cm$^{-2}$ s$^{-1}$ in a $10^4$ second observation, with autonomous slewing capability to get on target within one minute(!).  At the time of this writing (April, 2004), it seems probable that all GRBs have associated SNe as identified by optically thick photospheric SN emissions appearing as reddened optical excesses in GRB afterglow light curves some 10 -- 30 days after the GRB event. In low redshift GRBs, and most clearly in the case of GRB 030329\\cite{hjo03}, the SN emissions are well fit by SN 1998bw light curves K-corrected and stretched to match the redshift of the GRB source. These emissions are powered by trapped shock energy and radioisotope heating (principally due to the $^{56}$Ni$\\rightarrow^{56}$Co$\\rightarrow ^{56}$Fe decay chain formed by explosive nucleosynthesis or accretion disk winds\\cite{psm04}).  The optical follow-on observations indicate that any delays between collapses to a neutron star and black hole (or strange star) last for less than a few days. This places constraints on the angular momentum and mass at formation, and the equations of state of GRB progenitors\\cite{sha04,wmd04}.  Lazzati\\cite{laz04} argues that the issue of SN excesses in late (10 -- 30 days) time optical afterglow light curves is settled, and he examines the low redshift $(z =0.17)$ GRB 030329 in detail. As yet, SN excesses have not been observed from a GRB which also displays X-ray features. One possibility is that X-ray features will be detected at high significance by Swift from a GRB with a nearly coincident (within $\\lesssim$ a few days) SN, in which case the original supranova model\\cite{vs98}, with delays between the two collapse events ranging from weeks to months, is ruled out. If Fe features are firmly detected, then this might imply ejection of material from an Fe core (which then does not explain what powers the SN Ic light curve). Another possibility is that the GRBs displaying Fe features form a disjoint set from those which display SN emissions. In this case, two separate population tracks for GRBs which produce SN emissions and those which produce X-ray line features from the illumination of a distant reprocessor are implied\\cite{wmd04}. It is also possible that the X-ray features are artifacts of analysis.  As is apparent, the observational situation, particularly with regard to the X-ray emission lines and absorption features and edges in GRB spectra, is murky. Then what about the theoretical situation?  The relativistic blast-wave external-shock model with both forward and reverse shock emissions is generally accepted as the basis to interpret afterglow observations, and even to apply during the prompt phase in those GRBs which display enhanced optical reverse shock radiation.  At least, this model provides a serviceable tool to interpret radiation fluxes and to derive physical and environmental parameters ($\\epsilon_e$, $\\epsilon_B$, $n_{CBM}$, and to distinguish uniform vs.\\ windfed CBM). The radiation processes that make $\\gamma$-rays during the prompt $(\\approx t_{90}$) phase is an open question, with external and internal shock models each proposed. An external shock model, powered by a single explosive event, is favored on the basis of the data, as argued below.  In this article, supranova-model predictions of signatures of SNR shells heated by plerionic emissions emitted by a newly born, rapidly rotating and highly magnetized neutron star are presented. The heating of the SNR, and the interaction of the GRB emissions with this shell of material imply many predictions for the supranova model. We critique the collapsar and supranova models and argue that they are both inconsistent with the data (even leaving aside the X-ray observations). The observations do not (yet) show the predicted plerionic signatures of a supranova model, nor do they imply standard reddened emissions at late times if a fraction of GRBs have delays of weeks to months between the two collapse events (that fraction, and the validity of the original supranova model, continues to decline in significance as more observations of reddened excesses accumulate). The collapsar model cannot explain the prompt GRB emissions, durations, efficiencies, nor the standard maximum-energy reservoir result. It predicts a wind-type CBM, which is only seen on rare occasion\\cite{pri02}. The baryon-dilute jet must be able to penetrate the massive stellar envelope without substantial mixing and reach $\\Gamma$ factors $\\gg 100$.  Some new ideas are needed. One possibility is that energy is dissipated by general relativistic effects within the ergosphere of a charged, rotating black hole as it is formed, rather than extracted when a magnetized disk is accreted onto a rotating black hole, as in the collapsar model\\cite{mwh01}. Numerical simulations of core collapse of massive stars show that material can be evacuated along the rotation axis of the stellar core, leaving a relatively baryon-clean environment\\cite{bur04}. Within a few minutes to 1 -- 2 days of the SN, a second collapse would drive a jet in the poleward direction of the newly formed, rapidly spinning black hole to make the GRB.  This model would make jet formation more feasible than in a collapsar scenario where nearly baryon-pure ejecta, whether in the form of a pair fireball or a Poynting-flux dominated flow, has to ``drill a channel\" through a massive stellar envelope and avoid the processes that are so effective at mixing $^{56}$Co in SN 1987A, namely strong Rayleigh-Taylor mixing instabilities along the jet head and Kelvin-Helmholtz instabilities along the jet spine and in the Rayleigh-Taylor fingers.  Such a scenario does not depart from the well-established connection of long-duration GRBs with high-mass stars born in star-forming galaxies that end their lives as Type Ic SNe. Indeed, rapidly spinning Wolf-Rayet progenitors stars are the most likely GRB progenitors, as in the collapsar\\cite{woo93} and some versions (though not all\\cite{vs99}) of the supranova models. The main difference is in the nature of the collapse process and the origin and meaning of the prompt GRB emissions, with important implications for the physics of stellar evolution and black-hole formation. Whether the GRB is formed by an engine or explosion is important for understanding the physics of black hole formation. Here we argue for an explosive event (cf.\\ Ref.[\\cite{kul04}]).  Such a model brings into focus the ``standard-energy reservoir result''\\cite{fra01} or, more accurately, the ``standard-energy upper limit''\\cite{bfk03}, and avoids many unpleasant issues that are often avoided in discussions of the collapsar scenario, such as low efficiency and absence of curvature signatures that are expected from colliding shells in the internal/external scenario, and the frequent presence of a rather uniform CBM as inferred from afterglow model fits to the external shock emissions\\cite{pk01}.  I also comment on current models for energy extraction from a rotating black hole, including the Blandford-Znajek effect and the Penrose mechanism.  Both the collapsar and supranova models have difficulties to explain the GRB phenomenon. We\\cite{dm99,dm03} have argued that the pulses in a GRB light curve are flashes made when a relativistic blast wave intercepts density inhomogeneities in the CBM.  Delayed low-level radiations from fallback onto the newly formed black hole are possible (and this could help explain X-ray features should they turn out to be real). The GRB light curves are compatible with an explosion or a magnetar discharge\\cite{uso92}. The relativistic ejecta shell can support a reverse shock structure to explain prompt optical and early radio afterglows.   Reviews of GRB physics and observations that focus on observations\\cite{pkw00} and theory\\cite{zm04}, both from the perspective of the external shock model\\cite{der02} and the internal/external scenario\\cite{piran99,mes02}, can be consulted for further background. ",
        "conclusions": ""
    },
    "0404/astro-ph0404264_arXiv.txt": {
        "abstract": "{As a result of a radial velocity follow-up of OGLE planetary transit candidates in Carina, we report the discovery of two new transiting planets with very short orbital periods: OGLE-TR-113 with m=1.35 $\\pm 0.22 M_{Jup}$, r=$1.08^{+0.07}_{-0.05} R_{Jup}$, P=1.43 day, and OGLE-TR-132 with m=1.01 $\\pm 0.31 M_{Jup}$,  r=$1.15^{+0.80}_{-0.13} R_{Jup}$, P=1.69 day. These detections bring to three the number of known \"very hot Jupiters\" (Jovian exoplanets like OGLE-TR-56 with periods much smaller than 3 days), indicating that the absence of such planets in radial velocity survey does not reflect an absolute limit. ",
        "introduction": "Since 1995, more than 120 planetary candidates have been discovered, mainly by radial velocity surveys. Although they provide a lot of information concerning the orbital parameters and the host star properties, such surveys do not yield the real mass of the planet (only $m\\sin i$) and do not give any information about its size. The discovery of HD209458 by 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 extra-solar planet, illustrating the real complementarity of the two methods. The OGLE survey (Optical Gravitational Lensing Experiment) announced the detection of 137 short-period multi-transiting objects (Udalski et al., \\cite{udalski1}, \\cite{udalski2}, \\cite{udalski3}, \\cite{udalski4}). Recently Konacki et al. (\\cite{konacki1}, \\cite{konacki2}), thanks to a Doppler follow-up, announced the characterization of the first extra-solar planet with the unexpected extremely short period of 1.2 day, much below the lower end of the period distribution of planets detected by Doppler surveys (Udry et al. \\cite{udry}). We present in this letter two new cases of extra-solar planets with very short orbital periods: OGLE-TR-113 and OGLE-TR-132. ",
        "conclusions": "\\begin{table} \\begin{tabular}{c c c c c} Name& M & R & m & r \\\\ & [$M_\\odot$] & [$R_\\odot$]& [$M_{Jup}$]& [$R_{Jup}$] \\\\ \\hline \\hline \\\\[-3mm] OGLE-113  & 0.77$\\pm$0.06 & 0.765$\\pm$0.025        & 1.35$\\pm$0.22 & 1.08$^{+0.07}_{-0.05}$ \\\\[1mm] OGLE-132  & 1.34$\\pm$0.10 & 1.41$^{+0.49}_{-0.10}$ & 1.01$\\pm$0.31 & 1.15$^{+0.80}_{-0.13}$ \\\\[1mm] \\end{tabular} \\caption{Summary table with the physical parameters $M$, $R$, $m$ and $r$.} \\label{tablezoo} \\end{table}   The parameters of OGLE-TR-113b are very accurately defined, first because the transit shape is clearly delineated by the photometric data, second because there is only a narrow range of possible parameters allowed by stellar evolution models for a cool K dwarf. As a result, $r$ could be computed with a very small formal uncertainty.   No radial velocity variations were detected in OGLE-TR-131. Moreover, the spectroscopy indicates that it is most probably a sub-giant ($\\log g=3.30\\pm 0.73$), which renders the existence of a close companion unlikely. OGLE-TR-131 is included here to show that our estimates of the radial velocity uncertainties  are coherent with the residuals in the absence of detectable orbital motion.   The photometric transit signal of OGLE-TR-132 is near the detectability limit, but the existence of a radial velocity variation at the precise period and phase of the transit gives confidence in the reality of the transiting companion. In contrast with OGLE-TR-113, the parameters are not very well constrained. The transit shape is too poorly defined to constrain the transit latitude, so that there is a strong degeneracy between the impact parameter $b$ and the primary radius $R$. Furthermore the values of temperature and gravity measured from the spectra are compatible with a wide variety of young to evolved F dwarfs, with radii ranging from 1.3 to 1.9 $R_\\odot$. As a result, the upper uncertainty interval on $r$ is wide, as indicated by the dotted error line in Fig. \\ref{massradius}.  The semi-major axis of OGLE-TR-113b and OGLE-TR-132b are $a=0.0228 \\pm 0.0006$ AU and $a=0.0306 \\pm 0.0008$ AU respectively.    The detection of OGLE-TR-113b and OGLE-TR-132b show that the case of OGLE-TR-56 is not isolated and that \"very hot Jupiters\" (i.e. Jovian exoplanets with periods much smaller than 3 days) are not extremely uncommon. Therefore, the accumulation of hot Jupiters near periods of 3 days (Udry et al. \\cite{udry}) does not reflect an absolute limit for the existence of planets. The parent stars of OGLE-TR-56, 113 and 132 are very different, ranging from F to K dwarfs, indicating that very hot Jupiters are possible around different type of stars.  Fig. \\ref{massradius} gives the mass-radius relation for the four known transiting exoplanets. It is noteworthy that the three OGLE objects do not seem to have such an inflated radius as HD209458b, despite their much shorter periods.  The two new detections bring to 3 the number of \"very hot  Jupiters\" detected in the OGLE transit survey, among 155'000 light curves examined for transits. Given the geometric probability of transit ($\\sim $17\\% for $P=1.5$ days), this results implies a total number of 21 $\\pm$ 10 very hot Jupiters among the targets. The detection completeness of the OGLE survey toward transiting very hot Jupiter should be computed with detailed simulations, but it may be quite high because for such low values of period the phase coverage is very good. Assuming a detection probability of 50\\%, then the total number is $42 \\pm 20$, therefore one in every 2500 to 7000 targets. Even if there is a proportion of giants stars among the OGLE candidate that would have been weeded out of the radial velocity surveys, the absence of \"very hot Jupiters\" among the $\\sim$3000 field dwarfs surveyed in radial velocity in the solar neighborhood is therefore not incompatible with our result. This estimate also indicates that \"very hot Jupiters\" are not out of reach of future radial velocity surveys. It is noteworthy though that for such low periods, photometric transit surveys are a more efficient detection method than radial velocity monitoring.  \\begin{figure} \\resizebox{8.5cm}{!}{\\includegraphics{bouchyfig4.eps}} \\caption{The four known transiting extrasolar planets plotted in the Mass-Radius diagram. The symbols indicate Jupiter and Saturn.} \\label{massradius} \\end{figure}"
    },
    "0404/astro-ph0404578_arXiv.txt": {
        "abstract": "In this paper, we construct full amplitude non-linear hydrodynamical models of fundamental mode Galactic Cepheids and analyze the resulting theoretical period-colour and amplitude-colour relations at maximum, mean and minimum light. These theoretical relations match the general form of the observed relations well. This agreement is, to some extent, independent of the mass-luminosity relations used, pulsation code, numerical techniques, details of the input physics and methods to convert theoretical quantities, such as bolometric luminosity and temperature, to observational quantities, such as V band magnitudes or $(V-I)$ colours. We show that the period-colour and amplitude-colour properties of fundamental mode Galactic Cepheids with periods such that $\\log (P)>0.8$ can be explained by a simple application of the Stefan-Boltzmann law and the interaction of the photosphere with the hydrogen ionization front. We discuss the implications of our results for explaining the behavior of Galactic Cepheid period-colour, and period-luminosity relations at mean light. ",
        "introduction": "\\citet[][paper I, hereafter KN]{kan04} presented new observational characteristics of fundamental mode Cepheids obtained from an analysis of the Galactic Cepheid data and the OGLE LMC/SMC Cepheid data. These were the period-colour (PC) and the amplitude-colour (AC) relations at maximum, mean and minimum light. KN analyzed these PC and AC diagrams in the context of the work of \\citet[][hereafter SKM]{sim93}, who used the Stefan-Boltzmann law and the fact that radius fluctuations are small in Cepheids, to derive the following equation, valid for optical pulsations of Cepheid variables:  \\begin{eqnarray} \\log T_{max} - \\log T_{min} = {1\\over{10}}(V_{min} - V_{max}), \\end{eqnarray}  \\ni  where $T_{max/min}$ is the photospheric temperature at maximum or minimum light. Consequently, if for some reason the PC relation is flat at maximum light, and the colour used is a good predictor of temperature, then equation (1) predicts a relationship between $V$ band amplitude and $T_{min}$, and thus a correlation between the $V$ band amplitude and the observed colour (after extinction correction) at minimum $V$ band light. \\citet{cod47} found that Galactic Cepheids exhibit a spectral type that is independent of period at maximum light. SKM analyzed existing Galactic Cepheid data to show that higher amplitude Galactic Cepheids are indeed driven to cooler temperatures, and thus redder $(B-V)$ colours, at minimum light in accordance with equation (1) and the observational findings of \\citet{cod47}.  The reason why Galactic Cepheids follow a spectral type that is independent of period at maximum light was explained in SKM: it is due to the interaction of the photosphere with the hydrogen ionization front (HIF). At maximum light, the HIF is so far out in the mass distribution that the photosphere, taken to be optical depth 2/3, occurs right at the base of the HIF. Together with the HIF is a sharp rise in opacity. At this point, the mean free path goes to zero. In the absence of any significant driving, even though the surrounding atmosphere has a non-zero inward velocity, this opacity wall prevents the photosphere from going deeper in the star and erases any ``memory'' of global stellar conditions. Thus for a large range of periods, the photosphere occurs at the base of the HIF at maximum light at a temperature which is independent of period. At maximum light, this leads to a flat relation between period \\& temperature, period \\& $(B-V)$ colour and period \\& spectral type, as seen in SKM. At other phases, the HIF lies too far inside the mass distribution to interact with the photosphere. This, together with equation (1), implies a relation between amplitude and $(B-V)$ colour or temperature at minimum light. Equation (1) suggests that if $T_{min}$ obeyed a flat relation with period, then there should be a relation between amplitude and colour at maximum light. Hence PC relations are connected to AC relations, and changes in the PC relations should be reflected in corresponding changes in the AC relations.  Because of the extensive data now available, KN analyzed recent Galactic and Magellanic Cloud Cepheid data in terms of PC and AC diagrams at maximum, mean and minimum light. Using $(V-I)$ colours, they performed a more detailed analysis of the Galactic Cepheid data and found that the Galactic PC relation at maximum light displayed a statistically significant break at 10 days, but was consistent with a single line at mean and minimum light. In the LMC, the PC relation displayed this break at all three phases, though the statistical significance at maximum light was marginal. The SMC PC relation displayed similar properties to that in the Galaxy.  Analysis of the Galactic Cepheid data in terms of AC diagrams confirmed the work of SKM, and extended it to the case of $(V-I)$ colour: there is a relation between amplitude and $(V-I)$ colour at minimum $V$ band light. For the LMC, short period Cepheids ($\\log (P) < 1.0$) show no relation between $V$ band amplitude and $(V-I)$ colour at minimum light, but long period Cepheids ($\\log (P) > 1.0$) are such that higher amplitude stars are driven to redder $(V-I)$ colours and hence cooler temperatures at minimum light. At maximum light, short period Cepheids are such that higher amplitude stars are driven to bluer $(V-I)$ colours and hence hotter temperatures but long period Cepheids do not show such a relation.  An understanding of the PC and AC relations derived by KN are important, not only for stellar pulsation and evolution studies of Cepheids, but also because the Cepheid period-luminosity (PL) relation depends on the PC relation \\citep[see, e.g., ][]{mad91}. The PL relation will reflect significant changes in the PC relation. Hence in studying PC and AC relations in different galaxies, we are studying the universality of the Cepheid PL relation. Recent work by \\citet{tam02,tam03}, \\citet{fou03}, \\citet{kan04}, \\citet{nge04}, \\citet{san04} and \\citet{sto04} have suggested that the PL relation in the Galaxy is significantly different from that in the LMC and, further, that the PL relation in the LMC is non-linear. In contrast, current observations indicate that the Galactic PL relation is linear \\citep{tam03,kan04,nge04}.  The purpose of this and subsequent papers is to confront these new observed characteristics of Cepheids with the latest stellar pulsation models and interpret the results in terms of the theory presented in SKM, which has been summarized above. This approach will ultimately yield a qualitatively deeper understanding some of the reasons behind the variation of the Cepheid PL relation from galaxy to galaxy. Our ultimate goal is to use our theoretical models to estimate the effect of this variation quantitatively. For this paper we concentrate on Galactic Cepheid models which we compare with the same Galactic Cepheid data used in KN. The LMC/SMC Cepheid models will be presented in the forthcoming papers in this series.  Our models and methods improve upon those in SKM in the following respects: First of all, they contain a formulation to model time dependent turbulent convection \\citep{yec99,kol02}, in contrast to the purely radiative models used in SKM. In addition, we construct more models so that we can examine in greater detail the PC and AC characteristics of Cepheids. Secondly, we investigate the pulsation properties as a function of other phase points. When investigating PC and AC relations at mean light, KN defined mean $(V-I)$ colour as $V_{mean}-I_{phmean}$ where $V_{mean}$ is the $V$ band magnitude closest to the value of $A_0$, the mean magnitude obtained from a Fourier decomposition, and $I_{phmean}$ is the $I$ band magnitude at this same phase. The reason for doing this is because the colour of the star at a certain phase is precisely the $V$ band magnitude minus the $I$ band magnitude at that same phase. In contrast, a definition of mean colour such as $<V> - <I>$ folds in the phase difference (albeit small) that exists between the $V$ and $I$ band light curves. We adopt the definitions given in KN for colour at maximum, mean and minimum light. However, there are two phases when the $V$ band magnitude is closest to the value of $A_0$: once on the ascending branch and once on the descending branch of the light curve. The colour (or the temperature), in both models and theory, need not necessary be the same at these two phases. In KN's analysis of the observed data, the descending branch mean was adopted as the ``mean''. We pay specific attention to this detail in our results and discussion section. Finally, we look at $(V-I)$ colours whereas SKM studied predominantly $(B-V)$ colours. The $(V-I)$ colour is a good indicator of temperature \\citep{gon96,tam03} and is a crucial colour for the existing calibration of the extra-galactic distance scale (e.g., \\citealt{fre01}).   \\begin{table} \\centering \\caption{Input parameters for Galactic Cepheid models with periods obtained from a linear analysis. The periods, $P_0$ and $P_1$, are referred to the fundamental and first overtone periods, respectively. Similarly for the growth rate, $\\eta$. Both of the mass and luminosity are in Solar unit, the temperature is in $K$ and the period is in days.} \\label{tab1} \\begin{tabular}{ccccccc} \\hline M & $\\log(L)$ & $T_{e} $ & $P_0$ & $\\eta_0$ & $P_1$ & $\\eta_1$ \\\\ \\hline \\hline \\multicolumn{7}{c}{ML Relation from \\citet{bon00}} \\\\ \\hline 11.5 & 4.279 & 4830 & 43.3924 & 0.042 & 27.58 & -0.164 \\\\ 11.1 & 4.228 & 4900 & 37.8090 & 0.039 & 24.40 & -0.157 \\\\ 10.8 & 4.188 & 4950 & 34.4156 & 0.037 & 22.25 & -0.133 \\\\ 10.5 & 4.147 & 5000 & 30.8189 & 0.035 & 20.26 & -0.119 \\\\ 9.55 & 4.009 & 5075 & 23.4835 & 0.025 & 15.69 & -0.094 \\\\ 9.45 & 3.994 & 5265 & 20.0414 & 0.028 & 13.55 & -0.051 \\\\ 8.60 & 3.856 & 5300 & 15.8313 & 0.022 & 10.78 & -0.040 \\\\ 7.70 & 3.696 & 5332 & 12.1076 & 0.015 & 8.291 & -0.036 \\\\ 7.30 & 3.618 & 5440 & 10.0124 & 0.014 & 6.886 & -0.017 \\\\ 7.30 & 3.618 & 5490 & 9.69213 & 0.014 & 6.670 & -0.007 \\\\ 7.00 & 3.557 & 5490 & 8.83394 & 0.013 & 6.087 & -0.010 \\\\ 6.80 & 3.515 & 5485 & 8.31328 & 0.011 & 5.733 & -0.014 \\\\ 6.45 & 3.438 & 5545 & 7.12090 & 0.011 & 4.921 & -0.007 \\\\ 6.10 & 3.357 & 5580 & 6.16131 & 0.009 & 4.267 & -0.006 \\\\ 6.00 & 3.333 & 5590 & 5.59661 & 0.006 & 3.884 & -0.007 \\\\ \\hline \\multicolumn{7}{c}{ML Relation from \\citet{chi89}} \\\\ \\hline 10.0 & 4.534 & 5150 & 66.9094 & 0.144 & 38.47 & -0.413 \\\\ 8.50 & 4.279 & 5100 & 44.8456 & 0.129 & 27.38 & -0.238 \\\\ 8.00 & 4.184 & 5090 & 38.5673 & 0.116 & 23.98 & -0.196 \\\\ 7.40 & 4.061 & 5050 & 32.4902 & 0.096 & 20.55 & -0.171 \\\\ 7.20 & 4.019 & 5250 & 26.1109 & 0.081 & 16.99 & -0.115 \\\\ 7.00 & 3.975 & 5400 & 21.8715 & 0.060 & 14.45 & -0.077 \\\\ 6.30 & 3.810 & 5390 & 16.9148 & 0.055 & 11.37 & -0.048 \\\\ 5.50 & 3.597 & 5350 & 12.4229 & 0.038 & 8.445 & -0.051 \\\\ 5.10 & 3.478 & 5396 & 10.0049 & 0.030 & 6.851 & -0.041 \\\\ 5.00 & 3.447 & 5420 & 9.37928 & 0.028 & 6.436 & -0.036 \\\\ 4.90 & 3.416 & 5440 & 8.80955 & 0.026 & 6.056 & -0.033 \\\\ 4.80 & 3.383 & 5470 & 8.21323 & 0.025 & 5.657 & -0.027 \\\\ 4.50 & 3.282 & 5460 & 7.06755 & 0.017 & 4.878 & -0.037 \\\\ 4.40 & 3.247 & 5490 & 6.56235 & 0.016 & 4.536 & -0.032 \\\\ 4.20 & 3.174 & 5560 & 5.60322 & 0.015 & 3.887 & -0.022 \\\\ \\hline \\multicolumn{7}{c}{Input parameters from SKM} \\\\ \\hline 4.57 & 3.306 & 5707 & 10.5472 & 0.035 & 7.234 & -0.000 \\\\ 5.44 & 3.578 & 5550 & 6.25949 & 0.022 & 4.340 &  0.020 \\\\ \\hline \\end{tabular} \\end{table}    \\begin{table} \\centering \\caption{Temperatures at maximum and minimum light from full-amplitude non-linear model calculations. The periods, luminosity and temperature are in days, $L_{\\odot}$ and $K$, respectively. \\label{tab2}} \\begin{tabular}{ccccc} \\hline $P$ & $L_{max}$ & $T_{max}$ & $L_{min}$ & $T_{min}$ \\\\ \\hline \\hline \\multicolumn{5}{c}{ML Relation from \\citet{bon00}} \\\\ \\hline 43.3924    & 23323.69 & 5301.44 & 15265.45 & 4653.25 \\\\ 37.8090    & 20547.96 & 5355.92 & 13844.78 & 4759.77 \\\\ 34.4156    & 18675.64 & 5399.02 & 12811.00 & 4681.60 \\\\ 30.8189    & 17046.03 & 5446.49 & 11592.02 & 4719.65 \\\\ 23.4835    & 12259.96 & 5456.64 & 8535.450 & 4820.82 \\\\ 20.0414    & 12503.63 & 5720.76 & 7897.603 & 4948.82 \\\\ 15.8313    & 8828.696 & 5664.44 & 6018.911 & 5025.59 \\\\ 12.1076    & 5616.756 & 5525.95 & 4402.925 & 5135.11 \\\\ 10.0124    & 4470.551 & 5531.94 & 3703.182 & 5274.68 \\\\ 9.69213    & 4484.612 & 5585.30 & 3647.089 & 5303.57 \\\\ 8.83394    & 3872.188 & 5619.60 & 3172.488 & 5312.02 \\\\ 8.31328    & 3501.343 & 5612.30 & 2906.126 & 5322.39 \\\\ 7.12090    & 2954.314 & 5673.40 & 2436.595 & 5396.37 \\\\ 6.16131    & 2434.291 & 5700.83 & 2059.419 & 5459.73 \\\\ 5.59661    & 2260.885 & 5691.47 & 2023.329 & 5513.33 \\\\ \\hline \\multicolumn{5}{c}{ML Relation from \\citet{chi89}} \\\\ \\hline 66.9094    & 39933.45 & 5468.07 & 25451.85 & 4807.32 \\\\ 44.8456    & 24179.71 & 5738.10 & 12478.80 & 4707.13 \\\\ 38.5673    & 19927.26 & 5762.38 & 9964.144 & 4696.57 \\\\ 32.4902    & 15175.07 & 5724.16 & 7818.562 & 4720.46 \\\\ 26.1109    & 13968.71 & 5886.81 & 7294.068 & 4949.57 \\\\ 21.8715    & 13020.35 & 6056.17 & 7076.027 & 5176.93 \\\\ 16.9148    & 8608.982 & 5999.40 & 4782.056 & 4961.41 \\\\ 12.4229    & 5070.216 & 5846.81 & 3081.266 & 4998.02 \\\\ 10.0049    & 3724.848 & 5785.06 & 2465.684 & 5074.73 \\\\ 9.37928    & 3427.493 & 5770.41 & 2320.390 & 5115.94 \\\\ 8.80955    & 3130.614 & 5748.80 & 2187.659 & 5156.00 \\\\ 8.21323    & 2843.303 & 5729.82 & 2062.474 & 5217.61 \\\\ 7.06755    & 2080.732 & 5572.38 & 1722.053 & 5295.72 \\\\ 6.56235    & 1885.999 & 5565.48 & 1587.056 & 5332.63 \\\\ 5.60322    & 1580.301 & 5697.21 & 1315.624 & 5383.56 \\\\ \\hline \\multicolumn{5}{c}{Input parameters from SKM} \\\\ \\hline 10.5472    & 2276.417 & 5894.64 & 1660.255 & 5430.83 \\\\ 6.25949    & 5084.652 & 6068.32 & 2974.031 & 5170.89 \\\\ \\hline \\end{tabular} \\end{table} ",
        "conclusions": "By looking at the way the Cepheid photosphere, the region where the Cepheid continuum is generated, interacts with the HIF, we have provided a simple qualitative physical explanation for the observed PC properties of fundamental mode Galactic Cepheids with $\\log (P) > 0.8$. This explanation relies on the fact that the opacity in a Cepheid becomes very high when hydrogen starts to ionize. This acts as a wall and prevents the photosphere going any deeper, and leads to a flat relation between period and temperature at the phase when the HIF interacts with the photosphere. For Galactic Cepheids this is observed at maximum light. At other phases, the photosphere is located away from the opacity wall and its temperature is related to the global properties of the star and hence its period. This explains, convincingly, why the Galactic Cepheid period-temperature relation is flat at maximum light and has a non-zero, single slope at mean and minimum light.  The qualitative nature of this idea remains true whether the pulsation code is purely radiative (as in SKM), or has a numerical recipe to model time dependent turbulent convection as in this work. The interaction between the photosphere and HIF may also provide some explanation for the suggestion made by \\citet{ker04}, that the region where spectral lines are formed do not necessary move homologously with the region where the K-band continuum is formed. In addition, because we have used two ML relations with a wide range of $L/M$ ratio in our study, and because the same physical effect is present in models constructed with either relation, the qualitative nature of our result is independent of the ML relation used. However, the ML relation used and its variation with metallicity, will dictate how the interaction of the HIF with the photosphere changes with Cepheids in a different metallicity environment. In the next paper in this series, we will investigate how these ideas can be used to explain the PC, AC and PL properties of fundamental mode LMC Cepheids.  We have also found evidence that the non-linear nature of the Galactic PC relation at maximum light reported by KN is due to short period fundamental mode Cepheids with $\\log (P) < 0.8$. These short period stars follow distinctly different PC relations and deserve further detailed study.  \\subsection{The effect on the PL relation}  What bearing do the results of this paper have on the Cepheid PC and PL relations at mean light? It is clear that PC relations at different phases contribute to the PC relation at mean light. If we choose $(V-I)_{i}$ as the extinction corrected colour at phase $i$, then a PC relation at this phase is:  \\[ (V-I)_{i} = a_{i} + b_{i}\\log (P), \\]  \\ni and taking the average over a pulsation period, $i=1...N$, this becomes:  \\begin{eqnarray} <V> - <I> = {1\\over{N}}\\sum_{i=1}^{i=N}a_{i} + {1\\over{N}}\\sum_{i=1}^{i=N}b_{i} \\log (P). \\end{eqnarray}  \\ni It is clear that the average intercept and slope will be affected by their values at individual phase points such as maximum or minimum light. \\citet{nge04b} compute PC relations for Galactic and Magellanic Cloud Cepheids at all phases between 0 and 1 and show that, for example for Galactic Cepheids, as the phase approaches maximum light, the slope of the PC relation becomes flatter. Since the mean light PC relation affects the mean light PL relation \\citep{mad91}, our aim of understanding the reasons behind changes in the PC and AC relations at different phases has a direct bearing on understanding at least one cause of the possible variation of the mean light Cepheid PL relation from galaxy to galaxy. Below we outline how our results are pertinent to studies of the variability of the Cepheid PL relation from galaxy to galaxy, though we emphasize that much of this discussion is dependent on theoretical model and data analysis currently being undertaken. Hence the following discussion is a road-map of some aspects of our future work.  The evidence that shows the slope of the LMC PL relation at mean light is significantly different to the slope of the mean light PL relation in the Galaxy has been provided in \\citet{tam03}, \\citet{nge04} and \\citet{sto04}. Furthermore, the mean light LMC PL relation suffers a change at a period of 10 days whilst the mean light Galactic PL relation data is consistent with a single slope with current data \\citep{tam03,fou03,kan04,nge04,san04,sto04}.  To understand the effect of the PC relation on the PL relation, consider the period-luminosity-color (PLC) relation valid at any phase:  \\[ M_v = a + b\\log (P) + c(V-I). \\]  \\ni If the PC relation is broken, say at a period $P_0$, then we have:  \\[ (V-I) = x + y\\log (P), P<P_0. \\] \\[ (V-I)=x'+y'\\log (P), P>P_0. \\]  \\ni  Substituting these two equations into the PLC relation leads to two PL relations:  \\[ M_v = a+cx + (b+cy)\\log (P),\\ P<P_0. \\] \\[ M_v = a+cx' + (b+cy')\\log (P),\\ P>P_0. \\]  \\ni  Hence changes in the slope and intercept of the PC relation have a direct bearing on the slope and intercept of the PL relation. However, the results from this paper and also from KN and \\citet{nge04b} imply that the mean light PC relation is affected by changes in the PC relation at different phases during a pulsation cycle, as given in equation (5). Thus the study of PC relations at various phases impacts on the variability of the mean light PL relation from galaxy to galaxy.  In this paper, we have updated and extended the work of SKM to provide an account of a simple physical mechanism, the interaction of the photosphere and HIF, which can change the properties of PC relations for Galactic Cepheids and so affect the PL relation. Specifically, a new result of our work is that for Galactic Cepheids with $\\log (P) > 0.8$, the maximum light PC relation is flat, i.e., the HIF and photosphere are engaged at maximum light. We also study the changes in AC relations because they are linked to PC relations through equation (1). A flat PC relation at maximum light leads to an AC relation at minimum light of slope $\\approx -0.1$ from equation (1), as seen in Figure \\ref{fig3} and Table \\ref{tab5}.  However, a Cepheid in the LMC will have a different ML relation, usually in the sense that LMC ML relations have higher $L/M$ ratios. Then in order to compare Galactic and LMC Cepheids of the same period, the LMC Cepheid needs to be hotter (KN). \\citet{kan95} and \\citet{kan96} found that such changes in the $(M,L,T_e)$ triplet describing the model changes the relative location of the HIF and hence the phase and the range of periods at which they interact. Our hypothesis is that for LMC models, it is only after $\\log (P) > 1.0$ that the HIF and photosphere are engaged at phases around maximum light, leading to a flat PC relation. This flat PC relation for long period ($\\log (P) > 1.0$) LMC Cepheids is one cause for the non-linear nature of the mean light LMC PL relation (KN). Empirically, \\citet{nge04b} provide preliminary evidence that the LMC PC relation is flat only for $\\log (P) > 1.0$ whereas this crossover period is $\\log (P) \\approx 0.8$ in the case of the Galaxy. In case of SMC Cepheids, KN also provide evidence that the maximum light PC relation is not flat even for long period Cepheids. Our contention is, using some of the results of this paper, that amplitudes in SMC Cepheids are not high enough to force an interaction between the photosphere and HIF, as the SMC Cepheids have lower amplitudes than the LMC \\citep{pac00}. This again must wait confirmation from a new set of SMC models.  The final effect on PL relations at mean light will depend on the behavior at other phases, such as minimum light. For example, preliminary calculations to be presented in a future paper imply that certain changes seen in the PC relation at minimum light (KN) may correspond to the first overtone mode becoming stable. Not surprisingly, a thorough quantitative study of this must await the analysis of LMC and SMC Cepheid models, which are currently under construction. The details of the influence of flat/non-flat PC relations on the mean PC and PL relations, the crossover period of $P_0$ that could vary from galaxy to galaxy, and the comparison of PL relations at maximum and minimum light for different galaxies are beyond the scope of this paper, and will be presented in the future papers.  \\citet{gro04} recently reported a metallicity dependence in the zero point of Cepheid PL relation at mean light. They used 34 Galactic Cepheids with individual metallicity measurements and then supplemented this sample with primarily long period Magellanic Cloud Cepheids to show the existence of a quadratic term in $\\log (P)$ in the PL relation. When they used primarily Galactic Cepheids in their sample, they found no evidence of a quadratic term. They interpreted these results as being due to a metallicity dependent zero point in the PL relation. These results are also consistent with the Cepheid LMC PL relation having different slopes for long and short period Cepheids as suggested by KN and \\citet{san04}. \\citet{san04} also plotted amplitude-mean colour relations in luminosity bins. However, KN's AC relations were along the instability strip. It will be interesting to apply the precepts behind equation (1) and multi-phase AC relations in luminosity and/or period bins, that is, across the instability strip."
    },
    "0404/astro-ph0404028_arXiv.txt": {
        "abstract": "We present a catalogue of $R_{\\rm c}I_{\\rm c}Z$ photometry over an area of 0.855 square degrees, centred on the young open cluster NGC 2547. The survey is substantially complete to limits of $R_{c}=21.5$, $I_{\\rm c}=19.5$, $Z=19.5$. We use the catalogue to define a sample of NGC 2547 candidates with model-dependent masses of about 0.05-1.0\\,$M_{\\odot}$. After correcting for incompleteness and estimating contamination by foreground field dwarfs, we investigate the mass function of the cluster, its binary content, and search for evidence of mass segregation among the lower mass stars.  There is ample evidence for mass segregation between high ($>3\\,M_{\\odot}$) and lower mass stars, but over the range $0.1<M<0.7\\,M_{\\odot}$, the data are consistent with no further mass segregation. By fitting King profiles we conclude that at least 60 percent of the low-mass stellar population are contained within our survey. The cluster mass function is remarkably similar to the Pleiades for $0.075<M<0.7\\,M_{\\odot}$.  Because of its age ($\\simeq 30$\\,Myr), we demonstrate that this mass function is robust to a number of systematic uncertainties likely to affect older and younger clusters and is therefore one of the best available estimates for the initial mass function in young disc populations.  For $0.05<M<0.075\\,M_{\\odot}$ there is some evidence for a deficit of brown dwarfs in NGC 2547 compared with other clusters. This deficit may extend to lower masses or may only be a dip, perhaps caused by an imperfect understanding of the mass-magnitude relationship at temperatures of around 2800\\,K. Incompleteness in both our survey, and the luminosity functions from which we estimate contamination by foreground objects, leave this question open. The binary fraction for systems with mass ratios greater than about 0.5 is 20-35 per cent for M-dwarfs in NGC 2547, quite consistent with that found in the field and other young clusters.  The full photometric catalogue and our lists of candidate cluster members are made available in electronic format. ",
        "introduction": "\\label{intro}   The open cluster NGC 2547 (= C0809-491) is important for efforts to understand the formation and evolution of low-mass, pre-main-sequence (PMS) stars.  Its isochronal age is about 30\\,Myr, it is relatively close ($\\simeq$ 400--450\\,pc), has low reddening ($E(B-V)=0.06$) and a rich population of identified early and late-type members (see Clari\\'{a} 1982; Jeffries, Totten \\& James 2000; Naylor et al. 2002 -- hereafter N02). Being close and quite compact it serves as an ideal laboratory to investigate evolutionary tracks and isochrones, rotation rates, magnetic activity and light element abundances of low-mass PMS stars as they approach the ZAMS. In particular, NGC 2547 occupies an important and comparatively little studied niche of parameter space with an age between those of regions of recent star formation, and clusters like the Pleiades, where low-mass stars have already reached the zero-age main-sequence (ZAMS).  Jeffries \\& Tolley (1998) first identified the low-mass PMS population of NGC 2547 as counterparts to {\\em ROSAT} X-ray sources. A number of these counterparts were positively identified as F to K-type cluster members by Jeffries et al. (2000) using radial velocities. A $BVI_{\\rm c}$ photometric study of NGC 2547 was presented by N02. Their survey covered an approximately square area of 0.32 square degrees and identified several hundred candidate cluster members with masses from 6\\,$M_{\\odot}$ down to about 0.25$M_{\\odot}$ at $V\\simeq20.5$. N02 concluded that the mass function (MF) was similar in shape to that of the Pleiades and of field stars over this mass range, but that the total mass of NGC 2547 could be a few times less than the Pleiades. Littlefair et al. (2003) have used the same list of candidate members to show that stars with mass $>3\\,M_{\\odot}$ in NGC 2547 are much more centrally concentrated than lower mass objects. Because NGC 2547 is no more than 10 dynamical crossing times old, it seems likely that most of this segregation must be primordial (e.g. Bonnell et al. 2001), rather than due to dynamical evolution.  The N02 survey did not cover a sufficiently large area or go to sufficient depths to discover whether the spatial distribution of low-mass ($<1\\,M_{\\odot}$) cluster members was independent of mass. The possibility of differential mass segregation between lower mass stars is important. Some recent theories concerning the formation of very low-mass stars and brown dwarfs suggest they are ejected as low-mass fragments from protostellar multiple systems before they have a chance to accrete significant material (e.g. Reipurth \\& Clarke 2001). Such objects may have a greater velocity dispersion than their higher mass siblings and hence be less spatially concentrated in a cluster (Sterzik \\& Durisen 2003).  This in turn would mean that a cluster MF measured over a limited volume would understimate the contribution from the lowest mass stars and brown dwarfs. However, alternative simulations of brown dwarf formation (e.g. Bate, Bonnell \\& Bromm 2003) predict that the initial velocity dispersions of low mass stars and brown dwarfs will be similar and little mass segregation or preferential evaporation of very low-mass objects would be expected prior to full dynamical relaxation.  In this paper we present a new $R_{\\rm c}I_{\\rm c}Z$ photometric survey of NGC 2547. This survey is deeper than that of N02, extending to just below the substellar boundary, encompassing the ``turnover'' of the observed MF seen in the field and in the Pleiades at $\\sim 0.2\\,M_{\\odot}$ (see Chabrier 2003 and references therein). The new survey also covers a wider area (0.855 square degrees) allowing us to explore the question of mass segregation among lower mass cluster candidates in more detail. The observational data and analysis are presented in section~2. The selection of candidate cluster members and an evaluation of completeness and contamination are addressed in section~3. In section 4 we evaluate the mass segregation, luminosity and mass functions for NGC 2547. We discuss the results and draw our conclusions in sections 5 and 6. ",
        "conclusions": "In this paper we have presented optical/near infrared photometry of NGC 2547 that is both deeper and covers a wider area (0.855 square degrees) than previously published surveys. The survey has been used to collate a catalogue of candidate members and investigate the degree of mass segregation and the mass function (MF) for the cluster.  The main conclusions of our work can be summarised as follows:  \\begin{enumerate}  \\item The cluster shows strong evidence for mass segregation, in that stars with $M>3\\,M_{\\odot}$ are {\\it much} more centrally concentrated than lower mass stars. There is some evidence for mass segregation above and below 1\\,$M_{\\odot}$, but lower mass stars have spatial distributions that are consistent with no further mass segregation for $0.1<M<0.7\\,M_{\\odot}$. By fitting King profiles we have estimated that at least 60 per cent of the cluster's low-mass population is included within our survey.  \\item The MF for $0.075<M<0.7\\,M_{\\odot}$ in NGC 2547 has been determined with roughly the same statistical precision as available for the best Pleiades studies (and much more precisely than field MF determinations) and is in remarkable agreement with the Pleiades MF (and other young clusters) over this mass range.  Because of it's age ($\\simeq 30$\\,Myr), NGC 2547 is not badly affected by a number of systematic uncertainties which could well affect studies of both younger and older clusters (age spreads and uncertainties, reddening, discs, very cool atmospheres, dynamical evolution) and neither are our results very dependent on which stellar evolutionary models are used. Our analysis therefore represents one of the most precise and robust determinations of a low-mass disc population MF, and possibly the initial MF if primordial velocity dispersions are reasonably mass-independent.  \\item Below the hydrogen burning limit we have identified an apparent dearth of high-mass brown dwarf candidates in NGC 2547 ($0.05<M<0.075\\,M_{\\odot}$). This deficit {\\it may} extend to lower masses, but increasing incompleteness in both our survey and our estimates of contaminating foreground M-dwarfs leave this question open. Alternatively, we may be seeing an example of the ``missing M dwarf'' phenomenon, identified by Dobbie et al. (2002), which could result from an imperfect understanding of atmospheres with temperatures of around 2700\\,K and would not reflect a true deficit of brown dwarfs.  A deeper, wider survey of NGC 2547 is called for to verify this result, to provide a better empirical estimate of contamination among the lowest mass cluster candidates, and hence show whether the deficit we have seen is merely a dip or a genuine lack of brown dwarfs.  \\item The total mass of NGC 2547 for stars with $M>0.06\\,M_{\\odot}$ is $(450\\pm 100)\\,M_{\\odot}$ and about a factor of two smaller than the Pleiades.  \\item The binary fraction of M-dwarfs in NGC 2547 is between 20 and 35 per cent for systems with mass ratios greater than 0.35 to 0.65. This fraction is consistent with values determined for populations of low-mass stars found in the field and other young clusters.  \\item Finally, we have provided in electronic format both our entire photometric catalogue as well as subsets of photometrically selected cluster candidates.  These catalogues contain data with robust and precisely determined photometric and astrometric uncertainties.  \\end{enumerate}"
    },
    "0404/astro-ph0404502_arXiv.txt": {
        "abstract": "We study the hydrodynamical evolution of galactic winds in disky dwarf galaxies moving through an intergalactic medium. In agreement with previous investigations, we find that when the ram pressure stripping does not disrupt the ISM, it usually has a negligible effect on the galactic wind dynamics. Only when the IGM ram pressure is comparable to the central ISM thermal pressure the stripping and the superwind influence each other increasing the gas removal rate. In this case several parameters regulate the ISM ejection process, as the original distribution of the ISM and the geometry of the IGM-galaxy interaction. When the ISM is not removed by the ram pressure or the wind, it loses memory of the starburst episode and recovers almost its pre-burst distribution in a timescale of 50-200 Myr. After this time another star formation episode becomes, in principle, possible. Evidently, galactic winds are consistent with a recurrent bursts star formation history.  Contrary to the ISM content, the amount of the metal-rich ejecta retained by the galaxy is more sensitive to the ram pressure action. Part of the ejecta is first trapped in a low density, extraplanar gas produced by the IGM-ISM interaction, and then pushed back onto the galactic disc. The amount of trapped metals in a moving galaxy may be up to three times larger than in a galaxy at rest. This prediction may be tested comparing metallicity of dwarf galaxies in nearby poor clusters or groups, such as Virgo or Fornax, with the field counterpart. The sensitivity of the metal entrapment efficiency on the geometry of the interaction may explain part of the observed scatter in the metallicity-luminosity relation for dwarf galaxies. ",
        "introduction": "Dwarf galaxies are key players in theories of galaxy formation.  In the standard cold dark matter picture they are the first forming objects, and larger galaxies are successively built by merging of these small systems (Blumenthal et al. 1984). Given their very low metallicity and small size, these poorly evolved objects are excellent laboratories to investigate the feedback of starbursts on the interstellar medium (ISM) and to study their chemical evolution. In models of dwarf galaxies formation the feedback from supernovae (SNe) and the consequent gas and metals loss is a crucial process (Dekel \\& Silk 1986, Dekel \\& Woo 2003).  The impact of starbursts in local dwarf galaxies is well studied observationally (Martin 1998, 1999; see Heckman 2003 for a recent review), but important theoretical questions remain unanswered.  A critical open problem is given by the dwarf galaxies chemical evolution. Chemical evolution models of blue compact galaxies (BCGs) (Matteucci \\& Tosi, 1985; Pilyugin 1992; Marconi, Matteucci, \\& Tosi 1994; Bradamante, Matteucci, \\& D'Ercole 1998; Larsen, Sommer-Larsen \\& Pagel 2001) indicate that the gas fraction-metallicity relations of BCGs is not compatible with the closed box scenario, suggesting that (differential) galactic winds carry away a large fraction of the metals produced by the young stars.  Numerical simulations (e.g. MacLow \\& Ferrara 1999, D'Ercole \\& Brighenti 1999, Strickland \\& Stevens 2000, Recchi et al. 2001, 2002) and analytic models (De Young \\& Heckman 1994) indeed show that metals are easily ejected in the intergalactic medium (IGM).  However, the details of the interaction between the metal rich hot gas and the cold interstellar medium or the IGM are yet not well understood. For instance, the mixing timescale for the metals produced in the starburst with the ISM is poorly known. Chemical evolution models often assume instantaneous mixing but HII regions nearby star formation sites do not appear to be enriched by the young stars which illuminate them (Kobulnicky \\& Skillman 1996, 1997, 1998).  This suggests that most of the freshly synthesized heavy elements do not mix right away with the surrounding ISM and instead reside for timescales $> 10^7$ yr in the hot ($\\ge 10^6$ K) phase, where they are indeed observed (Martin, Kobulnicky \\& Heckman 2002).  The question is then whether the hot gas leaves the galaxy and enriches the IGM or whether it eventually cools and mixes with the ISM on long timescales. The low but non-negligible metallicity of dwarf galaxies suggests that some form of mixing is effective. Possible mixing mechanisms are molecular diffusion (Roy \\& Kunth 1995, Tenorio-Tagle 1996, Oey 2003), condensation through thermal conduction (McKee \\& Begelman 1990) or turbulent mixing, likely the most important one (Bateman \\& Larson 1993, Roy \\& Kunth 1995). The close correlation between X-ray and $H\\alpha$ emission also indicates some degree of thermal mixing between hot and warm gas (e.g. Lehnert, Heckman \\& Weaver 1999, Strickland et al. 2002).  Babul \\& Rees (1992) and Tenorio-Tagle (1996) proposed a scenario in which the metal enriched superbubble powered by the wind is confined by a relatively high-pressure medium and then pushed back into the galaxy. Babul \\& Rees (1992) suggest that the confinment is exerted by the IGM (ICM) of the group (cluster) to which the galaxy belongs. This scenario was successively investigated in more detail by Murakami \\& Babul (1999), who estimate that the time-scale for the superbubble to collapse is of the order of a few $10^7$ yr. In the model proposed by Tenorio-Tagle (1996, see also Silich \\& Tenorio-Tagle 2001), instead, the superbubble is halted by a hypothesized gaseous halo surrounding the galaxy. The hot gas in the bubble then cools and falls back onto the galaxy in a time-scale $\\sim 1$ Gyr, much longer than in the models by Murakami \\& Babul (1999). While in their simple form these scenarios are probably too effective in retaining heavy elements (the low metal content of dwarf galaxies implies significant metal loss), some ``weaker'' version, where only a small but non-negligible fraction of the metal rich gas is able to cool, might help explaining the observed metallicity of dwarf galaxies.  A further complication is given by of ram pressure effect if, as expected, the galaxy moves through the IGM. It is difficult to  predict the final effect of the interaction of a galactic wind with  the IGM. The ram pressure  could be synergic to the SN heating in removing the ISM from the  galaxy. Gas lifted above the galactic plane by the SN energy may  be dragged away by the IGM flow. Ram pressure effects were considered by Murakami \\& Babul (1999) in a number of 2D numerical simulations. These authors find that when the ram pressure of the IGM is larger than its thermal pressure, the expanding shell driven by the galactic wind is deformed, fragmenting into dense clouds eventually dragged away from the galaxy. Analogously, in the model by Tenorio-Tagle (1996) the ram pressure may strip away the hypothesized extended, loosely bound ISM, removing the medium confining the superbubble and the metals contained in it.  On the other hand, may also happen that, depending on the inclination angle between the galactic plane and the orbital plane, at least one lobe of the expanding superbubble can be squashed back on the galaxy by the ram pressure.  In this case the enrichment process would be more efficient compared to the case where ram pressure stripping is negligible.  Motivated by the above arguments, we have studied the interaction of starbursting dwarfs with the surrounding IGM, and how such interaction influences the evolution of the superwind powered by the stellar burst. In a previous paper (Marcolini, Brighenti \\& D'Ercole 2003, hereafter Paper I) we investigated the effect of the ram pressure alone on the ISM of disky dwarf galaxies. Contrary to most of the other papers devoted to this argument, we considered a ram pressure typical of galaxy groups rather than of rich clusters, because most dwarf galaxies are found in the environment of loose groups (Tully 1987). In the present paper we follow for a long timescale (500 Myr) the evolution of a galactic wind originating at the galactic centre. We focus in particular on the ejection of the ISM and of the metals synthesised in the starburst, and how their circulation is influenced by the environment and by the ram pressure. ",
        "conclusions": "Here we briefly summarise and discuss the behaviour of the ISM and of the SN ejecta in our models.  For a better understanding of our results we also simulated the evolution of galactic winds occurring in the usual model galaxies but now assumed at rest relative to the IGM ($\\rho_{\\rm IGM}=2\\times 10^{-28}$ g cm$^{-3}$, $T_{\\rm IGM} = 10^6$ K; REST models), or not surrounded by any IGM (FIELD models). The aim of these models is to obtain a more direct insight of the role played by the ram pressure comparing interesting quantities such as the ISM and ejecta mass content (see Figures 5 and 7) in otherwise identical galaxies. The efficiency of metal ejection is known to be sensitive to the details of numerical simulations (cf. D'Ercole \\& Brighenti 1999 and MacLow \\& Ferrara 1999), and a consistent comparison must be done among similar models.  \\subsection{ISM evolution}  The ISM of SM galaxies subject to the high ram pressure is quickly dragged away and therefore we did not simulate winds for SM-HI models. The low ram pressure strips only $30-50$ \\% of the original ISM of small galaxies (SM-LO models). However, the starburst powered wind completely removes all the gas in $\\approx 100$ Myr, a result obtained also by the galaxy at rest. Evidently the ram pressure has a negligible influence on the galactic wind evolution for these models. For more massive galaxies the occurrence of a starburst does not influence significantly the ongoing mass loss due to the ram pressure in the MD-LO models and can only anticipate the complete removal of gas (e.g. for model MD-00-HI).  For the more massive LG galaxies the situation is complicated by the radiative nature of the gaseous halo which develops around the galactic disk, especially in the edge-on model with the larger ram pressure (LG-00-HI). Cooling of the ablated gas makes $M_{\\rm centr}$ to increase by a factor of few; $M_{\\rm gal}$, instead, is not affected. As for the other models, the occurrence of a starburst in models LG does not influence the time evolution of $M_{\\rm gal}$, which is determined by the ram pressure in both the HI and LO cases. We stress that in these models, as well as in any other model in which the galaxy is not rapidly deprived of gas by the wind, the ISM loses memory of the starburst after a few tens of Myr and recovers a distribution similar to the initial one. At this point the galaxy is in principle ready for another possible starburst episode.  \\begin{figure*} \\begin{center} \\psfig{figure=figure9col.ps} \\end{center} \\caption{Evolution of ${\\cal Z}$ (see text) in the central region. The left and right panel refer to MD and LG models, respectively. Solid blue lines: edge-on models; dashed-dotted green lines: $45^{\\circ}$ models; dashed red lines: face-on models.  Higher and lower dotted lines in the right panel represent the FIELD and REST models, respectively.  Light and heavy lines refer to LO and HI ram pressure models, respectively.} \\end{figure*}  In conclusion, our models show that the galactic wind either disrupts the whole ISM or has a negligible effect on the ISM content, which in turn is regulated mainly by the ram pressure. When $\\rho_{\\rm IGM} v^2_{\\rm IGM} \\sim P_0$, where $P_0$ is the central ISM thermal pressure, the ram pressure stripping and the wind act together to increase the ISM removal rate. Many parameters regulate the ISM dynamics in this case (intensity and duration of the ram pressure, wind mechanical luminosity, potential well depth, initial amount of the ISM, value of the inclination angle $\\theta$) and numerical simulations are needed to understand the ISM behaviour of a specific model.  \\subsection{Ejecta evolution}  For our low mass galaxies (models SM) the wind expels the whole ISM and no ejecta remains trapped into them. The more massive models with the lower ram pressure (MD-LO models) retain 5\\% of the total ejecta mass in the central region, and 10\\% in the whole galaxy. These are essentially the same values obtained by the analogous REST model, while for the analogous FIELD model the trapped quantities are lower by a factor of $\\sim 2$. We thus conclude that for these models the ram pressure has little effect in the entrapment of the SN metals, while the presence of a relatively high pressure ambient gas helps this task. These results are in qualitative agreement with those in Murakami \\& Babul (1999). A quantitative comparison with the models by Murakami \\& Babul however is not possible. In fact, having these authors considered a spherical galaxy, the (metal rich) superbubble preserves a spherical shape when the ram pressure is negligible, and collapses as a whole. Thus, all the metals produced by the starburst are retained by the galaxy. On the other hand, in our flattened galaxies a large fraction of the ejecta is always lost through the superbubble breakout.  With the larger ram pressure, the ejecta and the ISM in the edge-on model MD-00-HI (the only model which keeps some ISM at the end of the ram pressure phase) are stripped after $t\\sim 200$ Myr (Figure 8). For $t < 100$ Myr, however, the time evolution of the ejecta masses in both the central and galactic region is remarkably similar to that of model MD-00-LO. If the starburst occurs when the values of $M_{\\rm centr}$ is high (model MDbis) more ejecta gets trapped, peaking at $t \\simeq 30$ Myr with values of 20\\% and 55\\% in the central and galactic region, respectively. Then it decreases to 5\\% and 10\\% in the two regions.  The deeper potential well and the larger ISM amount of the LG-LO models are not the direct cause of the larger retention of metals with respect to the MD-LO galaxies, as can be verified comparing the FIELD and REST models for the LG and MD galaxies. Thus, the larger fraction of trapped metals occurring in the moving LG models is due to the presence of the extraplanar gaseous halo which develops in these models (in the edge-on case in particular), especially with the larger ram pressure. For the weaker ram pressure we find that the edge-on model keeps 20\\% of the ejecta in the central region at the end, though goes through a minima at $t$=70 Myr with 5\\%. The rise of $M_{\\rm ej,centr}$ after this time is due to the collapsing halo gas polluted by the metals carried by the wind. The halo is also responsible of the rather high value ($\\sim 70$ \\%) of the mass of the ejecta trapped in the whole galactic region of the LG-00-LO model. For different values of $\\theta$, the halo is less developed, and less metals are trapped: around 10\\% in the central region, and $\\sim 40$ \\% in the galactic region, with little differences for various values of $\\theta$.  For LG-HI models the halo is even more influential. Now the ejecta trapped in the central region is substantial: $\\sim 43$\\% for $\\theta=0^{\\circ}$,  $\\sim 30$\\% for $\\theta=45^{\\circ}$, and $\\sim 14$\\% for $\\theta=90^{\\circ}$. Very little ejecta is exterior to the central region, so the values of $M_{\\rm ej,gal}$ are very similar to those of $M_{\\rm ej,centr}$. Despite the presence of the more massive gaseous halo, $M_{\\rm ej,gal}$ in the LG-HI models becomes somewhat lower than in the LG-LO models because of the stronger ram pressure which continuously erodes the polluted ISM.  We conclude that, contrary to the ISM dynamics, the amount of the SN ejecta trapped into the galaxy results to be more affected by the action of the ram pressure. Part of the ejecta expelled by the superwind is pushed back onto the galaxy by the incoming IGM or remains trapped in the surrounding halo, and the fraction of metals retained by a moving galaxy can be up to 3 times (depending on the value of $\\theta$) larger than that retained by a galaxy at rest. This trend is opposite to that found by Murakami \\& Babul (1999) in their suite of models, where they also investigated the effect of group-like ram pressure on galactic winds. This discrepancy is due to the assumed spherical ISM distribution of their models which allows the complete retention of the metals in the simulations without ram pressure.  Overall, the main conclusion by De Young \\& Heckman (1994), D'Ercole \\& Brighenti (1999) and Mac Low \\& Ferrara (1999) are not changed by the ram pressure stripping: for galaxies of size comparable to SM and MD only a very small fraction on the metals (of the order of $\\approx 10$ \\%) remains trapped in the ISM. Only the large models LG can trap much more ejecta, an effect of the relatively large amount of gas located at high $z$. This is conveyed there by the ram pressure and depends on the flare in the ISM distribution (we discuss this point in section 5.4). The (substantial) metal mass not incorporated in the ISM at $t=500$ Myr is definitively lost by the galaxies and enriches the IGM, on spatial scales on the order of $30-50$ kpc. The general trend shown by our models is that the ram pressure increases $M_{\\rm ej,centr}$ and $M_{\\rm ej,gal}$.  To characterise the pollution degree we calculate the average ejecta fraction in the central region defined as ${\\cal Z}=M_{\\rm ej,centr}^{\\rm cold}/\\ M_{\\rm centr}$, where we consider only the cold fraction ($T<10^5$ K) of the ejecta which is effectively trapped in the ISM. In order to obtain the abundance of a specific X-element, one has to scale the ${\\cal Z}$ value by a factor $Z_{\\rm X,ej}$ representing the abundance of the X-element in the SN ejecta. Focussing on iron and oxigen, we assume $Z_{\\rm Fe,ej}=4.4 \\times 10^{-3}$ and $Z_{\\rm O,ej}=4.4 \\times 10^{-2}$, and obtain $Z_{\\rm Fe}/Z_{\\rm Fe,\\odot}=3.4 {\\cal Z}$ and $Z_{\\rm O}/Z_{\\rm O,\\odot}=4.6 {\\cal Z}$, respectively (cf. D'Ercole \\& Brighenti 1999 for more details).  In Fig. 9 we show the evolution of $\\cal Z$ in our models, and we note that in general larger ram pressures lead to an higher metal enrichment. Thus, the ability to retain metals appears to be sensitive to the parameters regulating the interaction ISM-IGM, and this may explain part of the observed scatter in the metallicity-luminosity relation (e.g. Lee, McCall \\& Richer 2003). We may also expect that dwarf irregulars in relatively high ram pressure environments have systematically larger metallicity than the field counterpart. Marginal evidence for such a trend has been claimed by Vilchez (1995) in his study of Virgo Irregulars. Elevated oxygen abundances for Virgo dwarfs has ben suggested also by Vilchez \\& Iglesias-Paramo (2003).    \\subsection{The gaseous halo}  As discussed above the gas which accumulates at large $z$ is a result of the ram pressure of the IGM with a flared ISM distribution. Its presence allows us to compare our results to those by Silich \\& Tenorio-Tagle (2001) and Legrand et al. (2001). These authors make a systematic study on the critical superwind luminosity needed for the superbubble to break out. They conclude that star-forming dwarf galaxies must have an extended gaseous halo in order retain their metals and enhance their abundances. A direct comparison between our edge-on models and the models by Silich \\& Tenorio-Tagle is prevented by the fact the galactic models are build following different criteria; the shape of the gaseous halo is also different. However, we consider two models by Silich \\& Tenorio-Tagle, their models M800.100 and M900.100, which have values of $M_{\\rm g}$ and $M_{\\rm h}$ similar to our MD and LG models, respectively. Moreover, the column density of the gaseous halo of these models during the edge-on stripping is also similar to the column density of the halo of the two quoted models by Silich \\& Tenorio-Tagle. Following these latter authors, the wind luminosity of our models is much larger than the critical luminosity needed by the superbubble to break out in the MD models. Thus the wind easily moves far from the galaxy bringing away most of the SNII ejecta. On the contrary, the wind luminosity would be only marginally sufficient to allow the breakout of the superbubble in the LG model. Actually, our simulations show that the superbubble breaks only marginally in this case, and the larger fraction of ejecta trapped in the galaxy (Fig. 7-8) indicates a more substantial role of the extraplanar gaseous halo.  \\subsection{Limitations of the models and future work}  In evaluating the above results one must bear in mind the two caveats discussed above: $i$) the numerical diffusion, and $ii$) the flared ISM distribution. The numerical diffusion prevents the possibility to give an accurate {\\it quantitative} estimate of the ejecta mixed with the ISM. However, we believe that the {\\it relative} differences among the models are still meaningful, and thus ram pressure may indeed increase the metal enrichment of the ISM due to a starburst.  The metal enrichment in our models is also influenced by the ISM flare. This influence is particularly evident for large galaxies moving edge-on through the IGM (LG-00 models), where a gaseous halo form around the galaxy, affecting the superwind expansion. The flare in our models derives from the assumption of isothermal ISM, which requires a rotation curve independent of $z$ in order to assure hydrostatic equilibrium (e.g. Tassoul 1978).  The real presence and the effective extension of flares in galaxies is still an open question, although evidence for flares in galaxies is claimed by several authors (e.g. Brinks \\& Burton 1984, Burton 1988, Olling 1996, Matthews \\& Wood 2003). Simulations analogous to those presented here but without flare in the initial ISM distribution would be interesting. We are currently devising non-isothermal models for the ISM to evaluate the real influence of the ISM distribution on the interaction with the IGM distribution."
    },
    "0404/astro-ph0404358_arXiv.txt": {
        "abstract": "{ Following Paper I, we provide extended tables of bolometric corrections, extinction coefficients, stellar isochrones, and integrated magnitudes and colours of single-burst stellar populations, for the Sloan Digital Sky Survey (SDSS) $ugriz$ photometric system. They are tested on comparisons with DR1 data for a few stellar systems, namely the Palomar~5 and NGC~2419 globular clusters and the Draco dSph galaxy. ",
        "introduction": "\\label{intro}  The Sloan Digital Sky Survey (SDSS; see York et al. 2000) is one of the most impressive astronomical campaigns ever carried out. It aims at providing photometry and subsequent spectroscopy for objects covering about one quarter of the entire sky. SDSS photometry is being obtained near-simultaneously in five broad-band filters in drift-scan mode, which results in highly homogeneous data.  The effective exposure time for imaging is approximately 54\\,s with a limiting magnitude of $r\\sim22.6$.  The primary science goals of the SDSS are at extragalactic and cosmological studies such as galaxy evolution as a function of redshift (e.g., G\\'omez et al.\\ 2003; Eisenstein et al.\\ 2003; Blanton et al.\\ 2001), the search for high-redshift quasars (e.g., Fan et al.\\ 2001, 2003),  weak lensing (e.g., Fischer et al.\\ 2000; McKay et al.\\ 2002), and the large-scale structure distribution of galaxies (e.g., Dodelson et al.\\ 2002; Zehavi et al.\\ 2002).  However, the SDSS also provides a huge database for various areas of stellar and Galactic astronomy, e.g., the search for rare or special stellar objects (e.g., Margon et al.\\ 2002; Hawley et al.\\ 2002; Helmi et al.\\ 2003), the study of resolved star clusters and nearby Milky Way satellites (e.g., Odenkirchen et al.\\ 2001a,b), and studies of Milky Way substructure (e.g., Chen et al. 2001; Newberg et al. 2002; Yanny et al. 2003).  For an overview of SDSS science, see Grebel (2001). Our paper aims at providing tools for the interpretation of SDSS stellar photometry data, i.e., isochrones transformed to the SDSS photometric system.  The SDSS began regular operations in 2000 and will run until 2005.  The SDSS data products include images in five passbands, photometrically and astrometrically calibrated object catalogs with a wealth of information on the properties of the recorded objects, and wavelength- and flux-calibrated spectra with redshifts. An Early Data Release (EDR) of SDSS commissioning data occurred in July 2001 (Stoughton et al.\\ 2002), and Data Release 1 (DR1) took place in May 2003 (DR1; Abazajian et al. 2003). The currently publicly available SDSS photometry data cover 2099 square degrees, about one fifth of the total anticipated survey area.  The subsequent data releases are planned to occur on a yearly basis.  The final one is scheduled for 2006.  It is estimated that the SDSS will ultimately comprise some $8\\cdot10^7$ stars with high-quality five-band photometry, including Galactic halo and disk field stars, stars in globular clusters, and in nearby resolved dwarf galaxies.   The SDSS has designed, defined, and calibrated its own photometric system (Fukugita et al.\\ 1996; Gunn et al.\\ 1998; Smith et al.\\ 2003), which is characterized by the following particularities: (1) the use of a modified Thuan-Gunn broadband filter system called $ugriz$; (2) a zero-point definition in the AB magnitude system; (3) the use of a modified, non-logarithmic, definition of the magnitude scale (Lupton, Gunn, \\& Szalay 1999). As a consequence, SDSS photometry cannot be easily transformed into traditional systems, or at least any transformation is likely to lead to a significant loss of the photometric information contained in SDSS data of fainter sources.  This presents a problem not only for users of the SDSS databases, but also for astronomers who obtain new data in the SDSS filters now offered at various observatories.  Therefore, in order to fully take advantage of the valuable and growing SDSS database and to enable the quantitative interpretation of data obtained in the SDSS filter system elsewhere, converting stellar models directly into the SDSS system is an obvious requirement. This is the goal of the present paper. This paper is a continuation of a series of papers dedicated to the conversion of theoretical stellar models to a wide variety of photometric systems. In Paper~I (Girardi et al.\\ 2002), we have described the assembly of a large library of stellar spectra, a simple formalism to compute tables of bolometric corrections from this, and the application of these corrections to a large database of theoretical isochrones. The systems there considered were the Johnson-Cousins-Glass, Washington, HST/WFPC2, HST/NICMOS, and ESO Imaging Survey ones (for the WFI, EMMI, and SOFI cameras in use at the European Southern Observatory at La Silla, Chile).  The basic procedures and the input stellar data involved in the present work are the same as in Paper~I. Thus, here we will skip most of the description, detailing just the particular aspects that are inherent to the SDSS photometric system.  In Section 2, we describe the SDSS AB and asinh magnitude systems and the synthetic photometry required to transform theoretical isochrones into these systems.  In Section 3, we apply these isochrones to multi-color DR1 data. Section 4 contains our conclusions. ",
        "conclusions": "\\label{sec_conclu}  We have computed bolometric corrections and extinction coefficients specific for the SDSS ABmag photometric system. They have been applied to transform previous Padova isochrones and integrated magnitudes of single-burst stellar populations into SDSS $ugriz$ absolute magnitudes. All data tables mentioned in this paper are available in the web site {\\tt http://pleiadi.pd.astro.it}. They have a structure similar to those already released with Paper~I, and extended descriptions are also included in the form of electronic files.  Comparisons between the present isochrone sets and SDSS DR1 data, together with the recent results by Rider et al. (2004), are encouraging. Our hope is that this database will be useful for the analysis of the huge amount of photometric data that has been, and will yet be, released by the SDSS project. Obvious applications go from the derivation of parameters for star clusters and nearby dwarf galaxies (distances, reddenings, ages and metallicities), to the isolation of particular objects in colour-colour diagrams, to the application in analyses of star counts and Galactic structure. The further use of SDSS filters in other observatories, in programs not related to the original survey, may well expand the possible range of use for these tables."
    },
    "0404/astro-ph0404444_arXiv.txt": {
        "abstract": "{\\it XMM-Newton} observations of seven QSOs are presented and the EPIC spectra analysed. Five of the AGN show evidence for Fe K$\\alpha$ emission, with three being slightly better fitted by lines of finite width; at the 99~per~cent level they are consistent with being intrinsically narrow, though. The broad-band spectra can be well modelled by a combination of different temperature blackbodies with a power-law, with temperatures between kT~$\\sim$~100--300~eV. On the whole, these temperatures are too high to be direct thermal emission from the accretion disc, so a  Comptonization model was used as a more physical parametrization. The Comptonizing electron population forms the soft excess emission, with an electron temperature of $\\sim$~120--680~eV. Power-law, thermal plasma and disc blackbody models were also fitted to the soft X-ray excess. Of the sample, four of the AGN are radio-quiet and three radio-loud. The radio-quiet QSOs may have slightly stronger soft excesses, although the electron temperatures cover the same range for both groups. ",
        "introduction": "\\label{sec:intro}  At energies below $\\sim$~2~keV, the spectra of most AGN show an upturn, away from the extrapolation of the high energy (2--10~keV) power-law. This so-called `soft excess' emission is thought to be common in both Seyfert galaxies and QSOs (Quasi-Stellar Objects). The first such soft excess was identified in Mrk~841 by Arnaud \\etal (1985); Turner \\& Pounds (1989) then found that $\\sim$~50~per~cent of their {\\it EXOSAT} sample showed steeper spectral slopes at low energies. Likewise, Walter \\& Fink (1993) and Schartel \\etal (1996) found that the {\\it ROSAT} PSPC spectral index tends to be significantly steeper than that measured above 2.4~keV (typically $\\sim$~1.9; e.g., Nandra \\& Pounds 1994). More recently, Pounds \\& Reeves (2002) have discussed the frequent presence of soft excesses in {\\it XMM-Newton} data.   Many papers have been published about the soft excess, covering both observational results -- with {\\it ROSAT} (e.g., Fiore \\etal 1994; Piro, Matt \\& Ricci 1997), {\\it Ginga}, {\\it EXOSAT} (e.g., Saxton \\etal 1993) and {\\it Einstein} (e.g., Masnou \\etal 1992; Zhou \\& Yu 1992) -- and theoretical work (e.g., Czerny \\& Elvis 1987; Czerny \\& {\\. Z}ycki 1994; Xia \\& Zhang 2001). It is generally thought that the soft excess may be linked to the hot tail-end of the Big Blue Bump (BBB), or is an extension of the UV band. The BBB/UV excess is likely to be due to thermal emission from the accretion disc surrounding the black hole (e.g., Shields 1978; Malkan \\& Sargent 1982). However, this thermal emission is not hot enough to account for the soft X-ray flux as well; hence, Comptonization is often invoked to explain the resultant emission. In this scenario, the direct thermal emission from the accretion disc is observed as the optical and UV spectrum. Some of the disc photons, however, undergo inverse Compton scattering with a population of hot electrons, thus gaining energy and producing a broader spectrum, which appears similar to a power-law over a limited energy range (assuming unsaturated Comptonization).  All seven objects in this paper, listed in Table~\\ref{z_wa}, have been previously observed by {\\it ROSAT} (Schartel \\etal 1996) and were each noted to have steep photon indices over the 0.1--2.4~keV {\\it ROSAT} band. The Galactic absorption in the direction of each low-redshift QSO is small and it has been previously found that there is no significant evidence for additional, intrinsic absorption in any of the objects. These QSOs, therefore, represent a useful sample to investigate the soft excess in both radio-quiet and radio-loud AGN.  \\begin{table*} \\begin{center} \\caption{Redshifts and Galactic absorption for the seven objects. Radio data were obtained from NVSS (Condon \\etal 1998). Optical data were taken from the NASA Extragalactic Database (NED), while the UV magnitudes were obtained from the Optical Monitor (OM) onboard {\\it XMM-Newton} where possible (Mason \\etal 2001). The corresponding wavelengths for the different bands are V -- 550~nm; UVW1 -- 291~nm; UVM2 -- 231~nm; UVW2 -- 212~nm} \\label{z_wa} $^{a}$ Radio fluxes from the NVSS. \\begin{tabular}{p{1.8truecm}p{0.9truecm}p{1.8truecm}p{1.0truecm}p{1.8truecm}p{1.0truecm}p{1.0truecm}p{1.0truecm}p{1.0truecm}} \\hline object & RL/RQ & 1.4 GHz &redshift & Gal. abs. & \\multicolumn{4}{c}{Optical and UV magnitudes}\\\\ & & flux (mJy)$^{a}$ & & (10$^{20}$~cm$^{-2}$) & V-band & UVW1 & UVM2 & UVW2\\\\ \\hline Q0056$-$363 & RQ & $<$~2.5  &0.162 & 1.93 & 16.7 & --- & --- & 13.9\\\\ PG 0804+761 & RQ & 3.3~$\\pm$~0.4 & 0.100 & 2.98 & 15.2 & --- & --- & ---\\\\ Mrk 1383 & RQ &  2.7~$\\pm$~0.5 &0.086 & 2.85 & 17.5 &  --- & 12.4 & 12.4\\\\ Mrk 876 & RQ & 3.9~$\\pm$~0.5 & 0.129 & 2.87 & 15.2 & 13.7 & 13.7 & 13.7\\\\ B2 1028+31 & RL & 230.5~$\\pm$~8 &  0.178 & 1.96 & 16.7 & --- & --- & ---\\\\ B2 1128+31 & RL & 369.7~$\\pm$~13  & 0.289 & 2.02 & 16.6& 15.3 & 14.4 & 14.3\\\\ B2 1721+34 & RL & 518.3~$\\pm$~19.7 & 0.206 & 3.11 & 16.5 & --- & ---  & ---\\\\ \\hline \\end{tabular} \\end{center} \\end{table*} ",
        "conclusions": "The soft excesses of four radio-quiet and three radio-loud QSOs are analysed. As a simple parametrization, two blackbody components fit the soft excess very well. More physically, Comptonization of the disc photons is invoked to explain the soft X-ray emission; this model also provides good fits. There is little obvious intrinsic difference between the soft excess in the radio-quiet and radio-loud objects, with the electron temperatures covering the same range. Five of the objects also showed evidence for iron emission, with three of them being better fitted with a somewhat broadened emission line. These widths are not very large, however, and are significant only at the 95~per~cent level."
    },
    "0404/astro-ph0404255.txt": {
        "abstract": "We present results from a series of high-resolution \\nbody\\ simulations that focus on the formation and evolution of eight dark matter halos, each of order a million particles within the virial radius. We follow the time evolution of hundreds of satellite galaxies with unprecedented time resolution, relating their physical properties to the differing halo environmental conditions.  The self-consistent cosmological framework in which our analysis was undertaken allows us to explore satellite disruption within live host potentials, a natural complement to earlier work conducted within static potentials.  Our host halos were chosen to sample a variety of formation histories, ages, and triaxialities; despite their obvious differences, we find striking similarities within the associated substructure populations. Namely, all satellite orbits follow nearly the same eccentricity distribution with a correlation between eccentricity and pericentre. We also find that the destruction rate of the substructure population is nearly independent of the mass, age, and triaxiality of the host halo.  There are, however, subtle differences in the velocity anisotropy of the satellite distribution.  We find that the local velocity bias at all radii is greater than unity for all halos and this increases as we move closer to the halo centre, where it varies from 1.1 to 1.4. For the global velocity bias we find a small but slightly positive bias, although when we restrict the global velocity bias calculation to satellites that have had at least one orbit, the bias is essentially removed. ",
        "introduction": "There is mounting evidence that the Cold Dark Matter (CDM) structure formation scenario provides the most accurate description of our Universe. Observations point towards a ``standard'' \\LCDM\\ Universe comprised of 28\\% dark matter, 68\\% dark energy, and luminous baryonic matter (i.e. galaxies, stars, gas, and dust) at a mere 4\\% (cf. Spergel~\\ea 2003). This so-called ``concordance model'' induces hierarchical structure formation whereby small objects form first and subsequently merge to form progressively larger objects (White \\& Reese 1978; Davis \\ea 1985; Tormen 1997). The outcome of such mergers, however, depends on many factors (e.g. the mass ratio of the merging halos, their relative velocities, etc.), the result of which is a varied mass accretion history for any given host system. While generally successful, the \\LCDM\\ model does face several problems, one such problem being the prediction that one-to-two orders of magnitude more satellite galaxies should be orbiting their host halos than are observed (Klypin~\\ea 1999; Moore~\\ea 1999).  The lack of observational evidence for these satellites has led to the suggestion that they are completely (or almost completely) dark, with strongly suppressed star formation due to the removal of gas from the small protogalaxies by the ionising radiation from the first stars and quasars (Bullock et~al. 2000; Tully et~al. 2001; Somerville 2002). Others suggest that perhaps low mass satellites never formed in the predicted numbers in the first place, indicating problems with the \\LCDM\\ model in general, replacing it with Warm Dark Matter instead (Knebe~\\ea 2002; Bode, Ostriker~\\& Turok 2001; Colin~\\ea 2000).  Recent results from (strong) lensing statistics suggest that the predicted excess of substructure is in fact required to reconcile some observations with theory (Dahle~\\ea 2003, Dalal~\\& Kochanek 2002), although this conclusion has not been universally accepted (Sand~\\ea 2003; Schechter~\\& Wambsganss 2002; Evans~\\& Witt 2003). If, however, the lensing detection of halo substructure \\textit{is} correct and the overabundant satellite population really does exist, it is imperative to understand the orbital evolution of these objects and their deviation from the background dark matter distribution.  The work described here focuses upon a set of numerical simulations of structure formation within the concordance model, analysing in detail the temporal and spatial properties of satellite galaxies residing within host dark matter halos. To date, typical satellite properties such as orbital parameters and mass loss under the influence of the host halo have primarily been investigated using \\textit{static} potentials for the dark matter host halo (e.g. Johnston \\ea 1996; Hayashi \\ea 2003).  We stress that each of these studies have provided invaluable insights into the physical processes involved in satellite disruption; our goals was to augment these studies by relaxing the assumption of a static host potential, in deference to the fact that realistic dark matter halos are not necessarily axis-symmetric.  Halos constantly grow in mass through slow accretion and violent mergers, possessing rather triaxial shapes (Warren et~al. 1992). While a self-consistent cosmological modeling of both hosts and satellites has long been recognised as optimal, the required mass and force resolution can be difficult to accommodate (hence the use of static host potentials in most previous studies).  The first fully self-consistent simulations targeting the subject were performed by Tormen (1997) and Tormen~\\ea (1998).  Both studies were landmark efforts, but lacked the temporal, spatial, and mass resolution necessary to explore a wide range of environmental effects. Unable to follow the satellite distribution within the host's virial radius, satellites were instead tracked only up to and including the point of ``accretion''.  This allowed an analysis of the infall pattern, rather than the orbital evolution of the satellites.  Ghigna \\ea (1998) also investigated the dynamics of satellite galaxies in live dark matter host halos. Although greatly increasing the mass and spatial resolution, they still lacked the temporal resolution to explicitly track the satellite orbits. Instead, the orbits were approximated using a spherical static potential. More recently, Taffoni~\\ea (2003) used \\nbody\\ simulations coupled with semi-analytical tools to explore the evolution of dark matter satellites inside more massive halos. However, they focus their efforts on the interplay between dynamical friction and tidal mass loss in determining the final fate of the satellites.  Kravtsov~\\ea (2004) also mainly concentrate on the mass loss history of satellites using fully self-consistent cosmological \\nbody\\ simulations.  In this paper we investigate the evolution of substructure and the orbital parameters of satellites using high spatial, mass, {\\it and} temporal resolution. As outlined in Paper~I (Gill, Knebe~\\& Gibson 2004; hereafter \\GKGI), our suite of simulations has the required resolution to follow the satellites even within the very central regions of the host potential ($\\geq$5--10\\% of the virial radius) and the time resolution to resolve the satellite dynamics with excellent accuracy ($\\Delta t \\approx$170~Myrs).  The outline of the paper is as follows. Section~\\ref{Computation} provides a description of the cosmological simulations employed.  The analysis of the host halo and environment can be found in Section~\\ref{HaloAnalysis}, with the satellite orbital parameters presented in Section~\\ref{SatAnalysis}. We then investigate the kinematic properties of the dark matter halos and satellites in Section~\\ref{velbias}. We finish with our summary and conclusions in Section~\\ref{Conclusions}.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %                         Methodology \t\t\t\t\t     % %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{Conclusions}  If the hierarchical model of structure formation is correct then the dynamics of satellite galaxies are an important ingredient to also understanding the formation and the evolution of galaxies. Therefore, in this paper we presented a series of self-consistent cosmological simulations of dark matter halos with the required mass and spatial resolution to follow satellite galaxies orbiting even within the central regions of the host potential. Moreover, the simulations had sufficient time resolution to actually resolve the satellite dynamics with high precision.  The first part of this study was dedicated to analysing and describing the differences in the host halos.  These host halos were chosen to sample a variety of triaxialities, formation times and mass/satellite accretion, despite being of comparable mass. The halos also had very different formation histories, from quiescent to violent. When investigating the halo environment we quantified a value of ``richness'' which was defined to be the fraction of satellite galaxies that the halo had accreted since its formation. One interesting result from this analysis was the similar rate of satellite disruption seen for all eight halos. Furthermore, over the history of each of our halos on average about 10\\% of its respective mass is locked up in the satellite galaxies.  Much of this mass can be attributed to massive satellites rapidly captured by dynamical friction.  But even though the eight dark matter host halos were quite different, their respective satellite population showed remarkable similarities. The average orbital eccentricity of the satellites was found to be $e \\approx$ 0.61 with minimal scatter ($\\sigma\\approx0.19$). Moreover, the average pericentre distance of the satellites was $p \\sim$ 35\\% of the virial radius for all halos, again with minimal scatter ($\\sigma\\approx 0.12$). Satellites that were disrupted while orbiting within the host's virial radius were replaced with a mass-less tracer particle and hence we were also able to present their orbital parameters at redshift $z=0$. We found that even though they have smaller eccentricities ($e^d\\approx 0.34$) than the surviving ones, their pericentre distributions are nearly identical.  Since the pericentres distributions of both surviving and disrupted satellites were similar, implication is that the disrupted satellites spend more time in the deeper regions of the potential well. As such, they experience stronger tidal forces for longer periods, and are thus being disrupted more readily. We also noticed that satellites with more orbits tend to have smaller eccentricities. Difficult to explain through the application of dynamical friction we attribute this to the satellite's response to the growing host halo.  We also found that the local velocity bias at all radii is greater than one and this increases as we move closer to the halo centre. Since this is a characteristic for each of our halos, it strengthens the case that this is a general pattern of the satellite population in dark matter halos.  For the global velocity bias we find an average $\\langle b_{v, \\rm global}\\rangle \\sim 1.103 \\pm 0.002$, a  slight, but significant, positive bias. Further, if we restrict the global velocity bias calculation to satellites that have had at least one orbit we observe a 7\\% decrease in bias to $\\langle b_{v, \\rm global}\\rangle \\sim 1.021 \\pm 0.002$. Thus when we just consider the ``virialised'' satellites, the bias nearly vanishes. Finally we recovered the $\\sigma_v \\sim M^{1/3}$ relationship between satellite velocity dispersion and halo mass.  Surprisingly, all the above stated results appear to be independent of the actual host halo and its history. We were unable to identify any trends with richness, triaxiality and/or formation time (other than the number of orbits). Such similarities are suggestive of potential additional underlying CDM universal laws.    %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % %ACKNOWLEDGEMENTS % %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0404/astro-ph0404567_arXiv.txt": {
        "abstract": "Many recent analyses have indicated that large scale WMAP data display anomalies that appear inconsistent with the standard cosmological paradigm.  However, the effects of foreground contamination, which require elimination of some fraction of the data, have not been fully investigated due to the complexity in the analysis.  Here we develop a general formalism of how to incorporate these effects in any analysis of this type.  Our approach is to compute the full multi-dimensional probability distribution function of all possible sky realizations that are consistent with the data and with the allowed level of contamination. Any statistic can be integrated over this probability distribution to assess its significance.  As an example we apply this method to compute the joint probability distribution function for the possible realizations of quadrupole and octopole using the WMAP data. This 12 dimensional distribution function is explored using the Markov Chain Monte Carlo technique. The resulting chains are used to asses the statistical significance of the low quadrupole using frequentist methods, which we find to be 3-4\\%. Octopole is normal and the probability of it being anomalously low or as low as WMAP reported value is very small. We address the quadrupole-octopole alignment using several methods that have been recently used to argue for anomalies, such as angular momentum dispersion, multipole vectors and a new method based on feature matching. While we confirm that the full sky map ILC suggest an alignment, we find that removing the most contaminated part of the data also removes any evidence of alignment: the probability distributions strongly disfavor the alignment. This suggests that most of the evidence for it comes from non-Gaussian features in the part of the data most contaminated by the foregrounds.  We also present an example, that of octopole alignment with the ecliptic, where the statistical significance can be enhanced by removing the contamination. ",
        "introduction": "The large scale structure of the WMAP data \\citep{2003MNRAS.346L..26E} has received a lot of attention since the first data release. Some authors have focused on the seemingly low values of quadrupole and octopole \\cite{2003MNRAS.346L..26E,2004MNRAS.348..885E,2003PhRvD..68l3523T,2004astro.ph..3073S}, and their alignment \\citep{2004PhRvD..69d063516,2004astro.ph..3353S}, while others considered various asymmetries in the data \\citep{2003astro.ph..7507E,2004astro.ph..3098E,2004astro.ph..2396H,2004astro.ph..4206H}. Some of these analyses are performed on one of the available full sky WMAP maps, either the original WMAP Internal Linear Combination Map (ILC) or an alternative map \\cite{2003PhRvD..68l3523T}, which we will refer to as TOH map. The full sky maps have the advantage that the harmonic analysis is unique, which facilitates investigation and assessment of statistical significance. We note that \\cite{2004astro.ph..3098E} prudently warn against the usage of the foreground corrected full-sky maps and use appropriate Monte-Carlo simulations.  However, full sky maps are not free of contamination, as was clearly emphasized by the WMAP team warning that their ILC map should not be used for science purposes. These are dominated by galactic foregrounds such as dust, synchrotron or free-free emission.  For this reason the power spectrum analysis is done on cut sky, where about 15-25\\% of most contaminated data in the galactic plane is removed.  Even outside this region there are residual uncertainties associated with imperfections of the foreground removal.  If ignored they may cause spurious alignments or other anomalies that appear statistically significant under the assumption that CMB is a Gaussian random field.  In this paper we revisit the statistical significance of these tests using a different approach.  Rather than ignoring the effects of foregrounds we try to take them into account explicitly by exploring the uncertainties they induce in the measurements of the multipole moments $a_{\\ell m}$.  We assess this uncertainty by determining the joint multi-dimensional probability distribution function for the true sky multipole moments. Once these are determined we can apply them to the statistics of choice to obtain their values in the presence of uncertainties associated with imperfect foreground removal or sky cuts.  We do not address the question of the meaning of a given statistic: all of the statistics are a-posteriori and their statistical significance is difficult to assess. Instead, our goal is to compare the values of these statistics with and without the inclusion of foreground uncertainties to see if including the latter changes the conclusions significantly.  In this paper we are interested in large scale features, so we focus exclusively on quadrupole and octopole.  This allows us to perform several tests.  First, we can revisit the question of whether the quadrupole and octopole are low in the presence of additional uncertainties associated with the foregrounds and sky cuts.  Secondly, we can test how robust results of methods for measuring the alignment of quadrupole and octopole are, once these uncertainties are taken into account.  We do this by applying several of existing statistics, as well as a new one we developed. Finally, we also explore the alignment of the large scale features with specific directions in the sky, such as that of ecliptic plane. ",
        "conclusions": " from our analysis it appears impossible for the octopole to be anomalously low and the same is true for quadrupole in our Wd analysis.  The difference may be a consequence of the noisier estimator used by WMAP.   We also discuss recent claims that the quadrupole and octopole are aligned. If one believes the ILC map, then the evidence for the quadrupole and octopole alignment is considerable. All three methods tested here, namely the maximum angular dispersion vectors, the multipole vectors and the feature matching method indicate that the two are suspiciously aligned. However, as soon as foreground uncertainties are included the evidence for this alignment disappears. It is not unexpected that the probability distributions broaden, but what is surprising is how rapidly the evidence vanishes and how strongly perfect or even partial alignment is excluded by the data. This strongly suggests that much of the evidence of the alignment comes from the portion of the data most contaminated by the galactic foregrounds.  Plots shown in this paper indicate a considerable difference between Wd and Vdfs cases. It is interesting to explore whether this difference comes from the number of templates being marginalised over or whether they are due to different frequency channels being employed. To investigate this we repeat the analysis using W channel data and marginalise over all three templates. The results are very similar to the Wd case, indicating that the main difference is due to different channel being employed.  This suggests there might be additional systematic effects that are not handled properly even by our more conservative treatment.  One possibility is that either V or W channel MEM derived foreground maps are contaminated on the largest scales, or that there are systematic contaminations in CMB maps \\cite{2004astro.ph..3353S}.  More work is needed to explore these various possibilities.  Finally, we also present an example where our method can enhance the statistical significance by removing the contamination which would otherwise mask the evidence.  We show that for the alignment of multipole vectors with ecliptic plane the statistical significance is not lowered by the foreground uncertainties.  Once again, the statistical significance of this effect is unclear, but at least it seems clear that it is not significantly affected by the foregrounds and may even be enhanced, once the foreground contamination is removed from the data.  Our method is statistical and relies on the foreground templates to be faithful representation of all components that contaminate the data. Systematic uncertainties in the foregrounds translate in systematic uncertainties in derived quantities. Although our method provides a statistical framework for assessing the statistical significance of various effects, the real improvement will come from better understanding and modelling of the foregrounds. This goal can be achieved using multi-frequency data combined with a better understanding of physical processes involved. In this case a less conservative treatment of the foregrounds may be possible. We have tried some of these examples in our tests. Analyzing the original V or W maps without sky cuts, but with foreground template marginalization, causes MCMC sampler not to converge, which is indicative of a complex likelihood in this case.  Similarly, using sky cuts but with no foreground removal shows clear evidence of contamination and causes the quadrupole and octopole to increase significantly \\cite{2004astro.ph..3073S}.  Subtracting the WMAP recommended foregrounds and using $\\lambda=1$ in equation \\ref{eq:lambda} gives results very similar to $\\lambda=\\infty$. Thus we believe our results reflect the current uncertainties in the foreground subtraction and suggest these may be responsible for many of the anomalies seen in the WMAP data on large scales."
    },
    "0404/astro-ph0404084_arXiv.txt": {
        "abstract": "We analyse the rest--frame (U$-$V) colour--magnitude relation for $2$ clusters at redshift $0.7$ and $0.8$, drawn from the ESO Distant Cluster Survey.  By comparing with the population of red galaxies in the Coma cluster, we show that the high redshift clusters exhibit a deficit of passive faint red galaxies.  Our results show that the red--sequence population cannot be explained in terms of a monolithic and synchronous formation scenario.  A large fraction of faint passive galaxies in clusters today has moved onto the red sequence relatively recently as a consequence of the fact that their star formation activity has come to an end at $z<0.8$. ",
        "introduction": "\\label{sec:intro} Red cluster galaxies form a tight sequence in the colour--magnitude diagram that, in nearby clusters, extends over a range of at least $5$--$6$ mag from the Brightest Cluster Galaxy (BCG).  The existence of a tight colour--magnitude relation (CMR) up to redshift $\\sim 1$, and the evolution of its slope and its zero--point as a function of redshift, are commonly interpreted as the result of a single formation scenario in which cluster ellipticals constitute a passively evolving population formed at high redshift ($z\\gtrsim2$--$3$) in a monolithic collapse (\\cite[Kodama et al. 1998]{K98}).  In this model, the slope of the CMR reflects metallicity differences and naturally arises taking into account the effects of supernovae winds.  An alternative explanation has been proposed by \\cite{KC} and confirmed recently by \\cite{DLKW}.  In this model, elliptical galaxies form through mergers of disk systems and a CMR arises because more massive ellipticals form by mergers of more massive, and hence more metal rich, disk systems.  On the other hand, it is clear that red passive galaxies in distant clusters constitute only a subset of the passive galaxy population in clusters today (\\cite[van Dokkum \\& Franx 1996]{DF}).  Distant clusters contain significant populations of galaxies with active star formation, that can evolve onto the CMR after their star formation activity is terminated, possibly as a consequence of their infall onto clusters (\\cite[Smail et al. 1998; Poggianti et al. 1999]{S98,P99}).  In this work we present the CMR for $2$ clusters at redshifts $0.7$ and $0.8$ from the ESO Distant Cluster Survey (EDisCS). ",
        "conclusions": "\\label{sec:discussion}  The results presented are robust both against the technique adopted for removing non--cluster members and against the photometric errors.  The red galaxy deficit is detected also when rejecting non--members using a purely statistical subtraction or a more stringent criterion for membership based solely on photometric redshifts.  In fact, a deficit is evident also in the {\\it full} photometric catalogue, when no field correction is attempted. Photometric errors in the EDisCS catalogue are comparable to the errors for the Coma data, therefore the differences observed in the distributions of Fig.~\\ref{fig:histo} cannot be a spurious result arising from photometric errors.  A decline in the number of red sequence members at faint magnitudes was first observed in clusters at $z=0.25$ by \\cite{S98}.  Evidence for a `truncation' of the CM sequence has been noticed in a cluster at $z=1.2$ by \\cite{K00}, who suggested that faint early--type galaxies might not have been in place until $z \\sim 1.2$.  In more recent work \\cite{K04} have come to the same conclusion using deep wide-field optical imaging data of the Subaru/XMM-Newton Deep Survey.  The CMR of the $2$ EDisCS clusters at $z\\sim0.8$ also shows a deficit of red, relatively faint galaxies.  Our investigation shows that the evolution of the red--sequence population at these redshifts cannot be explained as the result of a monolithic and synchronous formation scenario.  A large fraction of the passive $\\lesssim 0.1$L$_*$ galaxies must have moved onto the CMR at redshifts lower than $0.8$ as a consequence of the fact that their star formation activity has come to an end.  Our analysis enlightens the importance of studying the evolution of the cluster population as a whole, trying to understand how galaxies are accreted from the field, how the dense cluster environment affects their star formation rate, and how these galaxies fade and move on to the red sequence.  Only this kind of analysis can unveil the full evolutionary paths of galaxies that lie on the red--sequence today and give strong constraints on the relative importance of star formation and metallicity in establishing the the observed red--sequence.  We plan to investigate this in more detail in future work."
    },
    "0404/astro-ph0404421_arXiv.txt": {
        "abstract": "A survey of diffuse interstellar sight lines observed with the {\\em Far Ultraviolet Spectroscopic Explorer} has led to the serendipitous discovery of a high-ionization nebula around the hot white dwarf KPD~0005+5106. The nebula has an \\ion{O}{6} $\\lambda$1032 surface brightness of up to 25,000 photons~s$^{-1}$~cm$^{-2}$~sr$^{-1}$, making it the brightest region of extended \\ion{O}{6} emission in our survey. Photoionization models using the incident white dwarf continuum successfully reproduce the observed \\ion{O}{6} intensity. The \\ion{O}{6} emission arises in the highly ionized inner region of a planetary nebula around KPD~0005+5106. This newly discovered nebula may be one member of a class of high-ionization planetary nebulae that are difficult to detect in the optical, but which can be easily identified in the ultraviolet. ",
        "introduction": "Planetary nebulae (PNe) come in a variety of shapes, irregular, bipolar, or spherical, and exhibit emission from forbidden lines such as [\\ion{O}{3}], [\\ion{Ne}{3}], and [\\ion{Ne}{5}] as well as from the H$\\alpha$ recombination line. In addition, numerous high-excitation PNe exhibit most of the other Balmer lines as well as \\ion{He}{2}, \\ion{C}{3}], and other forbidden lines such as [\\ion{Ne}{4}], [\\ion{Ar}{4}], [\\ion{Mg}{5}]. The typical radial extent is a few tenths of a parsec. Only a few PNe have radii exceeding 1.5~pc. Because of their small size and low surface brightness, PNe are notoriously difficult to discover \\citep[e.g.,][]{kw,w97,rau}.  We report the discovery of an \\ion{O}{6} emitting nebula around KPD~0005+5106 with the {\\em Far Ultraviolet Spectroscopic Explorer (FUSE)} that we identify as a high-ionization PN (HIPN) around this white dwarf (WD). This would make KPD~0005+5106 only the third WD \\citep[out of 162 WDs listed by][]{nap} known to have a PN. Including KPD~0005+5106, only eight of the fifteen stars listed with $T_{\\rm eff}\\geq120,000$~K \\citep{nap} have a known PN. ",
        "conclusions": "With a surface temperature $T_{\\rm eff}=120,000$~K, KPD~0005+5106 is the hottest known DO (i.e. helium-rich) WD. The stellar spectrum has been observed at X-ray, ultraviolet, and optical wavelengths. Model fits yield stellar parameters of $\\log(g)=7$, $E(B-V)=0.13$, and a distance of about 270~pc (Werner, Heber, \\& Fleming 1994; Werner et al. 1996). The star possesses a soft X-ray corona \\citep{odw}, and appears to undergo mass ejection \\citep{w96,sion}.  In the $\\log(g)-\\log(T_{\\rm eff})$ diagram, KPD~0005+5106 lies between the highest-gravity PG~1159 stars known to have PNe and the two DO WD PN candidates PG~0108+101 and PG~0109+111. The discovery of a PN around the DO WD PG~1034+001 \\citep{hew} extends the range of known PNe in the $\\log(g)-\\log(T_{\\rm eff})$ plane along the evolutionary tracks toward lower temperatures and higher surface gravities ($\\log(g)=7.5, T_{\\rm eff}=100,000$~K; Werner, Dreizler, \\& Wolff 1995). KPD~0005+5106 lies well within this extended parameter space. However, previous searches for an optical PN around this WD have been unsuccessful \\citep{w97}.  Photoionized \\ion{O}{6} is expected from PN with hot central stars ($T_{\\rm eff}\\geq120,000$~K) and has been detected in a few such objects (Chu, Gruendl, \\& Guerrero 2004). We used the photoionization code CLOUDY to determine whether photoionization of a putative PN around KPD0005+5106 could explain the detected \\ion{O}{6} emission. For the ionizing source, we used a synthetic spectrum of the star calculated with the program TLUSTY \\citep{hl}, using the observed stellar abundances and parameters \\citep{w96}. The stellar model was scaled to the flux observed with {\\em FUSE}. Simple PN expansion models by \\citet{ost} yield densities of a few cm$^{-3}$ at the distances of the LWRS aperture positions for a $\\sim10,000$~yr old PN. We assumed solar abundances and a particle density of unity (i.e., $n({\\rm H})=1$~cm$^{-3}$) in the absorbing cloud. With no other free parameters, our model reproduces the observed \\ion{O}{6} flux to within a factor of a few. Based on this agreement between model predictions and observations, we conclude that the \\ion{O}{6} nebula around KPD~0005+5106 is photoionized. It may be one of a new class of HIPNe that are faint in standard optical lines."
    },
    "0404/astro-ph0404292.txt": {
        "abstract": "We present a %for the first time a large sample of spectroscopic study of candidate brown dwarf members of the Orion Nebula Cluster (ONC).  We obtained new $J$-- and/or $K$--band spectra of $\\sim$100 objects within the ONC which are expected to be substellar based on their $K,(H-K)$ magnitudes and colors. Spectral classification in the near-infrared of young low mass objects % using near-infrared spectra is described, including the effects of surface gravity, veiling due to circumstellar material, and reddening. %, about a third of which are determined to be substellar. %From these data From our derived spectral types and existing near-infrared photometry we construct an HR diagram for the %low mass members of the cluster. % and possibly discover a previously unknown population %of apparently older $\\sim$10 Myr old stars which we determine are likely cluster members. Masses are inferred for each object and used to derive the brown dwarf fraction and assess the  mass function for the inner 5.'1 $\\times$ 5.'1 of the ONC, down to $\\sim$0.02 M$_\\odot$.  The logarithmic mass function rises to a peak at $\\sim$0.2 M$_\\odot$, similar to previous IMF determinations derived from purely photometric methods, but falls off more sharply at the hydrogen-burning limit before leveling through the substellar regime. % and shows some evidence for %a secondary peak, at $\\sim$0.05 M$_\\odot$. We compare the mass function derived here for the inner ONC to those presented in recent literature for the sparsely populated Taurus cloud members and the rich cluster IC 348. % which were also determined using spectroscopic data, We find good agreement between the shapes and peak values of the ONC and IC 348 mass distributions, but little similarity between the ONC and Taurus results. %We argue that comparison of these %studies suggests the IMF is not universal but instead dependent on the %star formation environment within a cluster. ",
        "introduction": "%Recent photometric and spectroscopic surveys have led to the discovery of %an abundance of substellar objects over a wide mass range, %from just below the hydrogen burning limit ($\\sim$0.08 M$_\\odot$) down to the %planetary mass regime ($<$ 0.02 M$_\\odot$ = 20 M$_{Jup}$). %However, despite the increasing number of brown dwarfs (BDs) known, %details of their formation process and its possible variation %with external conditions remain little understood. %masses formed during a single epoch within a cluster, %also known as the initial mass function (IMF), across %the hydrogren-burning limit.  Major questions that arise are: %is the IMF universal?  We have no reason a priori to assume the IMF is %a unversal function applicable to %all starburts, however available data are consistent with %the same IMF for imtermediate mass stars (M $\\sim$ 1--few M$_\\odot$) %occuring in very different physical environments. %While one might expect instead that the IMF should vary with %star formation environment, we do not yet have enough evidence %to determine if such a variation exists.  This problem can be %most readily examined with the lowest mass objects whose formation %will be most influenced by external environmental factors, e.g., %stellar density and %the content/characterization of the massive star population within %the parental molecular cloud. %Observed The stellar mass and age distributions in young star clusters % of masses and ages of objects found within %a cluster %is crucial to our understanding the formation and early evolution of %stars and brown dwarfs. %With this information we can %will can help answer some of the fundamental questions of cluster formation theory: Do all cluster members form in a single burst or is star formation a lengthy process? %Could star formation be triggered? Is the distribution of stellar masses formed during a single epoch within a cluster, also known as the initial mass function (IMF), universal or does it vary with either star formation environment or time? %Understanding %how low-mass stars form and evolve %remains one of the fundamental challenges of stellar theory. %An important tool for gaining this understanding is %the distribution of stellar %masses formed during a single epoch within a cluster, %also known as the initial mass function (IMF), %remains one of the fundamental challenges of star formation theory. While the stellar mass function has long been studied (e.g., Salpeter 1955), we are only recently beginning to explore the very low mass end of the distribution into the substellar regime. Identification of large, unbiased samples of low mass objects, especially in star-forming regions, is crucial to our understanding of the formation and early evolution of low mass stars and brown dwarfs. %A key issue of this understanding is knowledge of %the distribution of stellar %Knowledge of the shape of the mass function across the hydrogen-burning limit %and its possible variation %with external conditions can provide critical insight into the formation processes %governing low mass stars and brown dwarfs (BDs). %In particular, diagnostic studies of stellar populations in %different locations and at varying stages of evolution %are needed to explore the possibility of a universal mass function. %While one might expect instead that the IMF should vary with %star formation environment, we do not yet have enough evidence %to determine if such a variation exists.  This problem can be %most readily examined with the lowest mass objects whose formation %will be most influenced by external environmental factors, e.g., %stellar density and %the content/characterization of the massive star population within %the parental molecular cloud.  Young stellar clusters are particularly valuable for examining the shape of the low mass IMF because %they provide %statistically significant samples of %stars with common age and metalicity. %In young clusters where the the lowest mass members have not yet been lost to dynamical evolution. %and populations are relatively complete even at the substellar end. Furthermore, contracting low-mass pre-main sequence stars and brown dwarfs are 2-3.5 orders of magnitude more luminous than their counterparts on the main sequence, and thus can be more readily detected in large numbers. The dense molecular clouds associated with star-forming regions also reduce background field star contamination.   %Several photometric studies have been carried out of young clusters in various %environments.  Preibisch, Stanke \\& Zinnecker (2003) used deep imaging data to construct a %$J$-band luminosity function (LF) for the young (2 Myr) cluster IC 348. %From this they determined the cluster to have a 'brown dwarf deficit', whereby objects in the %mass range 0.02-0.075 $M_\\odot$ constitute %at most $\\sim$10\\% of the total cluster population. %A similar lack of sub-stellar objects was found through an optical/IR survey combining %imaging and spectroscopy in the Taurus star forming region (1-2 Myr) %(Brice\\~{n}o et al. 2002). %These results are in contrast to %brown dwarf densities found %within the Orion Nebula Cluster (ONC).  Because the ONC is one of the nearest massive star-forming regions to the Sun and the most populous young cluster within $\\sim$2 kpc, it has been observed at virtually all wavelengths over the past several decades. However, only recently have increased sensitivities due to near-IR detectors on larger telescopes allowed us to begin to understand and characterize the extent of the ONC's young stellar and brown dwarf population which, at $\\lesssim$ 1-2 Myr, is just beginning to emerge from its giant molecular cloud.  Several recent studies have explored the ONC at substellar masses. Hillenbrand \\& Carpenter (2000) (hereafter HC00) present the results of an $H$ and $K$ imaging survey of the inner ~5'.1 x 5'.1 region of the ONC. % taken with NIRC %on Keck I. Observed magnitudes, colors, and star counts were used to constrain the shape of the ONC mass function across the hydrogen burning limit down to $\\sim$0.03 M$_\\odot$.  They find evidence in the log-log mass function for a turnover above the hydrogen-burning limit, then a plateau into the substellar regime. A similar study by Muench et al. (2002; hereafter M02) uses $J,H,K$ imaging of the ONC to derive an IMF which rises to a broad primary peak at the lowest stellar masses between 0.3 M$_\\odot$ and the hydrogen burning limit before turning over and declining into the substellar regime.  However, instead of a plateau through the lowest masses, M02 find evidence for a secondary peak between 0.03--0.02 M$_\\odot$. Luhman et al. (2000) use $H$ and $K$ infrared imaging and limited ground-based spectroscopy to constrain the mass function and again find a peak just above the substellar regime, but then a steady decline through the lowest mass objects. %A more complete and sensitive spectroscopic study is needed if we are to %reliably measure the ONC's IMF %across the hydrogen-burning limit.  Generally speaking, $J,H,K$ photometry alone is insufficient for deriving stellar/substellar masses, though may be adequate in a statistical sense for estimating mass distributions given the right assumptions. The position of a young star in a near-IR color-magnitude diagram (CMD) is dependent on mass, age, extinction, and the possible presence of a circumstellar disk. These characteristics affect the conversion of a star's infrared magnitude and color into its stellar mass.  Unless the distributions of these parameters are known a priori, knowledge of the cluster's luminosity function alone is not sufficient to draw definitive conclusions about its mass function. In addition, cluster membership is often poorly known and statistical estimates concerning the extent and characterization of the field star population must be derived.  In the case of the densely populated ONC, it has been suggested that the field star contamination is small but non-negligible toward fainter magnitudes. HC00 used a modified version of the Galactic star count model (Wainscoat et al. 1992) convolved with a local extinction map (derived from a C$^{18}$O molecular line map) to estimate the field star contribution, which they found to constitute $\\sim$5\\% of the stars down to their completeness limit at $K\\sim$17.5.   In order to study a cluster's IMF in more than just a statistical sense, %rather than just place statistical constraints on it, spectroscopy is needed to confirm cluster membership of individual stars and uniquely determine location in the HR diagram (and hence mass). We have obtained near-infrared spectra of 97 %85 ONC stars in the ONC. %in the area %surveyed by HC00, including %$\\sim$50\\% of the stars expected to be brown dwarfs based on their $K,H-K$ magnitudes %and colors, down to the completeness limit %of $K\\sim$17.5, and of 14 stars outside the survey region of HC00. This wavelength regime (1-2.5 $\\mu$m) is of extreme interest for very cool stars and brown dwarfs not only because ultracool objects emit the bulk of their energy in the near-infrared, but also because this regime contains temperature-sensitive atomic as well as broad molecular features. In addition, there are several diagnostic lines which can be used as surface gravity indicators. From analysis of these data combined with existing photometry and pre-main sequence evolutionary theory we construct the cluster's IMF across the substellar boundary.  We then compare our results to those found from previous studies, both of the ONC and of clusters similar in age to the ONC but which have different star-forming environments.  In Section 2 we describe our data acquisition and reduction.  In Section 3 we present our spectra and methods for spectral classification.  This section includes a discussion of the effects of extinction and veiling due to circumstellar disks. In Section 4 we create an HR diagram for the stellar and substellar objects for which we have new spectral types. Section 5 contains our derivation of the ONC's IMF and Section 6 our analysis and comparison to previous work. ",
        "conclusions": "\\subsection{Comparison to previous ONC IMF determinations}  The stellar/substellar IMF has been discussed in previous work on the ONC. A determination based on near-infrared photometric data was made by HC00 using $H$ and $K$ magnitudes and colors combined with star count data to constrain the IMF down to $\\sim$0.03 M$_\\odot$.  They find a mass function for the inner regions of the ONC which rises to a peak around 0.15 M$_\\odot$ and then declines across the hydrogen burning limit with slope N($\\log M) \\propto$ $M^{0.57}$.  M02 transform the inner ONC's $K$-band luminosity function into an IMF and find %They find a mass function which rises with decreasing mass to form a broad peak between 0.3 M$_\\odot$ and the hydrogen-burning limit before turning over and falling off into the substellar regime.  This decline is broken between 0.02 and 0.03 M$_\\odot$ where the IMF may contain a secondary peak near the deuterium burning limit of $\\sim$13 M$_{Jup}$. Luhman et. al (2000) combined near-infrared NICMOS photometry of the inner 140\" $\\times$ 140\" of the ONC with limited ground-based spectroscopy of the brightest objects ($K <$ 12) to determine a mass function which follows a power-law slope similar, but slightly steeper than the Salpeter value, until it turns over at  $\\sim$0.2 M$_\\odot$ and declines steadily through the brown dwarf regime. %rises slowly through the brown dwarf %regime until $\\sim$0.2 M$_\\odot$ where it rolls over into a power-law slope %similar, but slightly steeper than the Salpeter value. H97 presents the most extensive spectroscopic survey of the ONC, combining optical spectral data with $V$ and $I$-band photometry over a large area ($\\sim$ 2.5 pc$^2$) extending into the outer regions of the cluster.  The IMF determination covers a large spectral range and appears to be rising from the high (50 M$_\\odot$) to low (0.1 M$_\\odot$) mass limits of that survey.   %In Figure~\\ref{fig:mf1} we compare our IMF of the inner $\\sim$0.5 pc$^2$ of the ONC, %which we derived using a combination of infrared %and optical spectral data, to mass functions derived previously for %the region.  The filled circles show the mass function for stars with %A$_V$ $<$ 10 as determined %by HC00 using infrared photometry. %In order to compare the two sets of results %we must also limit %our sample to A$_V$ $<$ 10, thereby creating an equivalent set of data. %The bottom panel of Figure~\\ref{fig:maghist} shows our completeness as a function %of magnitude for all stars with A$_V$ $\\lesssim$ 10 mag.  Since we %do not have derived extinction values for the stars without spectra %in the HC00 survey, we used a color cut of %($H-K$) $<$ 1.0 which approximately corresponds to a 1 Myr star reddened by %10 magnitudes of extinction in the absence of infrared excess (see Figure~\\ref{fig:cmd}). %As in the top panel, the open histogram indicates all stars with photometric %data in HC00 and the hatched histogram indicates the subset for %which we have spectra.  The dotted line represents the fractional %completeness of the spectroscopic sample with $\\sqrt(N)$ errorbars. %We estimate completeness in this sample to $\\sim$50\\% excepting the %magnitude range $K$ $\\sim$ 12.5-14.  We can correct for the underrepresented %bins %as outlined in Section 5.2. % %Solid and dotted hatched histograms in Figure~\\ref{fig:mf1} represent the non-corrected and corrected IMFs for our sample with A$_V$ $<$ 10. %The IMF of HC00 has been scaled to match the total number of stars in our IMF less than 0.4 M$_\\odot$. %The solid squares indicate the IMF derived by Hillenbrand 1997 for the %inner and outer regions of the nebula ($r \\;<$ 2.5 pc).   All stars in this sample have A$_V$ $<$ 10.  Data have been shifted to match the IMF from the %current work at M = 0.2 M$_\\odot$.  All three IMF determinations were made using %the DM97 tracks.   The peak in our IMF, $\\sim$0.2 M$_\\odot$, matches remarkably well to those found from both the deep near-infrared imaging IMF studies (HC00 and M02) which cover similar survey areas, and the Luhman et al. (2000) study which covered only the very inner region of the cluster. %The IMF derived by Luhman et al. (2000) peaked at much higher  looked at only the inner $\\sim$0.3 pc of the ONC.  Hillenbrand %(1997) find the mass spectrum to be biased towards higher masses within %this area which could account for the discrepancy in turnover %mass between Luhman et al. (2000) and other studies which explore %the spatially more extended population. %XXXXXXXXXXXXXXX HC00 \\& H97 XXXXXXXXXXXXXXXXXXXX Our data also show a leveling off in the mass distribution through the substellar regime similar to that found by HC00. A significant secondary peak within the substellar regime has been claimed by M02.  While we see some evidence for such a peak in our data, this result is not robust to within the errors.  Furthermore, if real, we find the secondary peak to occur at a slightly higher mass than M02 ($\\sim$0.05 M$_\\odot$ vs. $\\sim$0.025 M$_\\odot$). %XXXXXXXXXXXXXXX HC00 \\& H97 XXXXXXXXXXXXXXXXXXXX The primary difference between the observed IMF derived in the current work and those presented in previous studies of the substellar ONC population is the steepness of the primary peak and the sharpness of the fall-off at the hydrogen-burning limit (see for comparison Figure 16 of M02). Most IMF determinations for this region exhibit a gradual turnover in the mass function around $\\sim$0.2 M$_\\odot$ until %and then %a leveling off, or turnover to a secondary peak, at $\\approx \\frac{2}{3}$ the level of the primary peak is reached at which point the IMF levels off or forms a secondary peak. However, we find a sharp fall-off beyond $\\sim$0.1 M$_\\odot$ to $\\approx \\frac{1}{2}$ the primary peak value. %This distinction is lessened somewhat when we considered %our magnitude-corrected IMF (Figure~\\ref{fig:mf2}). %This %discrepancy is likely %due the difference in models used.  The M02 work used the   \\subsection{Photometric vs. Spectroscopic Mass Functions}  As mentioned in Section 1, spectroscopy is needed to study a cluster's IMF in more than a statistical sense. %rather than simply put statistical constraints on it. The fact that many of the fainter objects in our survey expected to be substellar based on their location in the $K$, ($H-K$) diagram are in fact hotter, possibly older objects gives strong evidence in support of this statement. Independent knowledge of an individual star's age, extinction and infrared excess arising from the possible presence of a circumstellar disk is needed before definitive conclusions can be drawn about that object's mass. Despite the increasing numbers of brown dwarfs studied, both in the field and in clusters, details of their formation process remain sufficiently poorly understood that knowledge of near-infrared magnitudes and colors alone is not enough to accurately predict these characteristics for young low mass objects.  Near-infrared magnitudes alone also cannot distinguish between cluster members and nonmembers, and field star contamination must be modeled rather than accounted for directly. %In addition, work by Walker et al. (2004) suggests that CTTS occulted by a circumstellar %disk show similar optical/near-IR colors to face-on brown dwarf systems, which may lead to an %incorrect, over-identification of substellar objects in photometric surveys.  Previous studies of the ONC had spectroscopy available in general only for the stellar population, and relied on photometry alone to determine the mass function at substellar masses. %\\footnote{As mentioned, two other %studies have looked at the ONC spectroscopically, but H97 did not %extend to substellar masses and Luhman et al. (2000) took spectra only %of the brightest objects ($K >$ 12) in their sample.}. This may have caused  over-estimates in the number of brown dwarfs for reasons discussed above. %In addition, work by Walker et al. (2004) suggests that CTTS occulted by a circumstellar %disk show similar optical/near-IR colors to face-on brown dwarf systems, which may lead to an %incorrect, over-identification of substellar objects in photometric surveys. %The work by Muench et al. (2001) to look for infrared excess objects %within the ONC note that 21 of their 109 photometrically identified brown dwarf candidates ($\\sim$20\\%) %are coincident with optically resolved proplyds. However, the more significant cause of the shallower peak in previous IMF determinations for the inner ONC as opposed to the sharply peak IMF derived here is likely just the inherent nature of photometric vs. spectroscopic mass functions.  Determining masses for stars in a sample from temperatures derived from spectral types necessarily discretizes the data. Conversely, photometric studies are by their nature continuous distributions, and deriving masses from magnitudes and colors or luminosity functions results in a smooth distribution of masses. We emphasize that while neither situation is ideal, mass functions derived for young objects from infrared photometry alone represents only the most statistically probable distribution of underlying masses. Spectroscopy is needed in order to derive cluster membership and masses for individual objects.     %\\subSection{Comparison of the ONC IMF to IMF determinations for other star-forming regions} \\subsection{Comparison to Other Star-Forming Regions}  Diagnostic studies of stellar populations in different locations and at varying stages of evolution are needed to explore the possibility of a universal mass function. While one might expect that the IMF should vary with star formation environment, we do not yet have enough evidence to determine if such a variation exists. %This problem can be %most readily examined with the lowest mass objects whose formation %may be more heavily influenced by external environmental factors, e.g., %stellar density and %the content/characterization of the massive star population within %the parental molecular cloud. Aside from work already mentioned on the ONC, numerous studies have been carried out to characterize the low mass stellar and substellar mass functions of other young clusters in various environments. Because of the intrinsic faintness of these objects, most surveys are photometric. Authors then use a combination of theoretical models and statistical analysis to transform a cluster's color-magnitude diagram or luminosity function into an IMF which may or may not accurately represent the underlying cluster population (see Section 6.2).  However, the substellar populations of two other young star-forming clusters have been studied spectroscopically using techniques similar to those presented here. Luhman (2000), Brice\\~{n}o et al. (2002) and Luhman et al. (2003a) surveyed the sparsely-populated Taurus star-forming region, and Luhman et al. (2003b) studied the rich cluster IC 348. These clusters have ages similar to the ONC (1 and 2 Myr). Therefore, if the IMF is universal, similar mass distributions should be observed for all three clusters.  Contrary to this hypothesis, Luhman et al. (2003b) discuss the very different shapes of the substellar mass distributions in Taurus and IC 348. The IMF for Taurus peaks around $\\sim$0.8 M$_\\odot$ and then declines steadily to lower masses through the brown dwarf regime. The IC 348 mass function rises to a peak around 0.15 M$_\\odot$ and then falls off sharply and levels out for substellar objects. While direct comparison of the data requires caution given that different mass tracks were used for the two studies (Luhman et al. (2003b) used Baraffe et al. (1998) tracks whereas we have used DM97 tracks to infer masses), we find our that IMF for the ONC bears remarkable resemblance to that presented for IC 348 in Luhman et al. (2003b) (see for comparison Figure 12 in Luhman et al. 2003b).  Both IMFs peak at $\\sim$0.15--0.2 M$_\\odot$ and fall off rather abruptly at the substellar boundary. The fact the IMFs derived for these two dense, young clusters %using spectroscopic techniques bear such close resemblance to each other while exhibiting distinguishable differences from the IMF determined for the much more sparsely- populated Taurus cluster gives strong support to the argument put forth by authors such as Luhman et al. (2003b) that the IMF is not universal, but may instead depend on star formation environment. %It has been suggested that lower density clusters such as Taurus Alternatively, Kroupa et al. (2003) argue through numerical simulation that observed differences in the Taurus and ONC {\\it stellar} mass functions could be due to dynamical effects operating on initially identical IMFs; however, their model does not reproduce the observed differences in {\\it substellar} mass functions without invoking different initial conditions (eg., turbulence; Delgado-Donate et al. 2004).  Previous studies such as those mentioned above have found much higher brown dwarf fractions for the ONC in comparison to other clusters such as Taurus and IC 348 (see Luhman et al. 2003b). Brice\\~{n}o et al. (2002) defines the ratio of the numbers of substellar and stellar objects as: \\begin{displaymath} R_{SS} = \\frac{N(0.02 \\leq M/M_\\odot \\leq 0.08)}{N(0.08 < M/M_\\odot \\leq 10)}. \\end{displaymath} We have recomputed these numbers for Taurus and IC 348 using the DM97 models and find values of $R_{SS}$ = 0.11 (Taurus) and $R_{SS}$ = 0.13 (IC 348).  These values are very close to those found by Luhman et al. (2003b) using the Baraffe et al. (1998) models: $R_{SS}$ = 0.14 (Taurus) and $R_{SS}$ = 0.12 (IC 348). In contrast, Luhman et al. (2000) find a value of $R_{SS}$ = 0.26 from primarily photometric work on the inner ONC (using Baraffe et al. (1998) models). Considering the spectroscopic IMF presented here (Figure 14), we find a lower value of $R_{SS}$=0.20, indicating that although the ONC may have a higher brown dwarf fraction than Taurus or IC348, it is lower than previously inferred from photometric studies. %Combining our data with masses derived for the higher mass stars in H97, %we find a lower value of $R_{SS}$ = 0.20 indicating that the ONC %may actually have a lower brown dwarf fraction than previously thought.  We believe the numbers of substellar objects %in the ONC may have been over estimated in previous photometric surveys for reasons given in Section 6.2.  %Preibisch, Stanke \\& Zinnecker (2003) used deep imaging data to construct a %$J$-band luminosity function (LF) for the young (2 Myr) cluster IC 348. %From this they determined the cluster to have a 'brown dwarf deficit', whereby objects in the %mass range 0.02-0.075 $M_\\odot$ constitute %at most $\\sim$10\\% of the total cluster population. %A similar lack of sub-stellar objects was found through an optical/IR survey combining %imaging and spectroscopy in the Taurus star forming region (1-2 Myr) %(Brice\\~{n}o et al. 2002). %These results are in contrast to %brown dwarf densities found %within the Orion Nebula Cluster (ONC)."
    },
    "0404/astro-ph0404347_arXiv.txt": {
        "abstract": "Broad hydroxyl (OH) absorption-lines in the 1667 MHz and 1665 MHz transition towards the central region of NGC\\,3079 have been observed at high resolution with the European VLBI Network (EVN). Velocity fields of two OH absorption components were resolved across the unresolved nuclear radio continuum of $\\sim$ 10 parsecs.  The velocity field of the OH absorption close to the systemic velocity shows rotation in nearly the same sense as the edge-on galactic-scale molecular disk probed by CO(1--0) emission. The velocity field of the blue-shifted OH absorption displays a gradient in almost the opposite direction.  The blue-shifted velocity field represents a non-rotational component, which may trace an outflow from the nucleus, or material driven and shocked by the kiloparsec-scale superbubble. This OH absorption component traces a structure that does not support a counter-rotating disk suggested on the basis of the neutral hydrogen absorption. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404171_arXiv.txt": {
        "abstract": "{ We present the results of high S/N long-slit spectroscopy with the Multiple Mirror and the SAO 6-m telescopes, optical imaging with the Wise 1-m telescope and \\ion{H}{i} observations with the Nan\\c {c}ay Radio Telescope of the very metal-deficient (12+$\\log$(O/H)=7.64) luminous ($M_{\\rm B}$=--18\\fm1) blue compact galaxy (BCG)  HS~0837+4717. The blue bump near $\\lambda$4650, characteristic of WR stars is detected in the central supergiant \\ion{H}{ii} region, as well as the barely seen red bump at $\\lambda$5808. The derived number of WR stars in the region of the current starburst is $\\sim$1000. Evidence for fast motions in this region is also seen as broad, low-contrast components in the H$\\alpha$, H$\\beta$ and strong [\\ion{O}{iii}] lines $\\lambda\\lambda$4959,5007. While the extinction of the narrow emission lines from the supergiant \\ion{H}{ii} region is low, the very large Balmer decrement of the broad components suggests that the part of current starburst is highly obscured by dust. Abundance ratios X/O for X=Ne, Ar, S, Fe and Cl in the supergiant \\ion{H}{ii} region are in good agreement with the mean values of other very metal-deficient BCGs. Nitrogen, however, is overabundant by a factor of $\\sim$6. This implies an unusually efficient N enrichment in HS 0837+4717, and probably, a non-typical evolution scenario. The H$\\alpha$-line Position--Velocity (P--V) diagrams for directions approximately along the major and minor axes reveal disturbed motions of the ionized gas, mainly in peripheral regions. The SW part of the major axis P--V diagram looks like a rotation curve, with the velocity amplitude $V_{\\rm rot} \\sim$50--70~\\kms\\ at $r$=4.3 kpc. Its NE part displays, however, strong deviations, indicating either counter-rotation, or a strong outflow/supershell. If it is considered as indicating a shell-like feature its velocity amplitude of $\\sim$70 \\kms\\ (relative to the extrapolated rotation curve), and the apparent extent of $\\sim$4\\arcsec\\ (3.3 kpc) imply a dynamical age of $\\sim$14 Myr and the full energetic equivalent of $\\sim$2.3$\\times$10$^4$ SNe. The latter indicates continuing starbursts during at least this time interval. The long-slit spectra reveal a complex morphology for this galaxy. It consists of two compact regions at a distance of $\\sim$2 kpc. Their continuum flux differs by a factor of four. The brightest one is related to the current starburst with the age of $\\sim$3.7 Myr. The slightly redder fainter component could be an older starburst ($\\sim$25 Myr).  The  Wise 1-m telescope $UBVR$ integrated photometry reveals a high optical luminosity for this BCG, and the unusual $(B-V)$ and $(V-R)$ colours. The morphology of HS 0837+4717 is highly disturbed, with two small tails emerging to NNW and SSE. Such a disturbed overall morphology, a \"double-nucleus\" structure, significantly disturbed velocities of ionized gas, together with the very high power of the starburst suggests a possible explanation of the object as a recent merger. We compare the properties of this BCG and of similar objects known in the literature, and conclude that their high nitrogen excess is most probably related to the short phase of a powerful starburst when many WR stars contribute to the enrichment of ISM. ",
        "introduction": "Most low-mass gas-rich galaxies have low metallicities, with the typical range of 1$Z$\\sunn/15 to $Z$\\sunn/3  (e.g., review of Kunth \\& \\\"Ostlin \\cite{Kunth2000}). Only for a very small fraction of galaxies in the Local Universe is the metallicity extremely low, in the range of $Z$\\sunn/50 to $Z$\\sunn/20. This range is more characteristic of high-redshift damped Ly$\\alpha$ % absorption systems which could be young galaxies. The study of the extremely metal-deficient (XMD) galaxies allows one to test many theoretical ideas on massive star formation and evolution, and models of galaxy evolution in a very low metallicity environment.  In particular, the abundance patterns of $\\alpha$-elements and iron, that is, the ratios X/O (with X=Ne, S, Ar, Fe), allow one to check models of stellar nucleosynthesis. For metal-poor galaxies they were found to be remarkably constant over a wide range of O/H (e.g., Izotov \\& Thuan \\cite{IT99}, hereafter IT99). This implies that all these elements are primary and are produced in the same massive stars. The ratio N/O was also found to be fairly constant for XMD BCGs (IT99), implying mainly primary production of nitrogen in very low metallicity environment.  Wolf-Rayet (WR) stars -- a specific very short phase of massive star evolution, characterized by strong mass outflow, and enriched mainly in N and C, are often seen in high S/N spectra of starburst galaxies (e.g., Kunth \\& Joubert \\cite{KJ85}; Izotov et al. \\cite{Izotov96}; Guseva et al. \\cite{Guseva2000}). These stars are detected through the broad emission features of the so called ``blue bump'' near $\\lambda$4650~\\AA, and more rarely in the red, near $\\lambda$5808~\\AA. Models of stellar evolution predict that the number of WR stars in a young star cluster is a sensitive function of metallicity, IMF and  starburst age (e.g., Schaerer \\& Vacca \\cite{SV98}, hereafter SV98).  WR star outflows could produce a significant N overabundance at the locations of young starbursts. However, the analysis of a large sample of \\ion{H}{ii} galaxies by Kobulnicky \\& Skillman (\\cite{KS96}) showed that galaxies, with and without strong WR features in their integrated spectra, show identical N/O ratios. This is probably explained by nitrogen outflow in the hot phase. Therefore, only significantly after its ejection could the nitrogen be cooled and mixed into the optically observed ionized gas.  Significant nitrogen overabundance is detected in a few starburst galaxies. The most prominent one is Mkn~996 (Thuan et al. \\cite{Thuan96}), with 12+$\\log$(O/H)=8.0. It is a non-typical BCG due to the large number of WR stars and to a powerful outflow from the compact central SF burst. Its N/O in the region with diameter of 0.6 kpc is 4 to 25 times higher than the typical value for low-metallicity BCGs of $\\sim$1/40. The morphology of Mkn 996 implies that likely it is a remnant of a recent merger. The authors argue that the observed large N/O ratio is directly related to the powerful outflows of numerous WR stars.  Another well known case is a factor of three nitrogen overabundance, found in two compact regions ($d \\sim$20 pc) near the central starburst in the nearby ($D\\sim$4~Mpc) dwarf WR galaxy NGC~5253 (e.g., Kobulnicky et al. \\cite{Kobul97}, and the summary of previous results therein).  Only three galaxies: UM~420, Mkn ~1089, and UM~448, all luminous and with detected WR lines (Guseva et al. \\cite{Guseva2000}), from the sample of 50 BCGs in IT99, show nitrogen overabundance of a factor of $\\approx$3 relative to the mean N/O value of 1/30 for their range of O/H). Finally, the galaxy Haro 11 from the sample of luminous BCGs (Bergvall \\& \\\"Ostlin \\cite{Bergvall02}), also with strong WR bump, shows an N/O ratio that is 6 times higher than for the main BCG group.  In this paper we present the results of a comprehensive study of HS~0837+4717,  found in the Hamburg/SAO survey (Pustilnik et al. \\cite{Pustilnik99}). The very low metallicity of HS 0837+4717 ($Z\\approx$$Z$\\sunn/20) was first derived from the observations with the SAO RAS 6-m telescope (BTA) in 1996 (Kniazev et al. \\cite{Kniazev00a}). Classified as an XMD BCG, it is unusually luminous for its metallicity. This galaxy shows WR bumps and broad components of strong emission lines of hydrogen and oxygen. This is the next most nitrogen overabundant galaxy known after Mkn 996. But, in contrast to the latter, its optical spectrum does not look atypical for BCGs, and its very large nitrogen overabundance is seen on an otherwise more or less typical BCG background.  In Sect. \\ref{Obs} we describe the observations and their reduction. In Sect. \\ref{Results} all the results from the reduction and preliminary analysis are presented, including the heavy element abundances, the WR and broad components of strong lines, the ionized gas kinematics and the morphology and photometry. In Sect. \\ref{discussion} we discuss the obtained results and derived parameters of this galaxy, compare its properties with those of other luminous and XMD  BCGs, and draw conclusions. The adopted distance to the galaxy is 170.5 Mpc; respective scale is 827 pc per arcsecond. ",
        "conclusions": " \\begin{enumerate} \\item The oxygen abundance in HS~0837+4717 is 12+$\\log$(O/H)=7.64. The abundance ratios X/O for the elements Ne, Ar, S, Cl and Fe are consistent with the average values found for low-metallicity BCGs by Izotov \\& Thuan (\\cite{IT99}), implying the primary origin of these elements along with oxygen. \\item HS~0837+4717 is the most metal-deficient galaxy among the luminous BCGs with $M_{\\rm B}\\le$--18\\fm0. This BCG with its O/H and $M_{\\rm B}$ significantly deviates from the $Z$--$L_{\\rm B}$  relationship derived for the large BCG sample. \\item We detected the broad emission features characteristic of WR stars, and estimated the number of WR stars as $\\sim$1000. Their ratio to the number of O-stars is well compatible with that predicted by the current  models. \\item Broad low-contrast components of the H$\\alpha$, H$\\beta$, and [\\ion{O}{iii}]~$\\lambda\\lambda$4959, 5007 emission lines were detected, with FWHMs of $\\sim$1500 \\kms. The flux ratio of the H$\\alpha$ and H$\\beta$ broad components implies very strong obscuration ($\\sim$5 mag in $B$-band) in the emission region or along the line of sight. The extinction-corrected luminosity of the BCG should be higher than that observed, at least, by a factor of 2.5--5. \\item The position-velocity diagrams for the H$\\alpha$ narrow component are not compatible with regular rotation. The major axis P--V diagram indicates either counter-rotation, or the presence of a giant supershell with a diameter of  $\\sim$3.5  kpc and an expansion  velocity of about 70~\\kms. Half of the P--V diagram can be interpreted as rotation with an amplitude of $\\sim$50--70~\\kms\\ at a radius of $\\sim$4 kpc. \\item We find that this BCG is asymmetric both near the center (the inner structure consisting of two compact regions $\\sim$2 kpc apart) and at the periphery. While the disturbed morphology of HS 0837+4717 could suggest a recent strong interaction, no candidate galaxies are found in its vicinity. Along with the very disturbed gas kinematics this could be evidence for a hypothesis of a recent merging of two gas-rich dwarfs. \\item The nitrogen-to-oxygen abundance ratio in HS 0837+4717 is six times higher than that of other XMD BCGs. The properties common to HS 0837+4717 and five other non-XMD BCGs with the large nitrogen abundance excesses,  suggest that the large nitrogen overabundance could be connected with merger events and with a short phase of the related powerful starburst, when many WR stars contribute to the ISM enrichment. \\item HS 0837+4717 is located in a region with very low density of galaxies (void), the nearest of them with known redshift being situated at $\\sim$6.4 Mpc~h$^{-1}$. The strong isolation of the BCG or its progenitor(s) might be the reason for its slow chemical evolution. \\end{enumerate}"
    },
    "0404/astro-ph0404492_arXiv.txt": {
        "abstract": "INTEGRAL/SPI\\unmarkedfootnote{ corresponding author: schanne@hep.saclay.cea.fr \\\\ $^{2}$ Institut d'Astrophysique de Paris, 98 bis Boulevard Arago, 75014 Paris, France \\\\ $^{3}$ F\\'ed\\'eration de Recherche Astroparticule et Cosmologie, % Coll\\`ege de France, 11 Place Marcellin Berthelot, 75231 Paris, France } has recently observed a strong and extended emission resulting from electron-positron annihilation located in the Galactic center region, consistent with the Galactic bulge geometry, without a high energy gamma-ray counterpart, nor in the 1809 keV $^{26}$Al decay line. In order to explain the rate of positron injection in the Galactic bulge, estimated to more than 10$^{43}$ s$^{-1}$, the most commonly considered positron injection sources are type Ia supernovae. However, SN~Ia rate estimations show that those sources fall short of explaining the observed positron production rate, raising a challenging question about the nature of the Galactic positron source. In this context, a possible source of Galactic positrons could be supernova events of a new type, as the recently observed SN2003dh/GRB030329, an exploding Wolf-Rayet star (type Ic supernova) associated with a hypernova/gamma-ray burst; the question about the rate of this kind of events remains open, but could be problematically low. In this paper, we explore the possibility of positron production and escape by such an event in the framework of an asymmetric model, in which a huge amount of $^{56}$Ni is ejected in a cone with a very high velocity; the ejected material becomes quickly transparent to positrons, which spread out in the interstellar medium. ",
        "introduction": "The spectrometer SPI \\citep{schanne2002,attie2003,roques2003,vedrenne2003} on ESA's gamma-ray satellite INTEGRAL, in orbit since October 2002, has recently reported its first results on the observation towards the Galactic center region of the 511 keV gamma-ray line emission resulting from e$^{+}$ e$^{-}$ annihilation \\citep{jean2003}. The flux of annihilation detected by SPI is $\\Phi_{511}=$(0.99$^{+0.47}_{-0.21}$) $\\times$ 10$^{-3}$ ph cm$^{-2}$ s$^{-1}$, is concentrated in a narrow gamma-ray line at an energy of 511.06$^{+0.17}_{-0.10}$ keV and whose intrinsic line width is evaluated to 2.95$^{+0.45}_{-0.51}$ keV (FWHM), while the instrumental resolution at this energy is 2.16 keV. Furthermore, indications on the spatial shape of the 511 keV emission region have also been published by \\cite{jean2003} and \\cite{knodl2003}. A single point source located in the Galactic center region is excluded. When fitting a spherically symmetric Gaussian distribution to the emission region, the best fit is obtained (with a significance level of 12 $\\sigma$) for a Gaussian centered on the Galactic center, with an width of 10$^\\circ$ (FWHM). With 95\\% C.L. the emission region has an extension in the range between 6$^\\circ$ and 18$^\\circ$ (FWHM). A Richardson-Lucy deconvolution confirms this result. In addition, no significant emission from the Galactic disk has been yet detected by SPI.  We know the Galactic center to be located at a distance of $R_{o}=$8.0$\\pm$0.4 kpc with good accuracy, thanks to the precision obtained with a new geometric determination, a Kepler orbit fit to the star S2 orbiting the central black hole, by \\cite{eisen2003}. If we assume that the 511 keV emission takes place near the Galactic center, we can conclude that the spatial size of the emission region is $D=$1.4 kpc in diameter, which coincides roughly with the size of the Galactic bulge. Furthermore, from the 511 keV photon flux observed by SPI we conclude that the 511 keV photon production rate in the Galactic center region is $L_{511}= $7.7$\\times$10$^{42}$ ph s$^{-1}$.  The 511 keV gamma-ray line is produced by $e^{+}$ $e^{-}$ annihilation. \tHowever before annihilation with an $e^{-}$ of the interstellar medium is possible, the $e^{+}$ must first cool down to the thermal energy of the medium it propagates through. Direct $e^{+}$ $e^{-}$ annihilation (with the production of two 511 keV photons per annihilating pair) is possible. However if the temperature is not too high (T$<$10$^6$ K), the formation of an intermediate state, the positronium (Ps), before annihilation is more likely. One quarter of the Ps produced are para-prositronium states (with anti-parallel $e^{-}$ and $e^{+}$ spins), their annihilation produces two 511 keV photons, as does the direct $e^{+}$ $e^{-}$ annihilation. Three quarters of the Ps (ortho-positronium states, with parallel $e^{+}$ and $e^{-}$ spins) however annihilate with the production of 3 photons whose energy ranges from 0 to 511 keV. From the ratio of the fluxes measured for the 511 keV peak and the 3$\\gamma$ continuum, the fraction of $e^{+}$ and $e^{-}$ annihilating via the Ps intermediate state can be computed to be $f_{Ps}=8 (6 + 9 \\Phi_{511} / \\Phi_{3 \\gamma})^{-1}$. As a result of observations of the Galactic center region with CGRO/OSSE, \\cite{kinzer2001} measured a positronium fraction $f_{Ps}=$0.93$\\pm$0.04 and a narrow 511 keV line ($<$3 keV FWHM, confirmed by INTEGRAL/SPI). This result leaves not much room for direct $e^{+}$ $e^{-}$ annihilation, and shows that the bulk of the positron annihilation from the Galactic center direction occurs after positronium formation in a warm medium, excluding molecular clouds as in \\cite{guessum1991} and \\cite{ballmoos2003}.  From the SPI and OSSE measurements the rate of $e^{+}$ annihilation in the Galactic center region can be computed as $L_{e^+}=L_{511}(2 - 3 f_{Ps}/2)^{-1}$. The result is a huge number: every second near the Galactic center $L_{e^+}=$1.3$\\times$10$^{43}$ $e^+$ annihilate in a region comparable in size with the Galactic bulge. Under the assumption of a steady-state production/annihilation, the same amount of positrons must be produced each second near the Galactic center. This conclusion raises the question of the nature of the source capable of injecting such an amount of positrons near the Galactic center, as well as the question of the medium onto which those positrons annihilate, which does not necessarily coincide with the e$^+$ production region, since e$^+$ are transported out of their production site into the medium where they annihilate and which is what we actually observe. ",
        "conclusions": "INTEGRAL/SPI has recently detected positron annihilation in the Galactic center region, which is compatible in size with the Galactic bulge, and whose rate is of the order of 1.3$\\times$10$^{43}$ e$^+$ s$^{-1}$. Prompted by the claim that those positrons could be produced by the annihilation of a new kind of light dark matter particles in the Galactic center region, we have studied astrophysical candidate sources which could be capable of injecting the observed amounts of positrons in the Galactic center region. We have ruled out SNe~Ia as the dominant positron injectors, due to their too low explosion rate in the Galactic bulge. We have proposed an alternative solution, namely hypernovae, capable of injecting up to 25 times more positrons than a typical SN~Ia. Hypernovae are likely to occur in the Galactic center region, from which the produced positrons could escape and fill up the entire Galactic bulge. The rate of hypernova events remains uncertain, but first crude estimates show that it might also be too low to be compatible with the steady-state positron production/annihilation hypothesis. The observed positrons could therefore be the remains of a starburst in the Galactic center region, which occurred a few million years ago. Before excluding hypernovae as possible positron injectors in the Galactic center region, more observational constraints are required."
    },
    "0404/astro-ph0404201_arXiv.txt": {
        "abstract": "There is mounting observational evidence that the expansion of our universe is undergoing an acceleration. A dark energy component has usually been invoked as the most feasible mechanism for the acceleration. However, it is desirable to explore alternative possibilities motivated by particle physics before adopting such an untested entity. In this work, we focus our attention on an acceleration mechanism: one arising from gravitational leakage into extra dimensions. We confront this scenario with high-$z$ type Ia supernovae compiled by Tonry et al. (2003) and recent measurements of the X-ray gas mass fractions in clusters of galaxies published by Allen et al. (2002,2003). A combination of the two databases gives at a 99\\% confidence level that $\\Omega_m=0.29^{+0.04}_{-0.02}$, $\\Omega_{rc}=0.21^{+0.08}_{-0.08}$, and $\\Omega_k=-0.36^{+0.31}_{-0.35}$, indicating a closed universe. We then constrain the model using the test of the turnaround redshift, $z_{q=0}$, at which the universe switches from deceleration to acceleration. We show that, in order to explain that acceleration happened earlier than $z_{q=0} = 0.6$ within the framework of gravitational leakage into extra dimensions, a low matter density, $\\Omega_m < 0.27$, or a closed universe is necessary. ",
        "introduction": "The recent well known distance measurements of distant type Ia supernovae (SNeIa) suggest an accelerating universe at large scales (Riess et al. 1998, Perlmutter et al. 1999, Tonry et al. 2003, Barris et al. 2004, Knop et al. 2003, Riess et al. 2004). The cosmic acceleration has also been confirmed, independently of the SNeIa magnitude-redshift relation, by the observations of the cosmic microwave background anisotropies (WMAP: Bennett et al. 2003) and the large scale structure in the distribution of galaxies (SDSS: Tegmark et al. 2003a,b). It is well known that all known types of matter with positive pressure generate attractive forces and decelerate the expansion of the universe. Given this, a dark energy component with negative pressure was generally suggested to be the invisible fuel that drives the current acceleration of the universe. There are a huge number of candidates for the dark energy component in the literature, such as a cosmological constant $\\Lambda$ (Carroll et al. 1992; Krauss and Turner 1995; Ostriker and Steinhardt 1995; Chiba and Yoshii 1999), a decaying vacuum energy density or a time varying $\\Lambda$-term (Ozer and Taha 1987; Vishwakarma 2001), an evolving scalar field (referred to by some as quintessence: Ratra and Peebles 1988; Caldwell et al. 1998; Wang and Lovelace 2001; Weller and Albrech 2002; Gong 2002; Li et al. 2002a,b; Chen and Ratra 2003; Mukherjee et al. 2003; Gong 2004), the phantom energy, in which the sum of the pressure and energy density is negative (Caldwell 2002; Dabrowski et al. 2003; Wang, Gong and Su 2004), the so-called ``X-matter\" (Turner and White 1997; Zhu 1998; Podariu and Ratra 2001; Zhu, Fujimoto and Tatsumi 2001; Alcaniz, Lima and Cunha 2003; Lima, Cunha and Alcaniz  2003; Feng, Wang and Zhang 2004; Dai, Liang and Xu 2004), the Chaplygin gas (Kamenshchik et al. 2001; Bento et al. 2002; Alam et al. 2003; Alcaniz, Jain and Dev 2003; Dev, Alcaniz and Jain 2003; Silva and Bertolami 2003; Makler et al. 2003), and the Cardassion model (Freese and Lewis 2002; Zhu and Fujimoto 2002, 2003; Sen and Sen 2003; Wang et al. 2003; Frith 2004; Gong and Duan 2004a,b).  However, the dark energy has so far no convincing direct laboratory evidence for its existence, so it is desirable to explore alternative possibilities motivated by particle physics before adopting such a component. In this respect the models that make use of the very ideas of branes and extra dimensions to obtain an accelerating universe are particularly interesting (Randall and Sundrum 1999a,b). Within the framework of these braneworld cosmologies, our observable universe is assumed to be a surface or a brane embedded in a higher dimensional bulk spacetime in which gravity could spread, and the bulk gravity sees its own curvature term on the brane which accelerates the universe without dark energy (Randall 2002). Recently, based on the model of Dvali et al. (2000) of brane-induced gravity, Deffayet and coworkers (Deffayet 2001, Deffayet, Dvali and Gabadadze 2002) proposed a scenario in which the observed late time acceleration of the expansion of the universe is caused by gravitational leakage into an extra dimension and the Friedmann equation is modified as follows \\begin{equation} \\label{eq:ansatz} H^2 = H_0^2 \\left[ \\Omega_k(1+z)^2+\\left(\\sqrt{\\Omega_{rc}}+ \\sqrt{\\Omega_{rc}+\\Omega_m (1+z)^3}\\right)^2 \\right] \\end{equation} where $H$ is the Hubble parameter as a function of redshift $z$ ($H_0$ is its value at the present), $\\Omega_k$, $\\Omega_{rc}$ and $\\Omega_m$ represent the fractional contribution of curvature, the bulk-induced term and the matter (both baryonic and nonbaryonic), respectively. $\\Omega_{rc}$ is defined as $\\Omega_{rc} \\equiv 1/4r_c^2H_0^2$, where $r_c$ is the crossover scale beyond which the gravitational force follows the 5-dimensional $1/r^3$ behavior. From a phenomenological standpoint, it is a testable scenario with the same number of parameters as a cosmological constant model, contrasting with models of quintessence that have an additional free function, the equation of state, to be determined (Deffayet et al. 2002). Such a possible mechanism for cosmic acceleration has triggered investigations aiming to constrain this scenario using various cosmological observations, such as SNeIa (Avelino and Martins 2002; Deffayet, Dvali and Gabadadze 2002; Deffayet et al. 2002; Dabrowski et al. 2004), angular size of compact radio sources (Alcaniz 2002), the age measurements of high-$z$ objects (Alcaniz, Jain and Dev 2002), the optical gravitational lensing surveys (Jain et al. 2002) and the large scale structures (Multam\\\"aki et al. 2003). But the results are disperse and somewhat controversial, with most of them claiming good agreement between data and the model while some of them ruling out gravitational leakage into an extra dimension as a feasible mechanism for cosmic acceleration.  The purpose of this work is to quantitatively confront the scenario with the updated SNeIa sample compiled by Tonry et al. (2003) and to try to constrain the model parameters more accurately. It is shown that, although the two parameters, $\\Omega_{rc}$ and $\\Omega_m$, are degenerate and there is a range on the parameter plane to be consistent with the SNeIa data, a closed universe is prefered by this scenario. As is well known, the measurement of the X-ray gas mass fraction in galaxy clusters is an efficient way to determine the matter density, $\\Omega_m$, and hence can be used for breaking the degeneracy between $\\Omega_{rc}$ and $\\Omega_m$. When we combine the X-ray database published by Allen et al. (2002, 2003) for analyzing, we obtained a closed universe at a 99\\% confidence level, i.e., for the scenario of gravitational leakage into an extra dimension, a universe with curvature is favored by the combination of the two databases. We also analyze the turnaround redshift, $z_{q=0}$, at which the universe switches from deceleration to acceleration within the framework of the scenario. It is shown that, if the turnaround redshift happened earlier than $z_{q=0} = 0.6$, only a low matter density, $\\Omega_m < 0.27$, or a closed universe can explain this transition epoch. If, however, we consider the recent estimate by Riess et al. (2004), i.e., $z_{q=0} = 0.46 \\pm 0.13$, then a spatially flat scenario with $\\Omega_m =0.3$ (as suggested by clustering estimates) predicts $z_{q=0} = 0.48$, which is surprisingly close to the central value given by Riess et al. (2004). The paper is organized as follows. In the next section, we consider the observational constraints on the parameter space of the scenario arising from the updated SNeIa sample compiled by Tonry et al. (2003), as well as the combination with the X-ray gas mass fractions in galaxy clusters published by Allen et al. (2002, 2003). In section~3 we discuss the bounds on the model from the turnaround redshift, $z_{q=0}$. Finally, we present our conclusion and discussion in section~4. ",
        "conclusions": "The mounting observational evidences for an accelerating universe have stimulated renewed interest for alternative cosmologies. Generally, a dark energy component with negative pressure is invoked to explain the SNeIa results and to reconcile the inflationary flatness prediction ($\\Omega_T = 1$) with the dynamical estimates of the quantity of matter in the universe ($\\Omega_m \\sim 0.3$). In this paper we have focused our attention on another possible acceleration mechanism, one arising from gravitational leakage into extra dimensions. In order to be consistent with the current SNeIa and the X-ray clusters data, one would need a closed universe.  Recently Lue et al. (2004) derived dynamical equations for spherical perturbations at subhorizon scales and computed the growth of large-scale structure in the framework of this scenario. A suppression of the growth of density and velocity perturbations was found, e.g., for $\\Omega_m=0.3$, a perturbation of $\\delta_i=3\\times 10^{-3}$ at $z_i=1000$ collapse in the $\\Lambda$CDM case at $z\\approx 0.66$ when its linearly extrapolated density contrast is $\\delta_c=1.689$, while for the model being considered the collapse happens much later at $z\\approx 0.35$ when its $\\delta_c=1.656$. Furthermore, the authors showed that this scenario for cosmic acceleration gave rise to a present day fluctuation power spectrum normalization $\\sigma_8 \\leq 0.8$ at a 2$\\sigma$ level, lower than observed value (Lue et al. 2004).  As is shown in Figure~2 of Deffayet, Dvali and Gabadadze (2002), on the assumption of a flat universe, luminosity distance for $\\Lambda$CDM increases with redshift faster than that for the model being considered does (for the same $\\Omega_m$). Therefore it is natural that, if the $\\Lambda$CDM model with ($\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$, $\\Omega_k=0$) is consistent with the SNeIa data, the gravitational leakage model with ($\\Omega_m=0.3$, $\\Omega_{rc}=0.1225$, $\\Omega_k=0$) will not be as the data are becoming enough to determine the cosmological parameters more precisely. While Deffayet et al. (2002) showed that the gravitational leakage scenario was consistent with the 54 SNeIa of the sample C from Perlmutter et al. (1999) -- see also Alcaniz \\& Pires (2004) -- Avelino and Martins (2002) claimed that this proposal was disfavored by the dataset of 92 SNeIa from Riess et al (1998) and Perlmutter et al. (1999) [combining them via the procedure described in Wang (2000) and Wang \\& Garnavich (2001)]. We, however, think that only with a more general analysis, a joint investigation involving different classes of cosmological tests, it will be possible to delimit the $\\Omega_{\\rm{m}} - \\Omega_{r_c}$ plane more precisely, as well as to test more properly the consistency of these senarios. Such an analysis will appear in a forthcoming communication (Alcaniz \\& Zhu 2004)."
    },
    "0404/astro-ph0404037_arXiv.txt": {
        "abstract": "This Letter presents a frequentist analysis of the hot and cold spots of the cosmic microwave background data collected by the \\emph{Wilkinson Microwave Anisotropy Probe} (WMAP).  We compare the WMAP temperature statistics of extrema (number of extrema, mean excursion, variance, skewness and kurtosis of the excursion) to Monte Carlo simulations.  We find that on average, the local maxima (high temperatures in the anisotropy) are too cold and the local minima are too warm. In order to quantify this claim we describe a two-sided statistical hypothesis test which we advocate for other investigations of the Gaussianity hypothesis. Using this test we reject the isotropic Gaussian hypothesis at more than 99\\% confidence in a well-defined way. Our claims are based only on regions that are outside the most conservative WMAP foreground mask. We perform our test separately on maxima and minima, and on the north and south ecliptic and Galactic hemispheres and reject Gaussianity at above 95\\% confidence for almost all tests of the mean excursions. The same test also shows the variance of the maxima and minima to be low in the ecliptic north (99\\% confidence) but consistent in the south; this effect is not as pronounced in the Galactic north and south hemispheres. ",
        "introduction": "The \\emph{Wilkinson Microwave Anisotropy Probe} (WMAP) data provide the most detailed data on the full sky cosmic microwave background (CMB) to date. This information about the initial density fluctuations in the universe allows us to test the cosmological standard model at unprecedented levels of detail. (Bennett {et~al.} 2003a) A question of fundamental importance to our understanding of the origins of these primordial seed perturbations is whether the CMB radiation is really an isotropic and Gaussian random field, as generic inflationary theories predict (Starobinsky 1982; Guth \\& Pi 1982; Bardeen {et~al.} 1983).  A natural way to study the CMB is to look at the local extrema.  This was initially suggested because the high signal-to-noise ratio at the hot spots means they would be detected first (Sazhin 1985; Zabotin \\& Naselsky 1985; Vittorio \\& Juszkiewicz 1987; Bond \\& Efstathiou 1987). Heavens \\& Sheth calculate analytically the two-point correlation function of the local extrema  (Heavens \\& Sheth 1999). In addition, extrema trace the topological properties of the temperature map; this makes them good candidates for study (Wandelt {et~al.} 1998).  We pursue this investigation by simulating Gaussian Monte Carlo CMB skies and comparing the WMAP data to those simulations. We choose several statistics and then check to see if the WMAP statistics lie in the middle of the Monte Carlo distributions of statistics. We present results on the one-point functions of the local extrema: their number, mean excursion, and variance, and skewness and kurtosis of the excursion.  The literature contains many other searches for non-Gaussianity, in the WMAP data and other CMB experiments. For example, Vielva et~al.\\ detect non-Gaussianity in the three- and four-point wavelet moments (Vielva et al.\\ 2004), Chiang et~al.\\ detect it in phase correlations between spherical harmonic coefficients (Chiang et al.\\ 2003; see also Chiang et al.\\ 2002, 2004), and Park finds it in the genus Minkowski functional  (Park 2004). Eriksen et~al.\\ find anisotropy in the $n$-point functions of the CMB in different patches of the sky  (Eriksen et al.\\ 2004). Others discuss possible methods of detecting non-Gaussianity. Aliaga et~al.\\ look at studying non-Gaussianity through spherical wavelets and ``smooth tests of goodness-of-fit''  (Aliaga et al.\\ 2003). Cabella et~al.\\ review three methods of studying non-Gaussianity: through Minkowski functionals, spherical wavelets, and the spherical harmonics  (Cabella et al.\\ 2004). They propose a way to combine these methods.  Komatsu et~al.\\ discuss a fast way to test the bispectrum for primordial non-Gaussianity in the CMB  (Komatsu 2003a), and do not detect it  (Komatsu et al.\\ 2003b). Finally, Gazta{\\~n}aga et al.\\ find the CMB to be consistent with Gaussianity when considering the two and three-point functions  (Gazta{\\~n}aga \\& Wagg 2003; Gazta{\\~n}aga et al.\\ 2003). To this work, we add a strong detection of non-Gaussianity based on generic features: the local extrema.  The Letter is laid out as follows. The next section discusses our method for making Monte Carlo simulations of the CMB sky and calculating statistics on both the simulations and the WMAP data.  It also explains our statistical tests. Section 3 describes our results. We conclude in section 4. ",
        "conclusions": "In this Letter we generate simulated CMB skies.  We choose several statistics, and calculate them on both the simulations and the WMAP sky. We hypothesize that the WMAP statistics are drawn from the same distribution as the simulations' statistics, since we have attempted to accurately simulate the CMB sky. If the WMAP statistic is higher or lower than most of the simulations' statistics, this indicates that the WMAP statistic's underlying position $p\\in[0,1]$ in the distribution of Monte Carlo statistics is close to $0$ or $1$. If we are 95\\% confident that $p$ is within 0.025 of $0$ or $1$, then we claim the probability of the WMAP statistic happening by chance is sufficiently small to reject the hypothesis.  We find the WMAP data to have maxima that are significantly colder and minima that are significantly warmer than predicted by Monte Carlo simulation.  For almost all simulations, we have 95\\% confidence that the mean of the WMAP hot spots or cold spots is in a 5\\% tail of the Monte Carlo distribution. In one case, we are 99\\% confident the maxima statistic is in a 1\\% tail. Since we find the same lack of extreme temperature when we use the directly measured WMAP power spectrum, we are not simply restating that the WMAP power spectrum has a lack of power at large angular scales. The effect is independent of the Galactic mask or power spectrum used.  We also find some anisotropy between the ecliptic north and south hemispheres. The WMAP data in northern hemisphere have a low variance statistic (95\\% confident that the variance statistic is in a 5\\% tail).   In one case, we are 99\\% confident the variance of the maxima is in a 1\\% tail. There is less asymmetry between the north and south Galactic hemispheres.  Our results may not be a detection of primordial non-Gaussianity. They could still be an effect of the WMAP instrument or data pipeline not modeled in our simulations or an as yet undiscovered foreground. Our result is still highly significant. We have detected something, whether it is primordial non-Gaussianity or some other effect in the data. Having anomalous mean temperature values for the maxima and minima in both the north and south ecliptic hemispheres is unlikely to occur if the WMAP data were consistent with theoretical expectations. We will present a complete treatment of the one- and two-point extrema statistics for the WMAP data set in a future publication."
    },
    "0404/astro-ph0404509_arXiv.txt": {
        "abstract": "There is now firm evidence that the ICM consists of a mixture of hot plasma, magnetic fields and relativistic particles. The most important evidences for non-thermal phenomena in galaxy clusters comes from the diffuse Mpc-scale synchrotron radio emission (radio halos) observed in a growing number of massive clusters (\\cite{Fer03}) and from hard X-ray (HXR) excess emission (detected in a few cases) which can be explained in terms of IC scattering of relativistic electrons off the cosmic microwave background photons (Fusco-Femiano et al. 2003). There are now growing evidences that giant radio halos may be naturally accounted for by synchrotron emission from relativistic electrons reaccelerated by some kind of turbulence generated in the cluster volume during merger events (\\cite{Bru03}). With the aim to investigate the connection between thermal and non-thermal properties of the ICM, we have developed a statistical magneto-turbulent model which describes the evolution of the thermal and non-thermal emission from clusters. We calculate the energy and spectrum of the magnetosonic waves generated during cluster mergers, the acceleration and evolution of relativistic electrons and thus the resulting synchrotron and inverse Compton spectra. Here we give a brief description of the main results, while a more detailed discussion will be presented in a forthcoming paper (Cassano \\& Brunetti, in preparation). Einstein-De Sitter cosmology, $H_o=50$ km $s^{-1}$$Mpc^{-1}$, $q_o=0.5$, is assumed. ",
        "introduction": "\\noindent Giovannini, Tordi and Feretti (1999) found that $\\sim$5\\% of clusters from a complete X-ray flux limited sample have a radio halo source. The detection rate of radio halos shows an abrupt increase with increasing the X-ray luminosity and mass of the host clusters: about 30-35\\% of the galaxy clusters with X-ray luminosity larger than 10$^{45}$ erg s$^{-1}$ show diffuse non-thermal radio emission (\\cite{Fer03}). Recent papers (\\cite{Ensslin03}; \\cite{Kuo03}) have investigated the statistics of the formation of radio halos from a more theoretical point of view. In these works, however, the expected statistics are not derived from formation models of radio halos, but simply from the observed luminosity-mass correlations and mass thresholds.\\\\ We model the formation of radio halos and HXR tails in a self--consistent approach which follows, at the same time, the evolution of the thermal properties of the ICM and the triggering and evolution of the non--thermal phenomena assuming magneto-turbulent re--acceleration of relativistic particles. In particular, we follow the formation and evolution of clusters of galaxies, the generation of merger-driven turbulence and magnetosonic waves in the cluster volume, the acceleration and time--evolution of the relativistic particles, and of the related non--thermal emission. ",
        "conclusions": ""
    },
    "0404/astro-ph0404023_arXiv.txt": {
        "abstract": "{The X-ray luminosities of the hot halo gas around simulated, Milky Way like disk galaxies have been determined, as a function of redshift. The X-ray luminosity increases significantly with redshift, in some cases as much as a factor 30 going from $z$=0 to 2. Consequently, the optimal detection redshift can be $\\ga$0.5.  Results of fully cosmological simulations of galaxy groups and clusters, incorporating star formation, chemical evolution with non-instantaneous recycling, metal-dependent radiative cooling, galactic super-winds and thermal conduction are presented. X-ray luminosities at $z$=0 are somewhat high and central entropies somewhat low compared to observations. This is likely a combined effect of chemical evolution and metal-dependent, radiative cooling. Central ICM abundance profiles are somewhat steep, and the observed level of ICM enrichment can only be reproduced with IMFs more top-heavy than the Salpeter IMF. In agreement with observations it is found that the iron in the ICM is in place already at $z\\sim$1. The [Si/Fe] of the ICM decreases with time, and increases slightly with radius at a given time. The cluster galaxies match the observed ``red sequence'' very well, and the metallicity of cluster galaxies increases with galaxy mass, as observed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404215_arXiv.txt": {
        "abstract": "Evidence for a large scale flow of low density gas onto the Cepheus A young stellar cluster is presented. Observations of K-band near-infrared and multi-transition CS and \\n2hp\\ millimeter line emission are shown in relation to a sub-millimeter map of the cool dust around the most embedded stars. The near-infrared emission is offset from the dust peak suggesting a shift in the location of star formation over the history of the core. The CS emission is concentrated toward the core center but \\n2hp\\ peaks in two main cores offset from the center, opposite to the chemistry observed in low mass cores. A starless core with strong CS but weak \\n2hp\\ emission is found toward the western edge of the region. The average CS(2--1) spectrum over the cluster forming core is asymmetrically self-absorbed suggesting infall. We analyze the large scale dynamics by applying a one-dimensional radiative transfer code to a model spherical core with constant temperature and linewidth, and a density profile measured from an archival $850~\\mu$m map of the region. The best fit model that matches the three CS profiles requires a low CS abundance in the core and an outer, infalling envelope with a low density and undepleted CS abundance. The integrated intensities of the two \\n2hp\\ lines is well matched with a constant \\n2hp\\ abundance. The envelope infall velocity is tightly constrained by the CS(2--1) asymmetry and is sub-sonic but the size of the infalling region is poorly determined. The picture of a high density center with depleted CS slowly accreting a low density outer envelope with normal CS abundance suggests that core growth occurs at least partially by the dissipation of turbulent support on large scales. ",
        "introduction": "}  Most stars, particularly massive stars, form in groups (e.g. \\cite{carpenter00} 2000). It is therefore essential to study cluster forming regions in order to understand more completely the way in which the majority of stars are formed. Isolated low mass star formation occurs via the nearly isothermal free-fall collapse of a dense molecular cloud core, followed by the evolutionary phases Class 0, I, II and III objects (e.g. \\cite{evans99} 1999). However, the applicability of this paradigm to the formation of massive stars is debated (\\cite{garay+lizano99} 1999): for example, massive stars begin burning hydrogen and reach the main sequence while still accreting matter from the surrounding protostellar envelope and they can also develop strong winds, both of which will strongly affect the physical conditions, structure and chemistry of their surroundings. Due to the shape of the IMF and the fact that they evolve faster, massive protostars are rarer (and therefore more distant on average) than low mass protostars. Consequently fewer Class 0 massive protostar counterparts have been studied in detail. It is only recently that catalogs of high-mass protostellar objects have been made (e.g. \\cite{sridharan02} 2002).  The molecular cloud core Cepheus A East (hereafter Cep A) is a nearby site of massive star formation (\\cite{sargent77} 1977) located in the Cepheus OB association at a distance of 725 pc (\\cite{blauuw59} 1959). The far-IR luminosity is 2.4$\\times 10^4~L_\\odot$ (\\cite{evans81} 1981), corresponding to a small cluster of B stars. Cep A harbors one of the first molecular bipolar outflow sources discovered (\\cite{rodriguez80} 1980). Higher spatial resolution CO observations showed the outflow to be extremely complex, and it was termed quadrupolar (\\cite{torrelles93} 1993). The fastest components of this outflow are bipolar and oriented northwest-southeast (\\cite{rodriguez80} 1980), perpendicular to the low velocity CO structure. The slower and more extended component has been interpreted as the diverting and redirecting of the main outflow by the interaction with interstellar high-density gas, seen in NH$_3$ lines by \\cite{torrelles93} (1993). Ultra compact \\ion{H}{ii} regions and a diffuse thermal dust emission source have been identified from 20~$\\mu$m maps and 6 cm low-resolution VLA observations by \\cite{beichman79} (1979). Seven ionized hydrogen complexes lie in ``strings'' that form a ``Y'' tilted to the east (\\cite{hughes+wouterloot84} 1984), the bifurcation point of which is coincident with the exciting source of the molecular outflow. A cluster of compact radio sources have been identified as pre-main-sequence stars by \\cite{hughes88} (1988) due to their variability and the presence of OH and H$_2$O maser emission. From subsequent ammonia VLA observations, \\cite{torrelles93} (1993) proposed that one of these radio sources, HW2, is a $\\sim$10--20~$M_\\odot$ protostar. The larger core surrounding the cluster has temperature 35 K and mass $200-300 M_\\odot$ \\cite{ms91} (1991). On the basis of its protostellar content, high luminosity and low temperature, and following the bolometric temperature definition of \\cite{chen95} (1995), the Cep A core may be considered a high mass Class 0 source.  In order to examine the properties of this young, massive cluster forming region, we obtained near-infrared and millimeter wavelength multi-transition CS and \\n2hp\\ data. The observations are discussed in \\S\\ref{data}. We derive the density profile from $850~\\mu$m continuum measurements and fit core-averaged spectra using a 1-D radiative transfer model in \\S\\ref{analysis}. Our results indicate large CS depletion in the central core and an outer undepleted accreting layer. We discuss these results and conclude in \\S \\ref{discussion}. ",
        "conclusions": "}  The line observations reveal a complex chemistry in the core. The CS emission is concentrated toward the center near the peak of the $850~\\mu$m dust emission and the youngest, most embedded protostars. However, the \\n2hp\\ map shows two prominent cores offset on either side of the dust peak. The presence of CS and absence of \\n2hp\\ toward the star forming center of the core is likely due to the fact that neutral molecules released from the dust grains in the hotter region surrounding the protostars preferentially destroy ions such as \\n2hp\\ (\\cite{bergin00} 2000).  The maps also reveal a small starless core toward the west of the main core. The relatively strong CS and weak \\n2hp\\ emission toward this core suggests that it has only recently formed (\\cite{bl97} 1997). \\cite{williams+myers99} (1999) found a starless core with similar chemical properties in the Serpens NW cluster  Despite the chemical complexities on small scales, the average CS(2--1) spectrum is asymmetrically self-absorbed suggesting large scale collapse. Using the Hogerheijde \\& van der Tak radiative transfer code, we fit the average CS and \\n2hp\\ spectra with a spherical model consisting of an inner region with a Plummer-like density profile measured from archival SCUBA $850~\\mu$m data with constant temperature and linewidth. The best fit model that matches the three CS profiles has a low CS abundance in the inner region and an outer, infalling envelope with a low density and higher CS abundance. The depletion toward the center matches chemical evolution model expectations (\\cite{bl97} 1997) and the envelope CS abundance is similar to that in the extended ridge of Orion (\\cite{vdb98} 1998). The fit also matched the integrated intensities of the two \\n2hp\\ spectra with a constant abundance similar to that found in B68 (\\cite{bergin02} 2002). In practice the maps show that the \\n2hp\\ must deplete toward the core center (see also \\cite{doty02} 2002) but we have not attempted to match the complicated small scale chemistry and our model fit for the \\n2hp\\ abundance should only be considered an average, weighted by column density, over the core.  The CS(2--1) self-absorption requires a large scale outside-in collapse. The velocity of the collapse could be accurately measured but the depth of the collapse region could not, due to the absence of self-absorption in the higher transition CS lines. At 0.22~\\kms, the infall velocity is sub-sonic for a gas temperature of 35~K. The mass infall rate of the envelope can be estimated from the ratio of its mass, $M_{\\rm env}=240~M_\\odot$ determined from its size and density, and the time for the outer edge to reach the core center, $$\\dot M_{\\rm in} = {M_{\\rm env}v_{\\rm in}\\over r_{\\rm env}+r_{\\rm core}} = 7\\times 10^{-5}~M_\\odot {\\rm yr}^{-1}.$$ This may only be a lower limit to the total mass infall rate at the center if the inner region is also collapsing. Nevertheless, the envelope mass infall rate alone is more than an order of magnitude higher than typical mass infall rates for solar mass protostars (\\cite{zhou95} 1995).  Our data do not rule out inside-out collapse motions around individual protostars at higher densities on smaller size scales since the $\\tau=1$ surface of the CS(2--1) emission occurs at low densities and therefore at large scales. The multitude of sources, powerful outflows, and the complex chemistry would likely make an investigation of the small scale motions around individual protostars quite challenging.  On large scales, however, our picture is of a core with a Plummer-like density profile accreting low density gas sub-sonically. The CS abundance in the infalling envelope is similar to undepleted values in the ISM. The dissipation of turbulent support resulting in ``cooling flows'' may lead to core growth in this manner (\\cite{nakano98} 1998 ; \\cite{myers+lazarian98} 1998 ; \\cite{williams+myers00} 2000). The growing availability of dust continuum maps and multi-transition, multi-species line observations will lead to more refined structural and dynamical modeling and comparisons between different star forming environments in the future."
    },
    "0404/astro-ph0404486_arXiv.txt": {
        "abstract": "We investigate chemical evolution in Milky Way-like galaxies based on the cold dark matter model in which cosmic structures form via hierarchical merging.  We introduce chemical enrichment due to type Ia supernovae (SNe Ia) into the Mitaka semi-analytic galaxy formation model developed by Nagashima \\& Yoshii.  For the first time we derive distributions of stellar metallicities and their ratios in Milky Way-like galaxies treating chemical enrichment due to SNe Ia in a hierarchical galaxy formation model self-consistently.  As a first attempt, we assume all SNe Ia to have the same lifetime, and assume instantaneous recycling for type II supernovae (SNe II).  We find that our model reproduces well the metal abundance ratio [O/Fe] against [Fe/H] and the { iron metallicity distribution function} in the solar neighborhood.  This means that the so-called G-dwarf problem is resolved by the hierarchical formation of galaxies, and a gas infall term introduced in traditional monolithic collapse models to solve this problem is well explained by the mixture of some physical processes such as hierarchical merging of dark halos, gas cooling, energy feedback and injection of gas and metals into hot gas due to supernovae.  Our model predicts more oxygen-enhanced stars in bulges at [Fe/H] $\\simeq 0$ than in disks.  This trend seems to be supported by recent observations while they have still uncertainties.  More data in number and accuracy will provide independent and important constraints on galaxy formation.  For the better understanding of the chemical enrichment due to SNe Ia in hierarchical galaxy formation, we discuss how physical processes affect the metal abundance ratio by varying the lifetime of SNe Ia, the star formation timescale and the strength of supernova feedback.  We find that the supernova feedback plays a key role among them and therefore there is no one-to-one correspondence between star formation histories and stellar metallicity-ratio distributions. ",
        "introduction": "The chemical compositions of stars and gas, as is widely known, provide important clues to understanding the formation and evolution of galaxies.  In particular the abundance ratio of $\\alpha$-elements such as oxygen to iron is a useful probe of the formation history of galaxies.  The main sources of metals are considered to be supernovae (SNe), especially type II SNe (SNe II) and type Ia SNe (SNe Ia).  They have different abundance patterns, the former produces $\\alpha$-elements as well as iron and the latter produces hardly any $\\alpha$-elements. Therefore, the measurement of [$\\alpha$/Fe] is closely related to the ratio of SNe II to SNe Ia.  While the progenitors of SNe II are massive stars larger than $\\sim 10M_{\\odot}$, those of SNe Ia are considered to be white dwarfs in binary systems.  Therefore the explosion rate of SNe II is almost proportional to the star formation rate of their progenitor stars at that time, but there is a time lag between SNe Ia explosions and the formation of their progenitor stars.  Since the evolution timescales of those elements are independent and different, we can obtain useful information on galaxy formation and star formation histories from the abundance ratio of those elements.  Because SNe II explode almost instantaneously, the chemical evolution process has usually been formulated by assuming that massive stars immediately explode right after their formation (the instantaneous recycling [IR] approximation)\\citep{t80, lf83, lf85}.  In those works, galaxies are usually assumed to grow up in monolithic clouds, in which timescales of gas infall and star formation (SF) are typically treated as free parameters.  Since it is possible to observe and resolve individual stars in the solar-neighborhood, many studies have focused on evolution of the galactic disk.  By adjusting those parameters and by incorporating stellar population synthesis techniques \\citep{ay86, ay87}, detailed galaxy evolution models that reproduce observations have been developed \\citep{ayt92}.  In the framework of such traditional galaxy evolution models, several authors explicitly considered chemical enrichment due to SNe Ia \\citep{gr83, mg86, mf89, tnyhyt95, pt95, ytn96}.  They included delayed production of metals due to SNe Ia in the galaxy evolution models.  By comparing their models with solar neighborhood stars in { the metallicity distribution function (MDF)} and abundance ratio such as [O/Fe] against metallicity, \\citet{pt95} and \\citet{ytn96} inferred a typical lifetime of SNe Ia of 1.3$\\sim$1.5 Gyr.  Although \\citet{ktnhk98} has claimed that SNe Ia start to explode at the age of about 0.6 Gyr, \\citet{gm04} recently derived that the lifetime should be longer than about 1 Gyr at the 99 per cent confidence level by comparing the cosmic star formation history with the cosmic SNe Ia rate.   Considered from the viewpoint of cosmological structure formation, however, these models are still only phenomenological.  Recent theoretical and observational studies of structure formation have revealed that the universe is dominated in mass by cold dark matter (CDM) and that baryonic objects such as galaxies form in virialized objects of dark matter called dark halos.  In the CDM model, since smaller-scale density fluctuations in the early universe have larger fluctuation amplitudes, larger halos form through mergers of smaller dark halos (the hierarchical clustering scenario).  { This suggests that large galactic gas clouds would have formed at low redshift by mergers of subgalactic clumps which have formed at higher redshift. Therefore} we may have to modify our picture of chemical evolution in galaxies as well as that of galaxies themselves.  Recently, based on a semi-analytic (SA) approach, galaxy formation models in the hierarchical clustering scenario have been developed, in which the formation histories of individual dark halos are followed by using a Monte Carlo method and physical processes such as gas cooling and star formation are taken into account in the histories of dark halos.  This is a natural consequence of adopting the CDM model because we are not able to freely assume formation histories of dark halos as far as we would like to construct galaxy formation models consistent with the CDM model.  Many authors have found that such SA models reproduce well various characteristics of galaxies at the present and at high redshift such as luminosity functions, gas fractions, size distributions, and faint galaxy number counts \\citep[e.g.,][]{kwg93, cafnz94, clbf00, ngs99, ntgy01, nytg02, ny04, sp99, spf01}.  In recent SA models, chemical enrichment is considered, but generally only with the IR approximation.  \\citet{kc98} and \\citet{ng01} investigated color-magnitude and metallicity-magnitude relations of cluster elliptical galaxies.  This is extended to dwarf spheroidals, $M_{B}\\sim -10$, by \\citet{ny04}.  \\citet{k96}, \\citet{spf01} and \\citet{ongy04} consider chemical evolution in spiral galaxies, a part of which should be identified as damped Ly-$\\alpha$ systems.  Some of them found good agreement with observations.  Pioneering work taking into account chemical enrichment due to SNe Ia in a SA model was carried out by \\citet{t99} and \\citet{tk99}.  In those papers, they picked out averaged and individual formation histories of dark halos, and then, assuming the closed-box chemical evolution, they followed star formation and chemical enrichment histories.  Thus, while they took into account merging histories of galaxies, the outflow of gas caused by SNe or the supernova (SN) feedback that has been realized to be an primarily important process in galaxy formation, was not considered.  As shown by \\citet{kc98} and \\citet{ng01}, SN feedback significantly affects chemical enrichment, at least due to SNe II.  This suggests that the closed-box model has limitations in analyzing chemical evolution in a realistic situation.  As a complementary approach to SA modeling, chemo-dynamical simulations including SNe Ia have been developed \\citep[e.g.,][]{rvn96, b99, k01, lpc02, nm03, k03, oefj05}.  Although this approach has an advantage in resolving spatial structure, because of the limitation of numerical resolution and numerical techniques themselves, further improvements are still required \\citep{ojeqf03}.  On the other hand, the SA model is free from numerical effects and limitations.  This is a great advantage for understanding global properties of galaxy formation because in principle the SA model can investigate from dwarf galaxies to galaxy clusters simultaneously.  In this paper, we construct a fully self-consistent treatment of chemical enrichment with a SA model.  The basic model is the Mitaka model presented by \\citet{ny04}, which includes a Monte Carlo realization of the merging histories of dark halos, radiative gas cooling, quiescent star formation and starbursts, mergers of galaxies, chemical enrichment assuming the IR approximation, size estimation of galaxies taking into account the dynamical response to gas removal during starbursts and stellar population synthesis.  As a first attempt, we assume that all SNe Ia have the same lifetime.  This is, of course, a rather simplified model.  In reality the lifetimes of SNe Ia are considered to have a broad distribution, and there is even a claim that low-metallicity environments inhibit SNe Ia \\citep{ktnhk98}.  Our model, however, has an essential characteristic of SNe Ia, that is, the time lag between star formation and the explosion.  It enables us to see how such a delayed explosion of SNe Ia affects the chemical enrichment and abundance pattern in galaxies.  Using this model, we focus on the chemical enrichment in Milky Way (MW)-like galaxies.  This has a particularly meaning because there are many non-trivial effects on the chemical enrichment in the hierarchical formation of galaxies.  Massive galaxies such as the MW form not only via gas cooling but via mergers of pre-galactic sub-clumps with stronger efficiencies of the SN feedback than those in massive galaxies residing in deeper gravitational potential wells.  { Thus, the main purpose of this paper is to compare the SA model with a monolithic collapse model of \\citet{ytn96}. This will enable us to see how the hierarchical formation process affects galactic metal enrichment due to SNe Ia based on the CDM model.} Further comparison to other objects will be done in subsequent papers.  This paper is outlined as follows.  In \\S2 we describe our SA model.  In \\S3 we provide a detailed prescription of the chemical enrichment due to SNe Ia.  In \\S4 we show the luminosity function of galaxies that should be compared with that in the Local Group, and properties of MW-like galaxies.  In \\S5 we compare our model with observations in the [O/Fe]--[Fe/H] plane and the [Fe/H] distribution in a statistical sense. In \\S6 we briefly discuss individual galaxies.  In \\S7 we investigate parameter dependences of the main results.  In \\S8 we provide a summary and conclusions. ",
        "conclusions": "We have explored chemical enrichment due to both SNe II and SNe Ia in Milky Way-like galaxies in the semi-analytic galaxy formation model. Our treatment of SNe Ia is fully consistent with the galaxy formation model, that is, we solve for the recycling of each element among stars, cold gas and hot gas based on a $\\Lambda$CDM model.  It is important to follow the chemical enrichment in the framework of the hierarchical galaxy formation because there are many non-trivial effects on the chemical enrichment caused not only by mergers of galaxies but by varying efficiencies of the SN feedback dependent on the depth of gravitational potential wells of sub-galactic clumps.  As a first attempt at constructing such a consistent model, we have assumed that all SNe Ia have the same lifetime, $t_{\\rm Ia}$, that abundance patterns of metals from SNe Ia and SNe II are always the same, and that massive stars instantaneously explode as SNe II and release metals.  This is a natural extension of the work by \\citet{ytn96}, in which they self-consistently modeled the chemical enrichment due to SNe Ia and SNe II.  We have picked out galaxies in dark halos with $V_{\\rm circ}=220$ km~s$^{-1}$ and having a similar luminosity to the MW.  We have found that when we impose $t_{\\rm Ia}$=1.5 Gyr, the predictions of our model for such MW-like galaxies agree well with observations of the chemical composition of solar-neighborhood stars both in the stellar distribution in the [O/Fe]-[Fe/H] plane and in the { iron MDF}.  We would like to stress that the other model parameters such as the star formation timescale and SN feedback are the same as in the fiducial model of \\citet{ny04}, in which they found that the model reproduce well many aspects of observed galaxies, such as luminosity functions, cold gas fractions and sizes of galaxies for local galaxies, and the surface brightnesses, velocity dispersions, mass-to-light ratios and metallicities of local dwarf spheroidals.  This work also shows that the classical G-dwarf problem \\citep{v62, s63, pp75} is fully resolved in the framework of hierarchical formation of galaxies, in which the infall term introduced in the infall models to avoid the G-dwarf problem is naturally explained by the mixture of such physical processes as clustering of dark halos, gas cooling and SN feedback.  Our model passes the new tests, that is, { the iron MDF} and the abundance pattern of metals, taking into account chemical enrichment due to SNe Ia.  As shown in Section \\ref{sec:depend}, the abundance pattern provides an independent constraint on galaxy formation.  Therefore our results would support the scenario of hierarchical formation of galaxies.  While observational data used in this work have been limited to solar-neighborhood and bulge stars in the Milky Way, future observations of stars in other galaxies will provide statistically better constraints on galaxy formation.  In our model, disk and bulge stars are treated separately.  While the observations of bulge stars still have uncertainties, our model also reproduces well both the distribution of stars in the [O/Fe]--[Fe/H] plane and the [Fe/H] distribution.  In particular, although there are only a few bulge stars whose oxygen abundance is observed, our model predicts oxygen-enhanced stars at [Fe/H]$\\sim 0$, but they are not a dominant fraction.  Increasing observational data will provide a strong constraint on bulge formation.  As a by-product, we have obtained the luminosity function of galaxies in Local Group-halos.  Recent high-resolution $N$-body simulations have predicted many more dwarf scale dark halos than observed satellite galaxies, and this has been considered to be a serious problem \\citep{kkvp99, m99}.  Such an overabundance problem has been solved in the framework of the SA models taking into account effects of reionization and incompleteness due to significantly low surface brightness \\citep{s02, bflbc02b}.  Our model also shows an ability to solve this problem in a similar manner to those works.  The chemical yields we used are the same as those in \\citet{ytn96}, in which an infall model of monolithic collapse was used \\citep{ayt92}. The similarity between our hierarchical model and the monolithic cloud collapse model suggests that spiral galaxies in a hierarchical universe should have a similar formation history to that modeled by the monolithic cloud collapse model, as shown in \\citet{bcf96}.  Star formation histories we computed also suggests the similarities.  It does, however, not always mean that similar star formation histories provide similar chemical enrichment histories.  As shown in Section 6.3 and more directly by equation (\\ref{eqn:Zbeta}), SN feedback plays a critical role in chemical enrichment.  In this paper we have concentrated on investigating the [O/Fe] relation and the [Fe/H] distribution.  Since recent analyses of solar-neighborhood stars have revealed that the age-metallicity relation has a large scatter, it will be useful as a next step to see whether our model can reproduce this relation as well as the scatter in it \\citep{ia02}.  Furthermore, while the analysis in this paper focused only on MW-like galaxies, our SA model has the potential ability to investigate other systems simultaneously.  To investigate abundance ratios in such systems as the intracluster medium, the stars composing elliptical galaxies, and damped Ly-$\\alpha$ systems will provide independent, important clues to understanding galaxy formation.  The cosmic explosion rate of SNe Ia will also give a new insight into both galaxy formation and observational cosmology.  We will pursue these topics in future papers."
    },
    "0404/astro-ph0404353_arXiv.txt": {
        "abstract": "According to the Cold Dark Matter (CDM) hierarchical clustering theory of galaxy and large scale structure formation, there should be numerous low mass dark matter haloes present in the Universe today. If these haloes contain sufficient stars they should be detectable as low luminosity stellar systems or dwarf galaxies. We have previously described a new detection method for faint low surface brightness objects and shown that there are relatively large numbers of very faint dwarf galaxies in the nearby Virgo cluster. In this paper we present results from a similar survey carried out on the Millennium Galaxy strip which runs along the celestial equator and samples a very different galaxy environment.  We show that the dwarf-to-giant galaxy number ratio along this strip ranges from 0.7:1 to, at most, 6:1, corresponding to a flat luminosity function ($\\alpha \\approx -0.8$ to $-1.0$). This is very different to our value of 20:1 for the Virgo cluster. There is no population of low surface brightness dwarf galaxies in the field that have gone undetected by the redshift surveys. This result is exactly opposite to what CDM models predict for the environmental dependence of the dark matter mass function which is that there are proportionally more small dark matter haloes in lower density environments. ",
        "introduction": "Data from the recent large redshift surveys carried out by SLOAN and 2dF have been used to define the global (averaged over all environments) Luminosity Function (LF) of galaxies (Blanton et al. 2001, Norberg et al. 2002). These two surveys produce a consistent result for the faint-end slope of the LF, $\\alpha \\approx -1.2$. This value is somewhat flatter than typically predicted by most Cold Dark Matter (CDM) models of large scale structure and galaxy formation unless some form of dwarf galaxy formation suppression is invoked (Mathis et al. 2002, Cole et al. 2000). A challenge for the numerical modellers is the observed environmental dependence of the relative dwarf galaxy numbers discussed in this paper.  Dwarf galaxies have been found in large numbers in a variety of rich, high density environments (Virgo cluster: Binggeli et al. 1984, Coma cluster: Milne $\\&$ Pritchet 2002, Fornax cluster: Kambas et al. 2000) but the evidence is growing that the large number of dwarfs predicted by standard CDM theory \\footnote {By standard CDM we mean a model that does not invoke dwarf galaxy suppression mechanisms, as discussed later in this introduction.} (mass (luminosity ?) function faint-end slope $\\alpha \\approx -2$) fail to appear in lower density environments. According to the standard CDM model, dwarf galaxies form when initial Gaussian density fluctuations in the primeval Universe grow linearly, collapse and virialize to produce what we see as dwarf galaxies. Simulations and semi-analytic models have been looked at to see what predictions CDM theory makes about the local dwarf galaxy population. For example Kauffmann et al.(1993) used semi-analytic models to look at the formation of galaxies within this hierarchical clustering theory (see also Mathis et al. 2002, Cole et al. 2000). Using a standard CDM model  they looked at both a dark matter halo with a circular velocity, $V_{circ}$ $\\approx$ 200  km s$^{-1}$ and compared its LF to observations of the Milky Way (MW), and also a dark matter halo with $V_{circ}$ $\\approx$ 1000  km s$^{-1}$ and compared this LF  to observations of the Virgo cluster. From their model of the MW sized halo, their calculations predicted 5-10 times more faint, low mass galaxies than observation showed.   Moore et al. (1999) have also conducted numerical simulations of CDM hierarchical galaxy formation to compare predictions with observations of the MW and Virgo cluster. The  circular velocity (mass) distribution of the haloes they simulated for both the MW and Virgo cluster were very similar, differing only by the scaling factor of the halo mass, though the cluster halo was 2500 times larger than the galaxy halo and formed 5 Gyrs later. They found that their simulations agreed well with Virgo cluster observations - a plot of the abundance of haloes as a function of their circular velocity showed that the simulated and observed Virgo cluster numbers were very similar. However, the simulated galaxy haloes, when compared to that of the Local Group (LG) dwarf galaxies, overpredicted the total number of satellites larger than dSphs by a factor of about 50 (see also Klypin et al. 1999).  Although the above papers highlight the discrepancy between simulation and observation we should be careful with this comparison. In the main the simulations are of dark matter haloes and it is these that are overproduced in the simulations. To relate dark matter haloes to observations of luminous galaxies requires some modelling of the way in which baryonic material falls into the dark halo and how it is subsequently converted into stars. These physical processes are not so straightforward to model as those used in a standard CDM simulation. Attempts to make the observations and predictions match up include suppressing the formation of dwarf galaxies with a photoionizing background (Efstathiou 1992, Dijkstra et al. 2004), inhibiting star formation by expelling gas, a 'feedback' mechanism (Dekel \\& Silk, 1986) and merging the fainter galaxies so their number decreased. Kauffmann et al.(1993) concluded that it was very difficult to suppress the formation of so many dwarf galaxies compared to observation - this is often referred to as the sub-structure problem. Other possible solutions to the sub-structure problem that do not fit so well within the standard CDM model are that the initial power spectrum is wrong (Kamionkowski $\\&$ Liddle, 2000), baryonic material does not fall into small haloes - they remain dark (Bullock et al. 2000), baryonic material falls in, but fails to form stars or stars do form, but there are so few they have so far failed to be detected. It is the last of these solutions that we intend to investigate as part of the work described in this paper. Our motivation is that recent determinations of the field galaxy luminosity function (for example 2dF, see Norberg et al. 2002) have relied upon data obtained from photographic plates that are only sensitive to relatively high surface brightness objects (isophotal limit of $\\approx 24.5$ B$\\mu$). In the LG and in nearby clusters there is a well defined surface brightness magnitude relation (Ferguson $\\&$ Binggeli, 1994) such that low luminosity objects also have low surface brightness - they are doubly cursed. Photographic surveys would miss many of these faint low surface brightness (LSB) dwarf galaxies and even if detected, it is then very difficult to obtain redshifts, even with the largest telescopes. Thus potentially there may be many dwarf galaxies missing, due to selection effects, from the data used to derive the LF. This issue has also been discussed extensively by Cross et al. (2001), Cross \\& Driver (2002) and Liske et al. (2003). What we bring new to this discussion is a detection algorithm that is optimised to find LSB dwarf galaxies and a direct comparison with surveys sampling the galaxy population in different environments. So, our second motivation is that there appears to be a strong environmental affect on the relative numbers of dwarf compared to giant galaxies. How can CDM and its associated dwarf galaxy formation suppression mechanisms explain this?  A further important point is that the large redshift surveys have only accurately measured the LF for $M_{B}<-17$ (Driver \\& de Propis, 2003). It is not at all clear whether the extrapolation of the LF to fainter magnitudes is valid. The only environment where the LF appears to be well measured fainter than $M_{B}=-17$ is the Local Group (Mateo, 1998, Pritchet \\& van der Bergh, 1999) and this gives a flat faint-end slope ($\\alpha=-1.1$) down to the faint magnitudes ($M_{B}=-10$) we explore in this paper.  Various other surveys have previously been carried out to quantify the population of dwarf galaxies in different environments (Trentham $\\&$ Tully 2002, Trentham 1997, Chiboucas $\\&$ Mateo 2001). These studies usually take the form of finding the faint end slope of the LF (described by a Schechter function) for a sample of galaxies in some field, group or cluster environment. Comparing surveys is very difficult because they are often in different bands and have different magnitude and surface brightness limits. For example, Trentham $\\&$ Hodgkin (2002) find the B-band faint-end slope of the LF of the Virgo cluster to be $\\approx$ -1.4 for galaxies fainter than $M_{B} = -18$, and compare it to the value obtained by Phillipps et al (1998) who found a steeper value of -2.2 in the R-band, using a very different method to identify cluster galaxies. In their paper, Trentham $\\&$ Hodgkin also comment on the shallow LF obtained for the Ursa Major cluster (Trentham et al. 2001), but their data for the 2 clusters was obtained using different instruments and different filters. The method of selecting galaxies is also carried out in different ways for different surveys. Of particular concern is deciding which galaxies are cluster members and which are background, redshifts being difficult to obtain for faint LSB objects. Trentham et al. (2001) in their study of the UMa cluster find a condition for membership of the cluster based on measured light concentrations of the galaxies. They use the magnitude vs. central surface brightness relation of Ferguson $\\&$ Binggeli (1994) and say that for a given apparent magnitude, the concentration of light for cluster dwarf galaxies will be less concentrated than for background galaxies of the same apparent magnitude due to the dwarf's lower surface brightness and larger sizes. Trentham et al. state, that any dwarf galaxies which satisfy both these criteria are possible cluster members, although there is some contamination from background objects (see their paper for further details). They give no independent demonstration that their selection method works. Phillipps et al. (1998) use an entirely different method. They subtract galaxy counts obtained from fields outside of the cluster away from those inside the cluster to be left with the residual (small) cluster contribution. These methods have consistently led to luminosity functions much steeper than those derived by other methods. It is not difficult to see why - the 'clumpiness' of the background and the subtraction of one large number from another to leave a small residual. If the background count slope is $0.6m$ and this remains in the residual then the inferred luminosity function faint-end slope would be a very steep -2.5 (see also Valotto et al. 2001). In our previous work (Sabatini et al. 2003) we demonstrate (decreasing number density with distance from the cluster centre) that with the correct selection criteria we are able to preferentially select cluster dwarf galaxies.  To be able to make proper comparisons of the LFs in different environments, all variables (e.g. instrument, band, exposure times, selection criteria) should ideally be the same. This is what we have tried to do with the three 'environments' described in this paper. Our 3 surveys were conducted using the same instrument, technique (filter band, exposure time), and selection criteria. We can be confident therefore, that, unlike similar studies, we are comparing 'like with like'. Throughout this paper we use $H_{0}=75$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "As stated above, only 4 (5) of the objects with redshifts lie within 21 Mpc. Roughly accounting for those objects without redshifts we can have no more than 6 (18) objects within 21 Mpc in total. In section 3.1 we described a model of the numbers expected for various LF faint-end slopes. Given the observed numbers per sq deg we would have expected a LF faint-end slope of about -1.4 and so approximately 45\\% of our detections were predicted to lie within 21 Mpc. This corresponds to 23 (50) objects. This discrepency leaves us with a bit of a dilemma. As stated in section 4.1, the volume sampled by the MGS to 21 Mpc is over dense in bright ($M_{B}<-19$) galaxies by a factor of 20, yet it is certainly under dense in dwarfs ($-14<M_{B}<-10$) compared to the model expectations.  \\begin{figure*} \\begin{minipage}{12cm} \\psfig{file=test_dist.ps,width=12cm,height=8cm} \\caption{\\small{The distribution of measured distances.}} \\label{dist} \\end{minipage} \\end{figure*}  The explanation appears to be two-fold. Firstly there are a number of very high redshift objects that the model does not predict should be there. The model does not take account of galaxy evolution. The second reason is that galaxies are not distributed uniformly in the Universe. In Fig. \\ref{dist}  we show the distribution of observed galaxy distances, this can be directly compared with the model predictions of Fig. 2. Although the predicted peak at about 21 Mpc can clearly be seen, there is also an excess of galaxies at distances greater than 70 Mpc. The model has been useful in that it enabled us to clearly specify the problem and to define the consequences of our selection criteria, but now we have the distances to so many objects it is not required for the interpretation of the data.  In the previous section we compared our result to that of the Virgo and Ursa Major clusters. In this section we want to compare with the predictions of CDM galaxy formation models.  To do this we want to be as optimistic as possible about the numbers of objects we might be missing. There are 6 bright galaxies ($M_{B}<-19$) within 21 Mpc. We have 4 dwarf galaxy ($-14<M_{B}<-10$) detections (12, 13, 31, 33) within 21 Mpc giving a DGR of 0.7. Including the possible detection of object 48 of Table \\ref{unsure} increases this to 0.8. Adjusting now for the possibility that some of the unobserved and undetected objects lie within 21 Mpc produces a further 2 possible objects in Table \\ref{sure} and 11 in Table \\ref{unsure}, giving a DGR of 3. Now in section 3.1 we said that we can detect all galaxies with $-14<M_{B}<-10$ within 21 Mpc. This is actually only true if they follow the Driver, (1999) surface brightness relation. At fainter magnitudes some galaxies of higher surface brightness will be missed because they are too small. The volumes over which dwarf galaxies can be detected compared to the volume out to 21 Mpc are listed in Table \\ref{tab9}, this is the visibility function. \\footnote{Note that this does not affect our comparison with the Virgo data because both are observed over a similar depth.}  \\begin{table*} \\begin{minipage}{115mm} \\caption{Relative volumes, expressed as a percentage, that galaxies of different surface brightnesses ($\\mu_{0}$) and magnitudes ($M_{B}$) can be detected within - the visibility function.} \\label{tab9} \\begin{center} \\begin{tabular}{|c|c|c|c|c|c|c|} {\\bf $\\mu_{0}$}  &       &     &     & {\\bf $M_{B}$} &     &      \\\\ \\hline &             & -10 & -11 & -12     & -13 & -14  \\\\ \\hline \\hline 26 &   &  99   &   100  & 100        & 100    & 100    \\\\ \\hline 25  & & 25   & 99     & 100         & 100    & 100  \\\\ \\hline 24 &  &  6   & 25    &  99       &  100   &  100 \\\\ \\hline 23 &   &  2   &  6   &   25      &  99   &  100  \\\\ \\end{tabular} \\end{center} \\end{minipage} \\end{table*}  As can be seen for higher surface brightnesses and fainter magnitudes we do not sample the whole volume - the objects are too small at larger distances. Our observations do not rule out a population of faint galaxies of higher surface brightness in the field or in the Virgo cluster. Is there any evidence of such a population? Our first comment is that given the sparse numbers of detections for those magnitudes and surface brightnesses that we have full volume coverage, the LF would have to do something very strange if the numbers predicted by CDM are to be accounted for. In the Local Group there are 10 galaxies that satisfy our magnitude and surface brightness selection criteria. Of these, half lie in the region where we do not have full volume coverage. If this was also true for the MGS region then the DGR would at most double from 3 to 6. Observations by Deady et al. 2002 have been specifically made to try and identify higher surface brightness dwarfs in both the Fornax (Deady et al., 2002) and Virgo clusters (Drinkwater, private communications). In Virgo this amounts to about 3 per sq degree or about an additional 15\\% of our original total number. We conclude that there is no large population of higher surface brightness dwarf galaxies that have been missed in the MGS data and that at most the DGR is 6.   There is not a large population of faint field LSB dwarf galaxies that have been missed by the redshift surveys (see also Cross et al. 2001). We have measured the local LF down to three magnitudes fainter than the major redshift surveys, which produce LFs that are accurate over the range $M_{B}<-17$ (Driver \\& de Propris, 2003). With a DGR of, at most 6, the LF is flatter or declining ($\\alpha> -1$) compared to an extrapolation of the redshift surveys measured faint-end slopes. Within the CDM paradigm the suppression of star formation in field dwarf galaxies has been extremely efficient.  The observed environmental effect on both the numbers of dwarf galaxies and their relative number compared to bright galaxies in the field and cluster is completely opposite to that predicted by CDM for dark matter haloes. Lemson and Kauffmann, (1999) specifically consider the environmental influences on dark matter haloes and their associated galaxies. They conclude that the halo mass function (LF ?) '{\\it{is skewed towards high-mass objects in overdense regions of the Universe and towards low-mass objects in underdense regions'}}. Thus the CDM simulations predict that the ratio of low to high mass objects in the field should be higher than in clusters, completely opposite to what is observed. However, we must be careful with this comparison. CDM predicts how many dark haloes there should be - this should not be confused with the number of faint galaxies searched for in our surveys. Nevertheless, if these haloes do contains stars, thus making them visible as dwarf galaxies, then a mechanism must be used to suppress their formation in the field in order to reconcile their predicted numbers with observations - this is often referred to in the simulations as a 'feedback' mechanism. The normal 'feedback' mechanism invoked in most models is to expel gas from small dark matter potentials by the injection of energy by the first supernovae. This suppresses the formation of stars in these haloes and they remain undetected (dark). This should apply equally in all environments (Virgo, UMa and the MGS) suppressing the formation of dwarf galaxies everywhere. A possible solution is that the intra-cluster gas in environments like Virgo prevents the gas escaping (Babul \\& Rees, 1992). This would only apply to galaxies within the cluster core where the gas density is relatively high, but within the core, dwarf galaxies are subject to tidal destruction (resulting in intra-cluster stars, planetary nebulae and inter-galactic light) Sabatini et al. 2004. Ram pressure stripping is again only effective in the cluster core and suppresses rather than enhances star formation. Tully et al. (2002) have proposed that the environmental dependence is due to the time at which larger scale structures form in relation to the epoch of re-ionisation. They propose that the dwarf galaxy population of Virgo formed early, before reionisation, and was able to retain gas and form stars. In the lower density environments (UMa, MGS) the dark matter haloes form later, after re-ionisation, and the gas is too hot to collapse. Tully et al. say that there is only ``qualitative'' agreement between their idea and observations. Their argument is further weakened by the recent result by the WMAP team that places the epoch of re-ionisation at a much more distant redshift of $z \\approx 20$ (Spergel et al. 2003) \\newline A test for the existence of dark haloes (DM haloes with no stellar systems) would be to use gravitational lensing as a probe of substructure. This is an ideal tool to use since light is deflected gravitationally by matter, whether it is light or dark, thus if there were small dark haloes present in the Universe, they could be detected by this means. Such studies have been carried out (Metcalf $\\&$ Zhao, 2002, Bradac et al. 2002) and preliminary results show evidence for the presence of substructure. Dalal $\\&$ Kochanek (2002) studied seven four-image lens systems, six of which had flux anomalies which they commented could be due to the effects of substructure. They also rule out the possibilities of other affects causing the flux anomalies in a further study of their data (Kochanek $\\&$ Dalal, 2003), concluding that '{\\it{low mass haloes remain the best explanation of the phenomenon'}}. However, if these low mass DM haloes do exist in the numbers predicted by CDM, then as they fall through the disk of their parent galaxy, they should heat the disk and cause it to thicken (T{\\'o}th \\& Ostriker, 1992, Moore et al, 1999b). This is contrary to some observations of old thin disk systems or galaxies with no thick disk components, although it is now being argued that the amount of heating and thickening has been overestimated (Font et al., 2001, Vel{\\'a}zquez \\& White, 1999). This is clearly a matter for further investigation.  The Virgo cluster cannot have been assembled out of objects like the LG without some additional physical mechanism being involved that increases the ratio of dwarf to giant galaxies. Virgo is a very dense environment where many galaxy-galaxy interactions are likely to have occurred due to its short dynamical crossing time compared with UMa ($\\approx$0.1$H_{0}$ and $\\approx H_{0}$ respectively, Trentham et al, 2001, Trentham $\\&$ Tully, 2002). Virgo is also an X-ray cluster so galaxies in the cluster core move through a relatively dense inter-galactic gas. UMa is also probably in a much earlier stage of formation than Virgo. The question is are these the differences that lead to Virgo being so different?  The large dwarf galaxy population found in Virgo seems to lend some credence to the theory of dwarf galaxy formation by galaxy harassment, an idea put forward by Moore et al (1999b). In this scenario, dE galaxies are formed when infalling LSB spiral galaxies are harassed in the cluster by the giant galaxies, and lose their gas resulting in a morphological transformation into a dE. Further evidence to support this theory comes from a study of the Virgo cluster dwarfs, conducted by Conselice et al. (2001). They show that the dwarf ellipticals found in Virgo have a velocity distribution closer to that of the spirals than that of the earlier type galaxies. The dwarf velocity distribution is quite wide, and is non-Gaussian with a total velocity dispersion of 726km s$^{-1}$. This is similar to that of the spirals, which is 776km s$^{-1}$. The dwarf galaxies appear not to be relaxed and are less dynamically evolved than the Virgo cluster core elliptical population. However, in Sabatini et al. (2004) we show that the dwarf galaxies we detect in the Virgo cluster are too small to be the result of the harassment process proposed by Moore et al. (1999). We propose that the dE galaxies are the result of an earlier infalling dI galaxy population. These galaxies may be associated with the faint blue galaxies seen at higher redshift ($0.5<z<1.5$). Sabatini et al. (2004) suggest that the star formation of these small infalling haloes is enhanced by the weak tidal interactions with the cluster potential and other cluster galaxies - these types of interactions are not available to galaxies in UMa or the MGS. These haloes have their evolution advanced by the cluster environment. Maybe small dark matter haloes require these sorts of tidal interactions to light up and reveal their baryons."
    },
    "0404/gr-qc0404128.txt": {
        "abstract": "We provide a method for analytically constructing high-accuracy templates for the gravitational wave signals emitted by compact binaries moving in inspiralling eccentric orbits. By contrast to the simpler problem of modeling the gravitational wave signals emitted by inspiralling {\\it circular} orbits, which contain only two different time scales, namely those associated with the orbital motion and the radiation reaction, the case of {\\it inspiralling eccentric} orbits involves {\\it three different time scales}: orbital period, periastron precession and radiation-reaction time scales. By using an improved `method of variation of constants', we show how to combine these three time scales, without making the usual approximation of treating the radiative time scale as an adiabatic process. We explicitly implement our method at the 2.5PN post-Newtonian accuracy. Our final results can be viewed as computing new `post-adiabatic' short period contributions to the orbital phasing, or equivalently, new short-period contributions to the gravitational wave polarizations, $h_{+,\\times}$, that should be explicitly added to the `post-Newtonian' expansion for $h_{+,\\times}$, if one treats radiative effects on the orbital phasing of the latter in the usual adiabatic approximation. Our results should be of importance both for the LIGO/VIRGO/GEO network of ground based interferometric gravitational wave detectors (especially if Kozai oscillations turn out to be significant in globular cluster triplets), and for the future space-based interferometer LISA. ",
        "introduction": "\\label{IntroSec} Inspiralling black hole binaries are considered to be the most probable source of detectable gravitational radiation for the  first generation laser interferometric gravitational-wave detectors that are operational or nearing the completion of their construction phase \\cite{GWIFs}. The above understanding is based on both astrophysical considerations \\cite{GLPPS} and the availability of  highly accurate general relativistic theoretical waveforms required to pluck the weak gravitational wave signal from the  noisy interferometric  data \\cite{DIJS}. Inspiralling compact binaries are usually modeled as point particles in {\\it quasi-circular} orbits. For long lived compact binaries, the  quasi-circular approximation is quite appropriate, as the radiation reaction decreases the orbital  \\ecc to negligible values by  the  epoch the emitted gravitational radiation  enters the sensitive bandwidth of the interferometers. It is easy to deduce that for an isolated binary, the \\ecc goes down roughly  by a factor of three, when its semi-major axis is halved \\cite{P64}.   In recent times, however, scenarios involving compact {\\it eccentric} binaries are being suggested  as  potential gravitational wave sources {\\em even}  for the terrestrial gravitational-wave detectors. For instance, one such  proposed astrophysical scenario \\cite{MH02}, involves hierarchical triplets (say, 123), usually modeled to consist of an inner (say, 12) and an outer binary (say, 03, where 0 denotes the center of mass of the 12 binary). If the mutual inclination angle between the orbital planes of the  inner and of  the outer binary is large enough, then the time averaged tidal force on the inner binary may induce oscillations in its eccentricity, known in the literature as the Kozai  mechanism \\cite{K62}. It was shown that in globular clusters,  the inner binaries of hierarchical triplets undergoing Kozai oscillations can merge under gravitational radiation reaction \\cite{MH02}. Later, it was shown that %roughly $30\\%$ a good fraction of such systems will have \\ecc $\\sim 0.1$, when emitted gravitational radiation from these binaries passes through $10 \\,{\\rm Hz} $ \\cite{W03}.  It is also believed that  the effect of  orbital \\ecc will have to be modeled accurately while computing theoretical waveforms for compact binaries relevant for the space-based laser interferometric gravitational-wave detector LISA \\cite{LISA}. The above statement is supported by a recent finding that by observing stellar mass black hole binaries in highly eccentric orbits - which may be common in  globular clusters - one can estimate accurately not just the  intrinsic binary parameters like  masses and eccentricity, but even the position of  the host cluster \\cite {B02}. It is also well known that the cosmological supermassive black hole binaries, embedded in surrounding stellar populations, would be powerful gravitational wave sources for detectors like LISA \\cite {TB76}. However, the question whether these binaries will be in a {\\it quasi-circular} or an  {\\it  inspiralling eccentric} orbit by the time their gravitational waves   are detectable by LISA is  not yet settled \\cite{S_MOG}. %Last year, Recently, it was suggested that if the  Kozai mechanism were operative, these supermassive black hole binaries, in highly eccentric orbits, would merge within the Hubble time \\cite{BLS03}. Very recently, it was shown that, using very-long-baseline interferometer (VLBI),  the  unresolved core of the radio galaxy 3C 66B executes  well-defined elliptical motions with an yearly  period, which was interpreted as the first direct evidence for the detection of a supermassive black hole binary \\cite{SIMT}. The   above observation raises the interesting possibility of  being able to  detect gravitational waves from supermassive black hole binaries in  eccentric orbits using  LISA.  Even though various versions of `ready to use' high accuracy  search templates for  inspiralling compact binaries with arbitrary masses in {\\it quasi-circular} orbit exist \\cite{DIJS}, so far none is available for compact objects in {\\it inspiralling eccentric} orbits. Before characterizing the strategy and new results presented in this paper, let us summarize the relevant existing literature on the influence of eccentricity on the  gravitational wave polarizations, $h_+ $ and $h_{\\times}$, with and without including the gravitational radiation reaction. After the seminal work of Peters and Mathews \\cite{PM63}, it was in the context of spacecraft Doppler detection of gravitational waves from compact binaries that the first explicit expressions for Newtonian accurate gravitational wave polarization states were derived \\cite{W87}. %Ref.~\\cite{LW90} used Using the method of osculating orbital elements and a numerical integration approach %to study the the effects of eccentricity and dominant radiation damping on $h_+ $ and $h_{\\times}$ was studied in \\cite{LW90}. First- and First-and-half-post-Newtonian accurate analytic expressions for far-zone fluxes and gravitational wave polarizations, for compact binaries in eccentric orbits, were computed in a series of papers in \\cite{GS89}. Using Newtonian accurate orbital motion, Refs.~\\cite{MBM95,MBM94} studied the  effect on gravitational wave polarizations of introducing by hand some secular effects either in the longitude of the periastron \\cite{MBM95} or in the semi-major axis and eccentricity \\cite{MBM94}. Using such waveforms, the  influence of eccentricity on the signal to noise ratio in  gravitational wave data analysis, was examined in \\cite{MBM94,MP99, PSLR}. These waveforms were also used to show that LISA will be sensitive to eccentric Galactic binary neutron stars \\cite{GIKPP} and that  by measuring their periastron advance, accurate estimates for the total mass of these binaries may be obtained \\cite{S01}. However, the widely used  gravitational wave templates, to detect gravitational waves from compact binaries in {\\it quasi-circular orbits}, are based on the  {\\it second post-Newtonian} accurate expressions for $h_+ $ and $h_{\\times}$, supplemented by expressions giving {\\it adiabatic} time evolution for the orbital phase and frequency also to the second post-Newtonian order \\cite{BDIWW,BIWW96} [However, see \\cite{DIJS_N} for the (numerical) construction of gravitational wave templates going {\\it beyond} the adiabatic approximation in the case of {\\it quasi-circular} orbits]. The second post-Newtonian order, usually referred to as the 2PN order, gives corrections to  leading order contributions in  gravitational theory, to  order of $\\sim (v/c)^4 \\sim (G\\,m\\,/c^2\\, r)^2$, where $m, v $ and $r$ are the total mass, orbital velocity and  the separation of the binary. The 2PN contribution to the gravitational waveform, required for the construction of $h_+ $ and $h_{\\times}$, and the associated far-zone fluxes for binaries moving in general {\\it eccentric} orbits, in harmonic gauge,  were computed in \\cite{WW96,GI97}. Employing the 2PN accurate generalized quasi-Keplerian parametrization of Damour, Sch\\\"afer and Wex \\cite{DS88,SW93} available in Arnowitt, Deser and Misner (ADM) coordinates, to represent relativistic motion  binaries in eccentric orbits, 2PN corrections to the rate of decay of the orbital elements of the representation as well as the explicit expressions for $h_+ $ and $h_{\\times}$ were provided in \\cite{GI97,GI02,GI03er}. The above mentioned expressions for $h_+$ and $h_\\times$ represent gravitational radiation from an eccentric binary, during that stage of inspiral where the  gravitational radiation reaction is so small that the orbital parameters can be treated as essentially unchanging over a few orbital periods (`adiabatic approximation'). The effects of eccentricity, advance of periastron and orbital inclination on power spectrum of the dominant Newtonian part of the polarizations were also presented in \\cite{GI02}.  The aim of this paper is to provide a method for explicitly constructing high-accuracy waveforms emitted by compact binaries moving in {\\it inspiralling eccentric orbits}. Compared to the  existing high-accuracy waveforms for {\\it inspiralling circular orbits} \\cite{DIJS}, the inclusion of orbital eccentricity into such templates is a non-trivial task as eccentricity brings along a new physical aspect, the precession of the periastron, and thus one must contend with precession and radiation reaction at the same time. In the quasi-circular case, there are two time-scales related to the orbital period and the radiation reaction. In the quasi-eccentric case, one has  a third  additional time-scale related to the precession of periastron. The technical problem we shall tackle and solve below is that of combining, in a consistent framework these three time scales, without making the usual approximation of treating the radiation-reaction time scale as an adiabatic process. We then explicitly implement our method at the `2PN+2.5PN' accuracy, {\\it i.e.} the effect of perturbing a `2PN-accurate analytic' description of eccentric orbits by the 2.5PN level radiation-reaction. It is useful to note that the gravitational wave observations of  inspiralling compact binaries, are  analogous to the high precision radio-wave observations of binary pulsars. %\\cite{TD82,DIS98}. The latter makes use of an accurate relativistic `timing formula'  which requires the  solution to the relativistic  equation of motion  for  a compact  binary moving  in an  elliptical orbit while the former  demands  accurate `phasing', i.e. an accurate mathematical modeling of the continuous time evolution of the gravitational waveform. The mathematical formulation, which resulted in an accurate `timing formula', as given in \\cite{DD85,DT92}, was obtained by one of us   many years  ago \\cite{TD83, TD_F85}. The present work will rely on techniques from \\cite{TD83, TD_F85} to combine the mentioned above three time scales % mentioned above to implement post-Newtonian accurate `phasing' for compact binaries moving in elliptical orbits.  The plan of the paper is the following. In Sec. \\ref{PhasingSec}, we outline the steps required to do the `phasing'. The method we use to find the domain of validity of our approach is presented in Sec. \\ref{DelineateSec}. In Sec. \\ref{MethodSec}, we formulate the procedure required to construct time evolving $h_\\times$ and $h_+$. We apply, in Sec. \\ref{2.5PNSec},  the formalism to do the 2.5PN accurate phasing. Sec.\\ref{PNadiabaticSec} displays the expressions required to extend the `phasing' to higher PN orders. Finally, in Sec. \\ref{ConcludeS}, we summarize our approach, results and point out possible future extensions. ",
        "conclusions": "\\label{ConcludeS} Let us summarize what we have proposed in this paper and point out possible extensions. We have provided a method for analytically constructing high-accuracy templates for the gravitational wave signals emitted by compact binaries when they move in inspiralling slowly precessing eccentric orbits. In  contrast to the simpler problem of modeling the gravitational wave signals emitted by inspiralling {\\it circular} orbits, which contain only two different time scales (orbital period and radiation reaction time scale), the case of {\\it inspiralling eccentric} orbits involves {\\it three different time scales}: orbital period, periastron precession and radiation-reaction time scales. An improved `method of variation of constants' is used to combine these three time scales, without making the usual approximation of treating %the radiative time scale as an adiabatic process. adiabatically the radiative time scale. %as an adiabatic process. By going to a suitable center-of-mass frame, the transverse traceless (TT) radiation field and hence the GW polarizations are %given as PN expansions of the form (\\ref{Eq4}) expressed as PN expansions of the form given in Eq. (\\ref{Eq4}) in harmonic coordinates involving only the relative position and %relative velocity. velocity. The polarisations can be rewritten in terms of the ADM positions and velocities by using the contact transformations available to move from the harmonic  coordinates  to the ADM coordinates. In the ADM coordinates, the unperturbed 2PN (3PN) motion may be explicitly solved by a  generalized quasi-Keplerian representation involving two angles $l$ and $\\lambda$ and four constants of the 2PN (3PN) motion $c_\\alpha= c_1,c_2,c_l$ and $c_\\lambda$. By a Lagrange method of variation of constants, this unperturbed solution is used to %obtain prescribe a general solution to the perturbed (by radiation-reaction) system of the form given by Eqs. (\\ref{Eq17}) and (\\ref{Eq18}), in terms of the varying $c_\\alpha$'s. Among the four $c_\\alpha$'s, two ($c_l$ and $c_\\lambda$) are found to be constants, while the other  two %$c_\\alpha$ $c_\\alpha$'s satisfy two coupled first order differential equations. %an evolution equation of the simple form (\\ref{Eq}). A two scale decomposition of the $c_\\alpha$'s is made to model the combination of the slow (radiation reaction, secular) drift and the fast (orbital time scale, periodic) oscillations. This allows us to decompose the TT gravitational wave amplitudes or polarizations into a part associated with the secular variations %and the other part and another part associated with  the fast oscillations. If one %were to expand expands in the fast oscillations, the  oscillatory contributions to the phasing would be  equivalent to new additional ${\\cal{O}}(v^5/c^5)$ contributions (which stay always small ) %to ${\\cal{O}}(v^5/c^5)$) to be added to the usually considered 2.5PN %the GW amplitude $h_{ij}^5$ in the adiabatic approximation. Therefore, as long as an amplitude correction ${\\cal{O}}(v^5/c^5)$ is not needed our results show that it is enough to use only secularly varying $\\bar{n}$ and $\\bar{e}_t$. Note that our description of secular effects automatically include an effect of a secular acceleration of the periastron precession, analogous to the usual secular acceleration of the orbital motion. Namely, the secular part of the angle $\\lambda -l$ measuring the periastron longitude varies as \\be \\bar{\\lambda}-\\bar{l}=\\int \\bar{k}\\bar{n}dt \\ee where $\\bar{k}(t)=\\delta\\Phi/2\\pi$ is the secularly varying periastron precession per orbital period.  Schematically,  thus our work may be summarized as follows: \\vspace*{0.3in} \\newline \\framebox(280,25)[c]{\\Large $ h_{+,\\times}(r^{\\rm harm},\\phi^{\\rm harm}, \\dot{r}^{\\rm harm},\\dot{\\phi}^{\\rm harm})$ } %$\\rightarrow$ \\newline \\mbox{ \\hspace{1.5in} \\Large $\\Downarrow$} \\vspace*{0.1in} \\newline \\framebox(280,25)[c]{ \\large $h_{+,\\times}(r^{\\rm ADM},\\phi^{\\rm ADM},\\dot{r}^{\\rm ADM}, \\dot{\\phi}^{\\rm ADM})$} \\newline \\mbox{ \\hspace{1.5in} \\Large $\\Downarrow$} \\vspace*{0.1in} \\newline \\framebox(280,25)[c]{ \\large %$h_{+,\\times}(n,e_t,c_l,c_\\lambda) $} $h_{+,\\times}(n,e_t,l,\\lambda) $} \\newline \\mbox{ \\hspace{1.5in} \\Large $\\Downarrow$} \\vspace*{0.1in} \\newline \\framebox(280,25)[c]{ \\large $ h_{+,\\times}(\\bar{n}+\\tilde{n},\\bar{e}_t+\\tilde{e}_t, \\bar{l}+\\tilde{l},\\bar{\\lambda}+\\tilde{\\lambda}) $} \\newline \\mbox{ \\hspace{1.5in} \\Large $\\Downarrow$} \\vspace*{0.1in} \\newline \\framebox(350,40)[c]{ \\large $h_{+,\\times}(\\bar{n},\\bar{e}_t, \\bar{l},\\bar{\\lambda})+ \\left\\{\\frac{\\partial h_{+,\\times}}{\\partial n} \\tilde{n}+ \\frac{\\partial h_{+,\\times}}{\\partial e_t} \\tilde{e_t}+ \\frac{\\partial h_{+,\\times}}{\\partial l} \\tilde{l}+ \\frac{\\partial h_{+,\\times}}{\\partial \\lambda} \\tilde{\\lambda}\\right\\} $} \\newline %\\vspace*{0.1in} %\\newline %\\mbox{ \\hspace{1.8in} \\Huge $\\equiv$} \\hspace*{1.8in} \\mbox{  \\Huge $\\equiv$} \\vspace*{0.1in} \\newline %\\newline %\\vspace*{0.1in} \\vspace*{0.2in} \\hspace*{0.2in} \\framebox(280,40)[c]{ \\Large $ %\\equiv h_{+,\\times}(\\bar{n},\\bar{e}_t, \\bar{l},\\bar{\\lambda})+ \\frac{1}{c^5}\\,\\delta h_{+,\\times}^5(\\tilde{n},\\tilde{e}_t, \\tilde{l},\\tilde{\\lambda}) %/c^5. $}  In this paper, the orbit of the binary was treated to be an inspiralling, slowly precessing ellipse, which prevented us from approaching the LSO. However, as mentioned earlier, using the EOB approach, we intend to explore the orbital dynamics and hence the evolution of gravitational wave polarizations near the  LSO in the near future. There are also quite a few generalizations which can be tackled using the formalism presented here and we list a few of them here. In this paper, the conservative dynamics was restricted to the 2PN order and therefore, a natural extension will be to incorporate the 3PN conservative dynamics.  We also restricted our approach to compact binaries consisting of non-spinning point masses. It is possible to extend the generalized quasi-Keplerian parametrization and hence our method to spinning compact binaries. To do that, first, one needs to extend the quasi-Keplerian representation to include the  effects due to spin-orbit and spin-spin interactions, which requires generalizing a restricted analysis done in \\cite{W95} and this is under investigation \\cite{KG04}. Finally,  starting from    the  gravitational wave polarizations from {\\it inspiralling eccentric} binaries in the time-domain, it should be possible to  perform  a  spectral analysis  and see how their power spectrum depends on various orbital elements like $n$, $e_t$ and $i$."
    },
    "0404/astro-ph0404435_arXiv.txt": {
        "abstract": "We report near-infrared photometry, spectroscopy, and speckle imaging of the hot, luminous star we identify as candidate LBV 1806-20 \\footnote{We note that this star has not been observed to undergo a major LBV outburst yet, and thus some may not consider it to be an LBV (though it is undoubtedly luminous, blue, and variable, and shares many spectral characteristics with {\\it bone fide} LBV stars, as we show here).  Hereafter, to avoid awkward phrasing, we refer to the candidate LBV as LBV 1806-20 as is done in the literature.} .  We also present photometry and spectroscopy of 3 nearby stars, which are members of the same star cluster containing LBV 1806-20 and SGR 1806-20.  The spectroscopy and photometry show that LBV 1806-20 is similar in many respects to the luminous ``Pistol Star'', albeit with some important differences.  They also provide estimates of the effective temperature and reddening of LBV 1806-20, and confirm distance estimates, leading to a best estimate for the luminosity of this star of $> 5 \\times 10^6 \\ L_{\\odot}$.  The nearby cluster stars have spectral types and inferred absolute magnitudes which confirm the distance (and thus luminosity) estimate for LBV 1806-20.  If we drop kinematic measurements of the distance ($15.1 ^{+1.8}_{-1.3}$ kpc), we have a lower limit on the distance of $>9.5$ kpc, and on the luminosity of $>2 \\times 10^6 \\ L_{\\odot}$, based on the cluster stars.  If we drop both the kinematic and cluster star indicators for distance, an ammonia absorption feature sets yet another lower limit to the distance of $>5.7$ kpc, with a corresponding luminosity estimate of $>7 \\times 10^5 \\ L_{\\odot}$ for the candidate LBV 1806-20. Furthermore, based on very high angular-resolution speckle images, we determine that LBV 1806-20 is not a cluster of stars, but is rather a single star or binary system. Simple arguments based on the Eddington luminosity lead to an estimate of the total mass of LBV 1806-20 (single or binary) exceeding $190 \\ M_{\\odot}$.  We discuss the possible uncertainties in these results, and their implications for the star formation history of this cluster. ",
        "introduction": "Mounting evidence gathered in recent years indicates that stars may be formed with masses much greater than previously thought possible.  The hot luminous ``Pistol Star'' near our Galaxy's center, for instance, has an estimated mass of $>150 \\ M_{\\odot}$ \\citep{Figer98}, and the stars R136a1 and R136a2 in the Large Magellanic Cloud each have masses of $140-155 M_{\\odot}$ \\citep{Massey}.  Since the luminosities of such massive stars exceed $10^6 \\ L_{\\odot}$, relatively small populations of these stars can dominate the power output of their host galaxy during their lifetimes. Furthermore, their deaths may spread chemically-enriched material into the galaxy and leave behind black holes as remnants.  The death events may be responsible for gamma-ray bursts in the ``collapsar'' scenario (e.g.  \\citet{price}; \\citet{hjorth}; \\citet{MWH}), and the relatively large remnant black holes may also explain the so-called ``intermediate-mass'' black holes currently being discovered in nearby galaxies (e.g. \\citet{IMBH} and references therein).  Thus, probing the upper limit for stellar mass has an important impact on our understanding of a wide range of astrophysical phenomena, including the chemical evolution of matter in our Galaxy and external galaxies, the history of galaxy and structure formation in the Universe, the formation of black holes in their dying supernova events, and possibly the origin of gamma-ray bursts via collapsars.  We report here new observations of a luminous star we identify as LBV 1806-20.  This star was first identified as a potential counterpart to the soft gamma-ray repeater SGR 1806-20 \\citep{Shri}, with high near-infrared brightness ($K = 8.4$ mag) despite significant absorption from interstellar dust in the Galactic Plane.  Subsequent moderate-resolution ($R \\sim 700$) infrared spectroscopy revealed it to be a candidate luminous blue variable (LBV) star and one of the most luminous stars in the Galaxy, with $L > 10^6 \\ L_{\\odot}$ (\\citet{vK95} ; \\citet{Corbel}).  This star is known to lie at the brightness peak of the radio nebula G10.0-0.3 (\\citet{Shri93}; \\citet{Vasisht}).  However, the revised Inter-Planetary Network (IPN) localization of SGR 1806-20 indicated that the SGR was significantly offset from the position of the candidate LBV star and the coincident core of G10.0-0.3 \\citep{Hurley}.  Recent {\\it Chandra} and infrared observations confirm that the SGR lies $\\sim 12 \\arcsec$ away from the candidate LBV (\\citet{kaplan}; \\citet{eiken01}).  In addition, \\citet{gaensler} show that G10.0-0.3 is not a supernova remnant at all, but is a radio nebula powered by the tremendous wind of the candidate LBV star at its core.  The apparent conundrum presented by this scenario -- why we find two such rare objects so close to each other on the sky without any apparent physical connection -- has been resolved by the observations of \\citet{Fuchs} who showed that LBV 1806-20 is a member of a cluster of massive stars, and of \\citet{eiken01} who showed that SGR 1806-20 appears to be a member of the same cluster.  Thus, while these two rare objects are not identical, they are related through common cluster membership. Interestingly, \\citet{vrba} show that one of the other SGRs in the Milky Way may also be associated with a massive star cluster.  The distance to the candidate LBV, SGR, and their associated cluster has also been a subject of some discussion in the literature.  Initial studies of CO emission from molecular clouds towards this line of sight and the detection of $\\rm NH_3$ absorption against the radio continuum indicated an extinction of $A_V = 35 \\pm 5$ mag and a best estimate for the distance of $14.5 \\pm 1.4$ kpc \\citep{Corbel} -- again confirming its status as one of the most luminous stars known. However, \\citet{blum01} present infrared spectra of members of a cluster in the nearby \\hii \\ region G10.2-0.3, which is also part of the (apparent) W31 giant molecular cloud complex containing G10.0-0.3, and find an apparently conflicting distance of $\\sim 3.4 \\pm 0.6$ kpc. This apparent conflict has also been recently resolved by \\citet{CE}, who present higher-resolution millimeter and infrared observations of G10.0-0.3 and G10.2-0.3.  They find that W31 is actually resolved into at least 2 components along the line-of-sight, with one component at $d \\sim 4$ kpc and with extinction $A_V \\sim 15$ mag (in excellent agreement with the \\citet{blum01} observations of G10.2-0.3) and another component at a (refined) distance of $d = 15.1^{+1.8}_{-1.3}$ kpc and with $A_V = 37 \\pm 3$ mag.  The radial velocity of LBV 1806-20 matches both of these components, but the $NH_3$ absorption towards the core of G10.0-0.3 (and thus LBV 1806-20) due to a cloud at $d = 5.7$ kpc unambiguously places the star in the ``far'' component of W31.  In addition, both the infrared extinction towards LBV 1806-20 (\\citet{vK95}; see also below) and the X-ray absorption column towards SGR 1806-20 (\\citet{eiken01}; \\citet{mereghetti}) match the expected extinction towards the ``far'' component of W31 and differ from that of the ``near'' component of \\citet{blum01} by $\\sim 15$ mag, thus confirming the association of the candidate LBV, SGR, and associated star cluster with the ``far'' component of W31.  In order to further investigate this intriguing object, we obtained near-infrared images and spectra of LBV 1806-20 and several nearby stars.  In Section 2, we describe these observations and their reduction.  In Section 3, we describe the analysis of the resulting data, including the spectral types of the candidate LBV and cluster stars, refined analyses of the reddening, confirmation of the distance estimate of \\citet{CE}, and the resulting luminosity estimate for LBV 1806-20.  In Section 4, we discuss the uncertainties in these measurements and their implications for our understanding of the formation and evolution of extremely massive stars and the birth environment of SGRs.  Finally, in Section 5 we present our conclusions. ",
        "conclusions": "We have presented near-infrared imaging and spectroscopy of the luminous star LBV 1806-20 and 3 other nearby luminous stars. Based on the results we conclude that LBV 1806-20 has spectral characteristics very similar to those of AG Car, the Pistol Star, and P Cyg -- all luminous blue variables -- and is thus likely to be an LBV itself.  The nearby luminous stars B, C, and D are Wolf-Rayet WC9d and possible blue hypergiant stars forming part of a cluster which includes LBV 1806-20.  Their absolute magnitudes and bolometric luminosities are consistent with other known stars with similar spectral types, confirming the distance and reddening estimates for LBV 1806-20 ($15.1 ^{+1.8}_{-1.3}$ kpc and $A_V = 29 \\pm 2$ mag). With a surface temperature in the range 18000-32000K, LBV 1806-20 has a bolometric luminosity $>5 \\times 10^6 \\ L_{\\odot}$.  If we drop kinematic measurements of the distance ($15.1 ^{+1.8}_{-1.3}$ kpc), we have a lower limit on the distance of $>9.5$ kpc, and on the luminosity of $>2 \\times 10^6 \\ L_{\\odot}$, based on the cluster stars.  If we drop both the kinematic and cluster star indicators for distance, an ammonia absorption feature sets yet another lower limit to the distance of $>5.7$ kpc, with a corresponding luminosity estimate of $>7 \\times 10^5 \\ L_{\\odot}$ for the candidate LBV 1806-20.  Our speckle imaging shows conclusively that LBV 1806-20 is {\\it not} an unresolved cluster of stars, though it may be a binary/multiple system.  If LBV 1806-20 is a single or multiple star, its total mass exceeds $190 M_{\\odot}$ (at the $\\sim 15$ kpc distance).  Finally, the presence of LBV 1806-20 with more evolved stars in the same cluster (i.e. the W-R WCL star and SGR 1806-20) implies that star formation may have occurred over multiple epochs in this region of space."
    },
    "0404/astro-ph0404603_arXiv.txt": {
        "abstract": "{A pointed observation of the low-luminosity galactic source 4U 2206+54 was carried out in November 1998 with {\\it BeppoSAX}. The light curve of 4U 2206+54 shows erratic variability on a timescale of $\\sim$1 hour; neither hardness variations nor time periodicities are detected throughout this 67 ks long observation. Thanks to the wide spectral coverage capabilities of {\\it BeppoSAX} we could observe the source X--ray continuum over three energy decades, from 0.6 to 60 keV. The spectrum could be equally well fitted either with a blackbody plus Comptonization or with a high energy cutoff power law. No iron emission around 6.5 keV was detected, while a tentative detection of a cyclotron resonant feature in absorption is presented. Comparison of the present {\\it BeppoSAX} data with the information available in the literature for this source suggests that 4U 2206+54 is a close binary system in which a (possibly magnetized) NS is accreting from the companion star wind. ",
        "introduction": "Massive X--ray Binaries (MXRBs) are double systems composed of a compact object, generally a neutron star (NS), orbiting an early-type star and accreting matter from it. In X--rays, MXRBs can be seen as persistently bright, with luminosities greater than 10$^{35}$ erg s$^{-1}$, or present a transient behaviour characterized by quiescent phases, with emissions around 10$^{34}$ erg s$^{-1}$ or less, followed by intense (up to $\\sim$10$^{38}$ erg s$^{-1}$ at peak) outbursts; in several cases, these outbursts show a periodic trend as a result of the orbital motion of the NS along a highly eccentric orbit (see White et al. 1995 and van Paradijs 1995 for a review). Usually, the former group is associated with compact objects steadily accreting from the companion via Roche lobe overflow and/or stellar wind, while in the latter one accretion is discontinuous and occurs when the compact source enters a disk-like envelope around the companion star or, more generally, interacts more closely with the companion as it approaches periastron (e.g., Corbet 1986).  There are however MXRBs which do not fit this classification, i.e. the so-called low-luminosity MXRBs, characterized by their relatively low persistent emission in the X--ray domain (10$^{34}$--10$^{35}$ erg s$^{-1}$) compared to those of persistent MXRBs and which do not display outbursts. The X--ray source 4U 2206+54 is one of these objects.  It was discovered by {\\it Uhuru} (Giacconi et al. 1972), was monitored with {\\it EXOSAT} between 1983 and 1985 by Saraswat \\& Apparao (1992) who reported aperiodic hard flares from the source lasting a few hundred seconds and producing variations of a factor 3 to 5, and long-term variations of a factor of $\\sim$20 in the 2--10 keV persistent luminosity ($\\approx$0.3--5$\\times$10$^{34}$ erg s$^{-1}$). These authors also reported a pulse period of $\\sim$400 s, which however has been recently questioned by Corbet \\& Peele (2001) from a re-analysis of the {\\it EXOSAT} data as well as of archival {\\it RXTE} data. Corbet \\& Peele (2001) further reported, on the basis of ASM observations, a 9.6-d periodicity in the X--ray flux of the source; they also modeled the {\\it RXTE} spectra obtained on two occasions (March 11 and 13, 1997) using a power law modified with an exponential cutoff. They found a flux decrease by a factor of three (from 3.12$\\times$10$^{-10}$ to 1.14$\\times$10$^{-10}$ erg cm$^{-2}$ s$^{-1}$) between the two pointings. Negueruela \\& Reig (2001) reanalyzed the {\\it RXTE} pointing of March 11, 1997 obtaining comparable results; they also confirmed the presence of flares during which they found a positive correlation between source hardness and flux. These authors also did not detect any X--ray pulsation from the object.  The X--ray spectral characteristics of this source are typical of accretion onto a NS from a wind coming from the companion star (Negueruela \\& Reig 2001), although the hypotheses of an accreting white dwarf (WD: Saraswat \\& Apparao 1992; Corbet \\& Peele 2001) or black hole (BH: Negueruela \\& Reig 2001) were also considered.  The optical counterpart was identified by Steiner at al. (1984) as the early-type star BD +53$^{\\circ}$2790, located at 2.5 kpc from the Earth. This star was subsequently thoroughly studied in its optical-UV spectroscopic properties by Negueruela \\& Reig (2001) who classified it as a peculiar late O-type active star. No radio counterpart has been detected so far (Nelson \\& Spencer 1988).  The information available on 4U 2206+54 is therefore not conclusive to understand the nature of the accreting source. In particular, the lack of knowledge of its X--ray spectrum above 30 keV hinders any hypothesis on the presence of a Cyclotron Resonant Feature (CRF) and thus any conjecture on the magnetic field of the accreting source as well as on its nature. Likewise, the poor sampling concerning the soft side of the X--ray spectrum never allowed an accurate estimate of the hydrogen column density $N_{\\rm H}$; also, this did not allow a sensitive search for a soft component in the emission from this source. Moreover, a further independent check of the presence of a periodic variability (or the lack thereof) in X--rays is also needed.  Therefore, to explore the timing and spectral behaviour of 4U 2206+54 over a broad spectral range, with particular attention to both soft ($<$2 keV) and hard ($>$20 keV) X--ray domains, we observed this source with {\\it BeppoSAX} (Boella et al. 1997a).  The paper is organized as follows: Sect. 2 will illustrate the observations and the data analysis, while in Sect. 3 the results showing the X--ray spectral and timing behaviours of 4U 2206+54 will be reported; in Sect. 4 a discussion will be given. ",
        "conclusions": "The {\\it BeppoSAX} observation of 4U 2206+54 reported in this paper allowed us to study the properties of the source in great detail, and for the first time, in the 0.6--60 keV range. The source exhibited an average 2--10 keV unabsorbed luminosity of 3.1$\\times$10$^{34}$ erg s$^{-1}$ during this pointing, thus comparable with the ``high-state\" emission observed by Saraswat \\& Apparao (1992) in 1983 and 1985 from {\\it EXOSAT} data, but $\\sim$10 times brighter than that observed by these authors in 1984, and a factor $\\sim$10 fainter than during the 1997 {\\it RXTE} pointing (e.g., Negueruela \\& Reig 2001). Thus the source can vary by about 10 fold, and as much as a factor of $\\sim$100, on year-long timescales. This can also be seen in Fig. 1 of Corbet \\& Peele (2001), where the entire {\\it RXTE}/ASM 1.5--12 keV light curve of 4U 2206+54 between years 1996 and 2001 is shown. The periodogram in Fig. 2 of these authors shows a peak at a frequency of $\\approx$4$\\times$10$^{-3}$ d$^{-1}$: if real, this would indicate the presence of a further (superorbital?) periodicity of $\\approx$250 d in the system.  On much shorter time scales, the 2--10 keV light curve shows flares of remarkable (a factor ten) intensity variations lasting about 1 hr, with finer variability down to timescales of 50--100 s. This flaring emission does not appear to imply spectral variation correlated with the intensity. Moreover, the PSD analysis indicates that the variable X--ray emission from the source is due to a stochastic phenomenon. These long- and short-term variability characteristics point to an explanation for this X--ray activity as due to random inhomogeneities in the accretion flow onto a compact object (e.g. van der Klis 1995).  The colorless variability result obtained from our {\\it BeppoSAX} data is at odds with that found by Saraswat \\& Apparao (1992) and by Negueruela \\& Reig (2001). The explanation for this is not clear. Concerning the findings by Saraswat \\& Apparao (1992) this discrepancy can be due to slight secular variations of the source hardness ratio in observations separated in time by about two years: indeed, if one considers the 1983 and 1985 {\\it EXOSAT} data in their Fig. 4 separately, no significant hardness-intensity dependence is observed. As regards instead the discrepancy with the data from the {\\it RXTE} pointing in Negueruela \\& Reig (2001) it may be possible that this hardness-intensity dependence becomes more evident at larger source luminosities.  Moreover, the PSD of {\\it BeppoSAX} data does not show any periodicity in the 10$^{-3}$--1 Hz range. We thus independently confirm the results by Corbet \\& Peele (2001) and Negueruela \\& Reig (2001) that no $\\approx$400 s X--ray modulation comes from 4U 2206+54.  X--ray spectroscopy with {\\it BeppoSAX} can help us in better understanding the accretion dynamics as well as the nature of the accretor. Indeed, we obtained for 4U 2206+54 an unprecedented simultaneous spectral coverage of the 0.6--60 keV range. Our results show that the ``classic\" model generally used to fit the X--ray spectra of MXRBs hosting an accreting NS (White et al. 1983a) works very well, as in the case of the March 1997 {\\it RXTE} data. In these observations the spectral parameters $N_{\\rm H}$, $\\Gamma$, $E_{\\rm cut}$ and $E_{\\rm fold}$ were in the range 1.12--1.71, (2.7--4.6)$\\times$10$^{22}$ cm$^{-2}$, 5.3--7.3 and 10.5--17.3, respectively, with the values at the lower edge of the interval holding at low source fluxes (Corbet \\& Peele 2001).  However, when we compare our best-fit parameters with the {\\it RXTE} spectral findings obtained when the source flux was highest (corresponding to a 2--10 keV luminosity of 2.3$\\times$10$^{35}$ erg s$^{-1}$), we see that in our data (i) the spectral slope $\\Gamma$ is substantially harder; (ii)  both characteristic energies of the model, $E_{\\rm cut}$ and $E_{\\rm fold}$, are lower by a factor $\\sim$2; (iii) the hydrogen column density $N_{\\rm H}$ is lower by a factor $\\sim$6. Comparison with the lower-intensity {\\it RXTE} observation (which implies a 2--10 keV luminosity of 8.5$\\times$10$^{34}$ erg s$^{-1}$) shows a consistency with the {\\it BeppoSAX} data barring the $N_{\\rm H}$ value, which is still $\\sim$3 times higher in the {\\it RXTE} results.  The results of points (i) and (ii) above may reflect an actually different spectral shape due to a different emission level from the source with respect to that observed in the March 1997 {\\it RXTE} data. Indeed, Table 1 of Corbet \\& Peele (2001) shows that $\\Gamma$ (which they indicate as $\\alpha$ in their paper), $E_{\\rm cut}$ and $E_{\\rm fold}$ appear to inversely correlate with the source flux, and a simple computation indicates that the source gets harder as the flux increases. Therefore, it appears that there is a switch to harder source spectra when 4U 2206+54 overcomes a threshold luminosity with a value lying somewhere between 1$\\times$10$^{35}$ and 2$\\times$10$^{35}$ erg s$^{-1}$. Unfortunately, in the light of the above data, we cannot say if this transition occurs smoothly or in the form of a ``parameter jump\".  Instead, we believe that the $N_{\\rm H}$ measurement obtained with {\\it BeppoSAX} is substantially different and more reliable thanks to the better spectral coverage at low energies afforded by the {\\it BeppoSAX} LECS. Indeed, if we consider only our data above 2.5 keV, we obtain that $N_{\\rm H}$ = 2.5$^{+0.9}_{-1.0}$ $\\times$10$^{22}$ cm$^{-2}$. Thus, in our opinion the $N_{\\rm H}$ value obtained with {\\it BeppoSAX} should be considered as the correct one. We note that the order-of-magnitude difference in flux between the {\\it BeppoSAX} and the {\\it RXTE} 2--10 keV measurements cannot be explained by the different $N_{\\rm H}$ estimate, which at most may account for a flux difference of about 30\\% only.  This new value of $N_{\\rm H}$, (0.8--0.9)$\\times$10$^{22}$ cm$^{-2}$), compares much better with the Galactic value along the 4U 2206+54 line of sight (0.59$\\times$10$^{22}$ cm$^{-2}$) and with the optical $V$-band reddening value ($A_V$ = 1.6) given by Negueruela \\& Reig (2001): although the empirical relation between $A_V$ and $N_{\\rm H}$ by Predehl \\& Schmitt (1995) implies the presence of further hydrogen local to the source, the difference is now to within a factor of two, and not an order of magnitude as from the previous $N_{\\rm H}$ estimates. This alleviates the conundrum, stressed by Negueruela \\& Reig (2001), of the non-detection of iron emission in presence of very optically thick material around the X--ray source. Concerning this issue, our observation allowed us to put tighter upper limits on the presence of any X--ray Fe emission with respect to the one determined by {\\it EXOSAT} (Gottwald et al. 1995).  The spectrum description by means of a more physically sound model, namely a BB+Comptonization, again points to the presence of a very compact object as the accretor: the spectral shape and the temperatures of the BB and the Compton cloud would suggest that the system most likely harbours a NS. Indeed, the presence of a WD in 4U 2206+54 is basically ruled out because of the X--ray spectral shape, which is completely different from that observed by e.g. Kubo et al. (1998) and Owens et al. (1999) in the system $\\gamma$ Cas, which is thought to host a WD. A better comparison between 4U 2206+54 and systems hosting a WD might come if we consider magnetic cataclysmic variables, such as intermediate polars (see e.g. de Martino et al. 2004 and references therein). However, these objects have spectra with a BB temperature $\\approx$30 times lower than found in 4U 2206+54, and practically no emission detectable above 30 keV. Thus, the high BB temperature of 4U 2206+54 is not compatible with that of a WD surface, which is expected to mainly emit in the UV rather than in the soft X--rays; additionally, the detection of X--ray emission up to 60 keV can hardly be explained by assuming a WD as the accretor in this system.  The same line of thought applies to disfavour a BH interpretation: the Comptonization component has a temperature and an optical depth unusual for a BH in its low-hard state and hosted in a MXRB (e.g. Frontera et al. 2001). So, all the above points to a NS as the accreting object in this system, even if no pulsations were ever detected.  Several observed properties of the source are naturally explained by the accreting NS model. Concerning the X--ray luminosity, accretion onto the NS from a stellar wind emitted by the O9.5V companion star can easily fit the observations: following Frank et al. (1992), if we assume that the companion emits $\\approx$10$^{-6}$ $M_\\odot$ in the form of a wind, one needs to hypothesize an accretion efficiency $\\eta_{\\rm acc} \\approx$ 10$^{-5}$ to produce a luminosity $L \\approx$ 10$^{35}$ erg s$^{-1}$. This value of $\\eta_{\\rm acc}$ may possibly be on the low side of the allowed values for accretion from stellar wind in close systems; however, according to Perna et al. (2003), if corrections to the standard formulae used to estimate the wind accretion rates are introduced, the accretion efficiency drops substantially.  Alternatively, as already suggested by Corbet \\& Peele (2001), partial accretion inhibition due to the ``propeller effect\" (Illarionov \\& Sunyaev 1975; Stella et al. 1986), according to which the magnetosphere of the NS acts as a barrier to accretion of matter onto the NS surface, can be at work. A fraction of the wind matter can nonetheless flow along the magnetic field lines and eventually can reach the NS surface.  As regards the secular X--ray variability over a timescale of $\\approx$1 year, we suggest that this might be due to modulations in the wind density, such as density waves produced by pulsations of the companion star envelope.  In spite of all the above, the NS interpretation rises some problems. In particular, as it evidently appears from Table 2, the BB radius. Clearly, a size of $\\sim$150 m is not acceptable if we assume that the whole NS surface is responsible for the BB X--ray emission. In order to correct for approximations in the BB model application to X--ray data, the hardening factor $f$ (Shimura \\& Takahara 1995), defined as the ratio between the color and the effective BB temperatures, can be introduced. This leads to a corrected BB radius equal to $f^{2} \\cdot R_{\\rm BB}$. However, we should assume that $f \\sim$ 8 to regain the correct BB size for a NS ($\\sim$10 km), while common values for $f$ are around 1.7 (Merloni et al. 2000) and extremes do not exceed $\\sim$3 (Borozdin et al. 1998). A further possible explanation for the small BB emitting area size, assuming isotropic emission from the NS surface, is the following: because of cooling and back-warming effects the spectrum at the NS surface, if fitted with a ``classic\" BB model, can lead to the net result of underestimating the emitting area by as much as 2 orders of magnitude (London et al. 1986).  The alternative to solve this shortcoming is to assume that the emission is not isotropic, i.e. that the accreting matter is either forming a disk around the compact object or is funneled onto the NS magnetic polar caps.  The first possibility appears unlikely because accretion in this system is most probably occurring via stellar wind emitted from the companion star. This comes from assuming that the 9.6 d periodicity determined by Corbet \\& Peele (2001) is the orbital period of the system and that the masses of the two components are 1.4 $M_\\odot$ for the NS and $M_{\\rm *}$~=~19~$M_\\odot$ (Lang 1992) for the O9.5V companion, as spectroscopically identified by Negueruela \\& Reig (2001). With these values, the Roche lobe radius of the companion is $R_{\\rm L} \\sim$ 32 $R_\\odot$, thus much larger than the radius of a O9.5V star, which is $R_{\\rm *}~\\sim$~7.8 $R_\\odot$ (Lang 1992). Thus, because the wind has very low intrinsic angular momentum, a large accretion disk is unlikely to be formed with this accretion mechanism. As remarked in the previous Section, the use of a DBB model instead of a BB produces more unstable fits to our X--ray dataset.  Alternatively, a magnetically-driven accretion scenario can be considered: in this case, one needs the magnetic field of the compact object to be strong enough to form two accretion columns. Indeed, the tentative detection of a CRF indicating the presence of a $\\approx$10$^{12}$ Gauss magnetic field (see Orlandini \\& Dal Fiume 2001) associated with the NS would suggest this possibility.  A further indication that the BB emission is indeed anisotropic (i.e., confined on a fraction of the NS surface) comes from the estimate of the size $r_{\\rm seed}$ of the region emitting the Comptonization seed photons. Following the prescription by in 't Zand et al. (1999) for the computation of $r_{\\rm seed}$ and assuming that the Comptonization seed photons in 4U 2206+54 are produced by the BB (therefore $kT_{\\rm seed} \\sim$ 1.6 keV) we obtain that $r_{\\rm seed} \\sim$ 0.12 km. This estimate is in quite good agreement with our independent determination of the BB emission region radius (see Table 2), thus suggesting that indeed the BB emitting area covers only a fraction of the NS surface.  However, as no X--ray pulsations are detected from this source, the magnetically-driven accretion interpretation needs at least one of the following possibilities to be tenable: an angle between NS spin and magnetic axes close to zero, or a very low inclination angle for the system. In this latter case, even if we assume a non-zero but small (e.g., few degrees) angle between the NS magnetic field and spin axes, the system geometry is such that we could continuously see X--ray radiation from a single polar cap of the NS. Clearly, the emission will not be modulated by the NS rotation in this case also.  This of course means introducing a fine tuning of the system parameters: however, a similar scenario has been proposed to explain the absence of pulsations from 4U 1700$-$37 (White et al. 1983b) and, more recently, from 4U 1700+24 (Masetti et al. 2002), which are believed to host an accreting NS. For 4U 1700+24 Galloway et al. (2002) further supported this description by finding a small amplitude (1 km s$^{-1}$) Doppler periodicity of $\\sim$400 d in their optical spectroscopic data: this period, earlier suggested by Masetti et al. (2002) from timing analysis of {\\it RXTE}/ASM data, is quite likely produced by the orbital motion of the system.  A test for the low-inclination hypothesis of 4U 2206+54 can come by determining, or at least by putting tight constraints on, the orbital Doppler shift of the companion star:  indeed, assuming the system parameters discussed above, we find that the orbital velocity of the companion is v$_{\\rm orb} \\sim$ 20 km s$^{-1}$; we note that this value is consistent with the scatter of the system radial velocities measured by Abt \\& Bautz (1963) from optical spectra.  Thus, summarizing, and despite the problems encountered in the analysis of the observational data on this source, the picture emerging is that 4U 2206+54 is a low-luminosity system composed of a NS and a `normal' blue main-sequence star; the NS is accreting from the wind coming from the companion and is orbiting it in a possibly low-inclination orbit. Tentative evidence of a strong magnetic field from the NS is found, but deep spectroscopic observations, e.g. with {\\it INTEGRAL}, of the hard X--ray tail of this source are needed to confirm (or disprove) this. Long-term variations in the X--ray flux from the source can be explained as due to oscillations in the wind density, possibly induced by slow pulsations of the companion star envelope. This hypothesis can be tested with long-term spectrophotometric monitoring of the optical counterpart."
    },
    "0404/astro-ph0404329_arXiv.txt": {
        "abstract": " ",
        "introduction": "The pioneering idea of the rotation of the Universe should be attributed to G. Gamow \\cite{Gamow46}, who expressed the opinion that the rotation of galaxies is due to the turbulent motion of masses  in the Universe,  and  ``we can ask ourselves whether it is not possible to assume that {\\it all matter in the visible universe is in a status of general rotation around some centre located far beyond the reach of our telescopes?}\". The idea of turbulence as a source of the rotation of galaxies was afterwards developed by  C.F. von Weizsaeker \\cite{Weizsaeker51}, Ozernoy and Chernin \\cite{Ozernoy68}, Ozernoy \\cite{Ozernoy78}, but presently is only of historical value. [If the angular momenta of galaxies had orginated in such a way, their spins should be perpendicular to the  main protostructure plane \\cite{Szandarin74}, which is not observed.] The exact solution of the Einstein equation for the model of a homogeneous universe with rotation and spatial expansion was proposed by Goedel \\cite{Goedel49, Goedel52}. The observational evidence  of global rotation of the Universe was presented by Birch \\cite{Birch82}. He investigated the position angles and polarisation  of classical high-luminosity  double radio sources, finding that the difference between the position angle of the source elongation and of the polarisation are correlated with the source position in the sky. Immediately  there appeared a paper by Phinney and Webster \\cite{Phinney83} concluding that ``the data are unsufficient to substantiate the claim\" and the statistics are  applied incorrectly. Answering, Birch \\cite{Birch83} pointed out the difference in the quantity investigated by him and that by Phinney and Webster, showing that their  data  exhibit the such effect. At the request ofBirch, Phinney and Webster \\cite{Phinney84} reanalysed the data, introducting new ``indirectional statistics\" and taking into account possible observational uncertainties. They  concluded ``that the reported effect ({\\it whatever may be its origin}) is strongly supported by the observations\". Bietenholz and Kronberg \\cite{Bietenholz84} repeated the analysis for a larger sample of objects, finding no effect of the Birch type. New statistical tests were later applied to the data  \\cite{Bietenholz86}.  Nodland and Ralston \\cite{Nodland97a} studied the correlation between the direction and distance to a galaxy and the angle $\\beta$ between the polarisation direction and the major galaxy axis. They found  an effect of a systematic rotation of the plane of polarisation of electromagnetic radiation, which  depends on redshifts. As usually, the result was attacked for the point of incorrectly applied statistics \\cite{Carrol97, Loredo97, Eisenstein97} see the reply \\cite{Nodland97b} with a claim that the new, better data do not support the existence of the effect \\cite{Wardle97}.   The problem of the rotation of the whole Universe  has attracted the attention of several scientists. It was shown that the reported rotation values are too big when compared with the CMB anisotropy. Silk \\cite{Silk70} pointed out that the dynamical effects of a general rotation of the Universe are presently unimportant, contrary to the the early Universe, when  angular velocity $\\Omega \\ge 10^{-13} rad/yr$.  He stressed that now the period of rotation must be greater than the Hubble time, which is a simple consequence of the CMB isotropy. Barrow, Juszkiewicz and Sonoda \\cite{Barrow85} also addressed this question. They showed that cosmic vorticity depends strongly on the cosmological models and assumptions connected with linearisation of homogeneous, anistropic cosmological models over the isotropic Friedmann Universe. For the flat universe, the value is ${\\omega\\over H_0} \\sim 2 \\cdot 10^{-5}$.   Another interesting problem was the discussions on the empirical relation between the angular momentum and mass of celestial bodies $J\\sim M^{5/3}$ \\cite{Brosche86}. Li-Xin Li \\cite{Li98} explained this relation for galaxies as a result of the influence of the global rotation of the Universe on galaxy  formation. ",
        "conclusions": ""
    },
    "0404/astro-ph0404090_arXiv.txt": {
        "abstract": "{\\begin{center} {\\it \\large Abstract\\vspace{-.5em}\\vspace{0pt}} \\end{center}} \\def\\endabstract{\\par} \\usepackage{amsfonts} \\usepackage{amssymb} \\usepackage{verbatim} \\usepackage{graphicx} \\usepackage{hyperref}  \\ifPDF \\hypersetup{% pdftoolbar=treu,% pdfmenubar=true} \\else \\usepackage{epsfig} \\usepackage{pslatex} \\fi  \\setlength{\\hoffset}{-1cm} \\setlength{\\voffset}{-1cm} \\setlength{\\textheight}{23.6cm} \\setlength{\\textwidth}{18.1cm} \\setlength{\\parindent}{0.7cm} \\setlength{\\parskip}{0ex} \\setlength{\\columnsep}{5.4mm} \\setlength{\\topmargin}{1.3cm} \\setlength{\\oddsidemargin}{0.0cm} \\setlength{\\evensidemargin}{0.0cm}  \\renewcommand{\\thesection}{\\rm \\Roman{section}. } \\renewcommand{\\thesubsection}{\\it \\Alph{subsection}. } \\renewcommand{\\thesubsubsection}{\\it  \\arabic{subsubsection}) } \\renewcommand{\\thetable}{ \\arabic{table}} \\renewcommand{\\thefigure}{ \\arabic{figure}}  \\renewcommand\\floatpagefraction{.9} \\renewcommand\\topfraction{.9} \\renewcommand\\bottomfraction{.9} \\renewcommand\\textfraction{.1}  \\newcommand{\\ssection}[1]{% {\\begin{center} {\\it \\large Abstract\\vspace{-.5em}\\vspace{0pt}} \\end{center}} \\def\\endabstract{\\par}  \\setcounter{totalnumber}{50} \\setcounter{topnumber}{50} \\setcounter{bottomnumber}{50}   \\title { \\vspace{-1cm} \\Large The IceCube Project} \\author{\\large C. Spiering for the IceCube Collaboration \\footnote{a full author list is given at the end of this paper}\\\\ \\\\ \\normalsize DESY Zeuthen, Platanenallee 6, D15739 Zeuthen, Germany\\\\ \\normalsize christian.spiering@desy.de}  \\begin{document} \\date{}  This talk gives a brief description of goals, design, expected performance and status of the IceCube project. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404573_arXiv.txt": {
        "abstract": "{We present the oxygen abundance determination for 90 \\ion{H}{ii} regions in the inner parts of the grand design galaxy M101. The abundances were derived employing the $P$ method (Pilyugin 2001a). A comparison is made with previous determinations using another calibration and direct measurements of electron temperature to derive the oxygen abundance. The results show  agreement with the abundances derived from the electron temperature method and also show that the older calibration is not as accurate as the $P$ method. ",
        "introduction": "The determination of  oxygen abundance is a critical stage prior to deriving the value for the metallicity in galaxies and the equivalent abundances for several other elements, such as sulfur, nitrogen or argon. The preferred method for determining the oxygen abundance in galaxies using H~{\\sc ii} regions is through electron temperature-sensitive lines (the so-called $T_{\\rm e}$ method), such as the [O~{\\sc iii}]$\\lambda$4363 or [\\ion{O}{iii}]$\\lambda$7325 auroral lines (Searle 1971; Rosa 1981; Garnett \\& Kennicutt 1994; Kennicutt et al.\\ 2003). However, these lines are not always available: for oxygen-rich regions, the oxygen line [O~{\\sc iii}]$\\lambda$4363 is weak and difficult to detect, so there are not many direct abundance determinations from the inner parts of galaxies.  Other methods are based on ``empirical'' calibrations of metallicity employing strong-line abundance estimators. These methods are based on direct measurements of the electronic temperature of low metallicity regions and in theoretical models for high metallicity regions. One method with widespread acceptance and use is the $R_{23}$-method, suggested by Pagel et al. (1979). It is based in the oxygen line ratio, $R_{23}=([{\\rm O II}]\\lambda\\lambda 3726,29 + [{\\rm O III}]\\lambda\\lambda 4959,5007)/{\\rm H}\\beta$. There are different calibrations using the $R_{23}$ ratio, such as those of Dopita \\& Evans (1986), Edmunds \\& Pagel (1984), McCall et al. (1985), McGaugh (1991) or Zaritsky et al.\\ (1994). However, this indicator presents one great disadvantage: the derived abundances depend strongly  on the $R_{23}$-\\ion{O}/\\ion{H} calibration (Kewley \\& Dopita 2002; Cedr\\'es 2003).  Moreover, for \\object{M101} Kennicutt et al.\\ (2003)  have recently found systematic differences up to a factor 3 between abundances derived from some empirical calibrations and those derived from the direct method, the latter being lower.  Smartt et al.\\ (2001) and Trundle et al.\\ (2002) have shown that in the Local Group spiral galaxy M31  oxygen abundances of B supergiant atmospheres are also systematically lower than those obtained by classical $R_{23}$-\\ion{O}/\\ion{H} calibrations. However, Monteverde et al.\\ (2000) found very good agreement between B supergiant abundances (obtained in a similar way as in the aforementioned references) and the abundances derived from the $T_{\\rm e}$ method in the interstellar medium in M33. Thus, it is clear that more comparisons are required.  New methods for  abundance determinations using strong lines have been developed recently. These methods achieve a better approximation to the results obtained with the $T_{\\rm e}$ method. One of these new calibrations, the $P$ method, is proposed by Pilyugin (2001a, 2001b).  In this paper we present an estimate of the oxygen abundance for the inner H~{\\sc ii} regions of M101 using the $P$ method with data from direct imaging observations, which give us a larger number of regions when compared with spectroscopic methods, these allowing  better sampling of the disc both spatially and in \\ion{H}{ii} region luminosity, and with less telescope time. ",
        "conclusions": "We have obtained the oxygen abundance through the $P$ method for 90 H~{\\sc ii} regions of the inner parts of M101.  There is very good agreement between our data and the derived abundances from the $T_{\\rm e}$ method for regions with high metallicity. Compared with spectroscopic methods, these results present a larger number of regions than any previous study. The dispersion of the data shows that, even for the larger uncertainties, the data are almost as reliable as those of spectroscopic studies employing the $p$ or the $T_{\\rm e}$ method. Moreover, the direct imaging method is less time-consuming because only two observing nights are required to obtain data for more than a hundred regions from one galaxy.  The $P$ method has proved to be a useful tool for determining  oxygen abundances in the inner zones of galaxies, where auroral lines are difficult to measure and the metallicity is moderately high. Moreover, direct imaging techniques proved  superior when considering observing time and the number of  regions observed."
    },
    "0404/astro-ph0404059_arXiv.txt": {
        "abstract": "We present results from our optical survey of the position error contours (``error boxes'') of unidentified EGRET sources at mid to high Galactic latitude. It is our intention to search for potential blazars that may have been missed in the original identification process of the three EGRET Catalogues and supplementary publications. We have first searched the error contours of unidentified sources at $b > \\vert 20^\\circ \\vert$ for flat spectrum radio sources using the NASA Extragalactic Database (NED). For each such radio source found we conducted optical searches for counterparts using the Palomar 60-inch telescope, and University of Wyoming's 2.3 and 0.6 m telescopes. Many of the radio sources have plausible optical counterparts, and spectroscopy will be conducted at a later date to determine which of these sources are quasars or active galaxies. Results show thats several sources are optically variable, and/or have flat or inverted radio to millimeter spectra and are thus potential blazars. ",
        "introduction": "Though the EGRET instrument detected 271 sources, most of these sources remain unidentified with counterparts at any other wavelength \\citep{hartman99}. However, the distribution of these sources on the sky does suggest at least two major populations: Those near the Galactic plane, where 90 \\% of the sources are unidentified, and those above the plane, where 50 \\% are unidentified \\citep{caraveo01}. \\citet{caraveo01} further suggest that those sources above the plane may have at least three subcomponents: extragalactic, Gould Belt, and Galactic halo. Though nearly all of the firmly identified extragalactic sources are blazars, there have been some recent studies suggesting that some sources are related to non-blazar AGN \\citep{mukherjee02} or galaxy clusters \\citep{colafrancesco02}. The task of investigating all of these sources seems daunting. Therefore, several investigators have gone the route of examining several individual sources in depth. An advantage of this approach is that such searches are like to yield abundant information on the several sources they investigate; however, it is more difficult to draw general conclusions from such results. We continue along these lines by searching within the error contours of high Galactic latitude EGRET unidentified sources for new blazars. We will then ascertain whether any such objects found are likely to be the source of the gamma rays. At the time of submission, we became aware of a parallel study by \\citet{sowards03}. Though similar in scope, ours studies differ, and are complimentary, in that we are continuing multi-epoch optical observations with B, V, R photometry and do include some sources at southern declinations. ",
        "conclusions": "We have searched the ``error boxes'' of EGRET sources, and found several potential blazar counterparts at other wavebands. The most compelling case, newly reported here, is that of 3EG J0038-0949 (matched with NVSS J993906-094247). These counterparts are generally characterized by flat or inverted radio spectra (in some cases up to 37 GHz), and in some cases, long and short term optical variability. We conclude that several of these objects are the source of the gamma rays. This study and others have shown that the EGRET instrument detected a class of dim blazars, in addition to those identified in the Third EGRET Cataloge (\\citep{hartman99}. NASA's upcomming Gamma Ray Large Area Telescope (GLAST) mission should confirm the existence of a relatively dim population of blazars thatare the source of gamma rays for many of the unidentified sources at high Galactic latitude."
    },
    "0404/astro-ph0404234_arXiv.txt": {
        "abstract": "{\\it Chandra} observations of distant ($D\\sim 10$ Mpc) elliptical galaxies have revealed large numbers of Low Mass X-ray Binaries (LMXBs) accreting at $\\dot M > 10^{-8} M_\\odot$ yr$^{-1}$.  The majority of these LMXBs reside in globular clusters (GCs) and it has been suggested that many of the field LMXBs also originated in GCs. We show here that ultracompact binaries with orbital periods of 8-10 minutes and He or C/O donors of $0.06-0.08M_\\odot$ naturally provide the observed $\\dot M$'s from gravitational radiation losses alone. Such systems are predicted to be formed in the dense GC environment, a hypothesis supported by the 11.4 minute binary 4U 1820-30, the brightest persistent LMXB in a Galactic GC.  These binaries have short enough lifetimes ($<3\\times 10^6$ years) while bright ($L>10^{38} {\\rm erg \\ s^{-1}}$) that we calculate their luminosity function under a steady-state approximation. This yields a luminosity function slope in agreement with that observed for $6\\times 10^{37}{\\rm erg \\ s^{-1}}<L<5\\times 10^{38} {\\rm erg \\ s^{-1}}$, encouraging us to use the observed numbers of LMXBs per GC mass to calculate the accumulated number of ultracompact binaries.  For a constant birthrate over 8 Gyrs, the number of ultracompact binaries which have evolved through this bright phase is $\\sim 4000$ in a $10^7M_\\odot$ GC, consistent with dynamical interaction calculations. Perhaps most importantly, if all ultracompacts become millisecond radio pulsars, then the observed normalization agrees with the inferred number of millisecond radio pulsars in 47 Tuc and Galactic GCs in general. ",
        "introduction": "The {\\it Einstein} observatory allowed for the discovery and study of X-ray emission from elliptical and S0 galaxies (e.g. Forman, Jones \\& Tucker 1985), finding that the emission likely had two distinct contributors (Trinchieri \\& Fabbiano 1985, Canizares, Fabbiano, \\& Trinchieri 1987).  The softer emission is from hot interstellar gas (Forman et al 1985, Trinchieri, Fabbiano, \\& Canizares 1986) and dominates the total X-ray emission from the most massive ellipticals.  Observations with ASCA (Matsumoto et al. 1997) showed that a harder (5-10 keV) spectral component scaled linearly with the bolometric luminosity of the galaxy, confirming its suggested origin as the unresolved contributions of numerous Low Mass X-Ray Binaries (LMXBs; Trinchieri \\& Fabbiano 1985; White, Nagase, \\& Parmar 1995; Irwin \\& Sarazin 1998).  The exquisite point source sensitivity of {\\it Chandra} led to the discoveries of the expected LMXBs (e.g. Sarazin, Irwin \\& Bregman 2000; Blanton, Sarazin \\& Irwin 2000; Sarazin, Irwin, \\& Bregman 2001; Finoguenov \\& Jones 2001; Angelini, Loewenstein, \\& Mushotzky 2001; Kraft et al. 2001; Kundu, Maccarone, \\& Zepf 2002; Di Stefano et al. 2003; Randall, Sarazin \\& Irwin 2004).  These observations (see Gilfanov 2004 and Kim \\& Fabbiano 2004) typically find 20-50 LMXBs with $L> 10^{38} {\\rm erg \\ s^{-1}}$ from elliptical galaxies with stellar masses of $(1-3)\\times 10^{11}M_\\odot$.  This has provided a large sample of very bright LMXBs that is unattainable from the Milky Way or even M31, allowing for a new probe of binary evolution.  Gilfanov (2004) showed that the LMXBs of these elliptical and S0 galaxies are consistent with a single luminosity function \\begin{equation} {dN\\over dL}\\propto {1\\over L^{\\alpha}}, \\end{equation} from $10^{37} {\\rm erg \\ s^{-1}}$ to a break value $L_B=5^{+1.4}_{-0.7}\\times 10^{38} {\\rm erg \\ s^{-1}}$.  with $\\alpha=1.64\\pm 0.22$.  Kim and Fabbiano (2004) recently found a similar result for $L>6\\times 10^{37} {\\rm erg \\ s^{-1}}$ of $\\alpha=1.80\\pm 0.2$ and $L_B=4.8\\pm 1.2 \\times 10^{38} {\\rm erg \\ s^{-1}}$.  For $L>L_B$, the luminosity function steepens to $\\alpha=2.7\\pm 0.5$.  The total number of X-ray sources scales with the galactic mass, with $\\approx 20$ sources of $L>10^{38} {\\rm erg \\ s^{-1}}$ per $10^{11}M_\\odot$ (Gilfanov 2004), or a total X-ray luminosity that scales linearly with the $K$ band luminosity as $L({\\rm LMXB})=1.5\\pm 0.6 \\times 10^{40} {\\rm erg \\ s^{-1}}$ for $L_K=10^{11}L_{K,\\odot}$ (Kim \\& Fabbiano 2004).  Starting with the original observation of NGC 4697 (Sarazin et al. 2001), it has become clear that many (20-70\\%) of these LMXBs are residing in globular clusters (Angelini et al. 2001; Kundu et al. 2002; Maccarone, Kundu \\& Zepf 2003; Minnitti et al. 2004). Kundu et al. (2002) showed that there was no difference in $dN/dL$ for LMXBs in and out of GCs (both had $\\alpha=1.55\\pm 0.15$ and $L_B\\approx 3\\times 10^{38} {\\rm erg \\ s^{-1}}$) in NGC 4472 , supporting the suggestion of White, Sarazin and Kulkarni (2002) that most of the bright LMXBs in ellipticals are made in GCs. Kundu et al. (2002) showed that nearly 4\\% of NGC 4472 GCs host a bright LMXB, with metal rich clusters about 3 times more likely to host an LMXB (confirmed in NGC 4365 and NGC 3115 by Kundu et. al. 2003). Sarazin et al. (2003) studied four galaxies and found that the specific incidence of LMXBs in GCs is about one source with $L>10^{38} {\\rm erg \\ s^{-1}}$ per $10^7 L_{\\odot,I}$, so that a star in a GC is about 1000 times more likely to be a donor than a star in the field.  Such a large enhancement of the incidence of LMXBs in GCs was first found in the Milky Way (Katz 1975; Clark 1975) and is indicative of the important role of interactions in creating mass transferring binaries in GCs (see Hut et al. 1992 for an overview).  Though numerous LMXBs have been found, the nature of the donor star in this old stellar population remains a mystery. Piro \\& Bildsten (2002) showed that for field LMXBs, the simplest way to reach $10^{38} {\\rm erg \\ s^{-1}}$ (or $\\dot M=LR/GM\\approx 10^{-8} M_\\odot$ yr$^{-1}$  for accretion onto a neutron star of $M=1.4M_\\odot$ and $R=10 {\\rm \\  km}$) is to have a red giant branch star fill the Roche lobe with orbital periods of days or longer. They noted that these wide binaries are nearly always transient accretors, and that multiple {\\it Chandra} observations would easily identify them (e.g. Kraft et al. 2001). However, calculating the resulting $dN/dL$ is impossible given our poor state of knowledge of the transient duty cycle.  We argue here that the most likely type of mass transferring binary responsible for the bright end of the luminosity function in distant elliptical galaxies (especially those in GCs) are ``ultracompact'' systems consisting of either a He or C/O white dwarf donor of mass $M_c\\approx 0.04-0.08M_\\odot$ which is filling its Roche lobe in a 5-10 minute orbit with a neutron star. In \\S 2, we summarize the observational clues that led us to this conjecture and derive the resulting $dN/dL$, showing that the observed power law is naturally explained. In \\S 3, we use the observed normalizations from {\\it Chandra} to compare to the predicted formation rate of ultracompacts in GCs from dynamical interactions. We also predict the expected number of millisecond radio pulsars (MSPs) assuming that all ultracompacts become MSPs. We close in \\S 4 by noting more observational tests, and discussing future work. ",
        "conclusions": "Our identification of these distant LMXBs with ultracompact binaries naturally explains the observed luminosity function from {\\it Chandra} observations and yields the LMXB birthrate of one new mass transferring binary every $2\\times 10^6$ years per $10^7M_\\odot$ of GCs. Rather remarkably, this derived birthrate from distant ellipticals agrees with both dynamical calculations (e.g. Ivanova \\& Rasio 2004) and the observed number of MSPs in galactic GCs, especially 47 Tuc. The much lower donor mass ($\\approx 0.06M_\\odot$) has alleviated the ``birthrate'' problem often discussed for LMXBs and MSPs in galactic GCs (e.g. Kulkarni et al. 1990).  The simplest way to prove this hypothesis is to find the 5-10 minute orbital periods. If we use 4U~1820-30 as our example, the level of orbital variability in the X-rays could be as low as a few percent, allowing for {\\it Chandra} searches amongst the few bright LMXBs in M31 GCs while {\\it XMM-Newton } could probe much deeper. The source 4U~1820-30 cycles in luminosity by about a factor of 3 over a 171 day cycle for unknown reasons (Chou \\& Grindlay 2001).  Such variability is easily detected by {\\it Chandra}, and indeed variability at this level has been reported for bright LMXBs in M31 (e.g. Trudolyubov \\& Priedhorsky 2004). However, attributing such behavior as unique to ultracompacts is much harder. Ultraviolet observations of 4U~1820-30 (King et al 1993; Arons and King 1993) confirmed the 11.4 minute orbit and similar work could be done with {\\it HST} amongst the bright systems in M31.  The lower luminosities that are visible in the GC sources in our galaxy and in M31 have yet to be probed by {\\it Chandra} in distant ellipticals. For ultracompact binaries, the expectation is that our derived $dN/dL$ will continue until a $L$ is reached where the systems become transients (presuming no dramatic episode associated with MSP turn-on). DB showed that in the absence of X-ray heating in the outer disk, this would occur at $L\\sim 10^{37} {\\rm erg \\ s^{-1}}$, whereas with X-ray heating, the disks can remain stable down to much lower X-ray luminosities. Unfortunately, the current populations of ultracompacts in our galaxy don't provide a stringent test of which case is correct. {\\it Chandra} may be able to identify this cutoff with deeper observations of nearby ellipticals.  There is still much to explain. Clearly, not all bright LMXBs are ultracompact binaries, both because the MSPs in wide binaries with He WDs (orbital periods longer than a day) cannot be made from ultracompacts and because roughly half the galactic GC sources are clearly hydrogen accretors (Kuulkers et al. 2003). However, the strong expectation is that most of the H donors with $L$ large enough to detect at 10 Mpc are transient accretors (e.g. Piro \\& Bildsten 2002), making secure predictions of $dN/dL$ difficult. Observations of the bright LMXBs in M31 (Di Stefano et al. 2002; Trudolyubov \\& Priedhorsky 2004) have motivated a stable mass transfer scenario involving thermal timescale mass transfer from evolved stars (Di Stefano et al. 2002), but $dN/dL$ was not derived. We have also not explained the higher incidence of bright LMXBs in metal-rich GCs (Kundu et al. 2002; Maccarone et al. 2004)."
    },
    "0404/astro-ph0404528_arXiv.txt": {
        "abstract": "We describe an analysis of the time-resolved measurements of the surface magnetic field in the roAp star \\equ. We have obtained a high-resolution and high S/N spectroscopic time-series, and the magnetic field was determined using Zeeman resolved profiles of the \\fet\\ 6149.25~\\AA\\ and \\feo\\ 6173.34~\\AA\\ lines. Contrary to recent reports we do not find any evidence of magnetic variability with pulsation phase, and derive an upper limit of 5--10~G for pulsational modulation of the surface magnetic field in \\equ. ",
        "introduction": "\\label{introduction}  After discovery of the conspicuous radial velocity (RV) pulsational variations in a sample of rapidly oscillating magnetic peculiar (roAp) stars (Kanaan \\& Hatzes \\scite{KH98}, Savanov, Malanushenko \\& Ryabchikova \\scite{SMR99}, Kochukhov \\& Ryabchikova \\scite{KR01a} for \\equ; Baldry et al. \\scite{BBV98}, Baldry \\& Bedding \\scite{BB00}, Kochukhov \\& Ryabchikova \\scite{KR01b} for \\cir\\ and HD~83368), attempts to search for magnetic field variations over the pulsational period have been made. First, Hubrig et al. \\cite{HKB04}  tried to measure pulsational variability of the longitudinal magnetic field \\bz\\ in six roAp stars. Their sample included \\equ\\ -- probably one of the most favourable stars for this kind of investigation. \\equ\\ is a bright northern roAp star with a strong magnetic field, and with one of the largest pulsational RV amplitudes which exceeds 1000~\\ms\\ in individual spectral lines. The extremely slow rotation of \\equ, leading to very sharp spectral lines, makes this star the best candidate for any study of the pulsational variability in spectroscopy. Hubrig et al. \\cite{HKB04} used low-resolution Zeeman time-series observations and measured \\bz\\  using hydrogen lines and unresolved blends of metal lines. They failed to detect any variability beyond the formal errors of their measurements which were 40--100~G.  According to a coarse theoretical estimate made by Hubrig et al. \\cite{HKB04} there should exist a linear relation between magnetic field variability over the pulsational cycle and the RV amplitudes.  In roAp stars the largest RV amplitudes are observed in the lines of first and second ions of rare-earth elements (REE), while they are usually below measurement errors for the lines of iron group elements. Thus, high-resolution spectroscopy and spectropolarimetry of selected REE spectral lines is a more promising tool for an investigation of possible rapid magnetic oscillations in roAp stars. Taking this into account Leone \\& Kurtz \\cite{LK03} obtained a high-resolution (R=115\\,000), high S/N time-series of observations of \\equ\\ with a circular polarization analyzer, and measured \\bz\\ using four \\nd\\ lines. They reported the discovery of pulsational variations with amplitudes between 112 to 240~G and, more surprisingly, with discrepant phases of magnetic maximum for different \\nd\\ lines. Leone \\& Kurtz's result was based on only 18 time-resolved spectra.  A year later Kochukhov, Ryabchikova \\&  Piskunov \\cite{KRP04} obtained a time-series of polarimetric observations of \\equ\\ with a smaller resolving power (R=38\\,000), but acquired more than 200 spectra over 3 nights, more than compensating for lesser quality of individual spectra. Kochukhov et al. \\cite{KRP04} used simultaneously 13 \\nd\\ lines for magnetic measurements which allowed them to achieve a substantial reduction of the error of the \\bz\\ determinations. They did not confirm longitudinal field variability over the pulsational period in \\equ\\ and gave a conservative upper limit of $\\approx$40~G for the amplitude of pulsational modulation of \\bz\\  determined from \\nd\\ lines.  Another attempt to search for possible rapid magnetic variability in \\equ\\ was made by Savanov, Musaev \\& Bondar \\cite{SMB03}. They measured the surface magnetic field \\bs\\ variations over the pulsational period using the \\fet\\ 6149.25~\\AA\\ line observed in unpolarized light. Due to a very simple Zeeman pattern (two equally separated $\\pi$- and $\\sigma$-components) this line is ideal for \\bs\\ measurements (see Mathys et al. \\scite{MHL97}). Savanov et al. reported a 1.8$\\sigma$ detection of \\bs\\ variability with an amplitude of 99$\\pm$53~G. At the same time they did not find periodic variations of RV measured for the individual Zeeman resolved components of the \\fet\\ 6149.25~\\AA\\ doublet exceeding their error limit (100--120~\\ms). The authors used high-resolution R=120\\,000 time-series observations, but the S/N of a single spectrum did not exceed 40--60. Clearly, the result of Savanov et al.~\\cite{SMB03} is marginal and needs to be confirmed or rejected with data of better quality.  In this paper we present the results of a new search for pulsational variations of \\bs\\ in \\equ\\ using high-resolution and high S/N time-resolved observations of this star. These observational data allowed us to obtain precise measurements of the magnetic field in \\equ, and strongly constrain possible changes of \\bs\\ over the pulsation cycle of the star.  \\begin{figure*} \\fifps{14.5cm}{me113_f2.eps} \\caption{Measurements of radial velocity and surface magnetic field obtained for \\equ\\ on 22/07/99. The \\nd\\ 6145.07~\\AA\\ pulsational RV curve (a) is folded with the best-fit oscillation period and is compared (c) with the variation of \\bs\\ measured from the separation of the Zeeman components of the \\fet\\ 6149.25~\\AA\\ (centre-of-gravity line position measurements). Panels (b) and (d) illustrate the amplitude spectra for the RV and \\bs\\ (thick line -- centre-of-gravity measurements, thin line -- multiple fit of gaussians). The vertical dashed lines show the photometric pulsational periods of \\equ\\ (Martinez et al. 1996).} \\label{fig2} \\end{figure*} ",
        "conclusions": "Our time-resolved magnetic measurements of \\equ\\ have achieved the highest precision for a roAp star, but reveal no evidence of pulsational modulation of the field strength. We constrain possible magnetic variability to be below $\\approx$5~G in Fe lines. These results suggest that the marginal detection of $\\approx$100~G \\bs\\ variability during the pulsation cycle  reported in \\equ\\ by Savanov et al. \\cite{SMB03} is spurious, and probably stems from insufficient precision of the \\bs\\ measurements in that study.  The null result reported in the present paper complements non-detection of the pulsational  variability of \\bz\\ determined from \\nd\\ lines (Kochukhov et al. \\scite{KRP04}). It should be recalled that the Fe lines studied in our paper and the REE lines showing strong pulsational RV modulation are formed at substantially different atmospheric depths due to the extreme stratification of chemical abundances in cool Ap stars and in \\equ\\ in particular. Stratification analysis presented by Ryabchikova et al. \\cite{RPK02} allows us to conclude that \\nd\\ lines sample very high atmospheric layers with optical depths $\\log\\tau_{5000}\\la-7$, while the Fe  lines are formed below $\\log\\tau_{5000}\\approx-1$. The striking difference in the RV amplitudes of Fe and \\nd\\ lines is then attributed to an outward increase of pulsational amplitude by roughly a factor of 50--100. This increase is none the less not accompanied by an increase or even the presence of any observable oscillations of the magnetic field structure. Combining the results of this paper with the study of Kochukhov et al. \\cite{KRP04}, and taking into account that for \\equ\\ $B_s\\approx  2.6 B_\\ell$, we find that at all observed atmospheric depths \\bs\\ changes by less than about $0.5$~G per \\ms\\ of corresponding velocity oscillations. Hence, possible magnetic variability is constrained to be below 1\\% of the field strength.  We note that \\equ\\ is distinguished among the roAp stars by its strong magnetic field, sharp lines and the exceptionally high amplitude of pulsational line profile  changes. Yet this star defies any attempts to detect magnetic variability with  pulsation phase. This suggests that rapid magnetic modulation may be even more  difficult to detect in other roAp pulsators. The outcome of our monitoring of \\bz\\ and \\bs\\ in \\equ\\ demonstrates that very accurate measurements of the Zeeman resolved lines can yield more precise magnetic time-series compared with the polarimetric \\bz\\ observations. Therefore, the most promising direction for future attempts to detect magnetic oscillations in \\equ\\ (which are in any case unlikely to exceed a few tens of gauss) would be to observe those \\ion{Nd}{ii}, \\nd\\ and \\pr\\ lines which are characterized by large pulsational RV shifts and at the same time show Zeeman resolved profiles. Unfortunately, the extra broadening of the pulsating lines (see Kochukhov \\& Ryabchikova \\scite{KR01a}) strongly smears observed Zeeman structure and makes proposed study of \\bs\\ oscillations extremely difficult."
    },
    "0404/astro-ph0404372_arXiv.txt": {
        "abstract": "Near-infrared and optical imaging of BL Lac host galaxies is used to investigate their colour properties. We find that the $R$--$H$ colour and colour gradient distributions of the BL Lac hosts are much wider than those for normal ellipticals, and many objects have very blue hosts and/or steep colour gradients. The blue colours are most likely caused by recent star formation. The lack of obvious signs of interaction may, however, require a significant time delay between the interaction event with associated star formation episodes and the onset of the nuclear activity. ",
        "introduction": "Optical imaging of low redshift (z $<$ 0.5) BL Lac objects (e.g. Falomo \\& Kotilainen 1999; Urry et al. 2000) has shown that virtually all of them are hosted in luminous ellipticals, with characteristics indistinguishable from those of inactive massive ellipticals. On the other hand, near-infrared (NIR) imaging has only been available for small samples of BL Lacs (Kotilainen et al. 1998 [K98]; Scarpa et al. 2000 [S00]; Cheung et al. 2003 [C03]). We present deep high spatial resolution ($\\sim$0.7 arcsec FWHM) NIR $H$--band (1.65 $\\mu$m) imaging of 23 low redshift (z $<$ 0.3) BL Lacs that were previously investigated in the optical $R$-band (references above). We combine them with previous data to form a sample of 41 BL Lacs with which to investigate the optical-NIR colour properties of the BL Lac hosts and to compare them with radio galaxies (RG) and inactive ellipticals. Full report is given in Kotilainen \\& Falomo (2004). ",
        "conclusions": ""
    },
    "0404/astro-ph0404144_arXiv.txt": {
        "abstract": "A number of Anomalous X-ray Pulsars (AXPs) have recently been detected in the optical/IR wavelengths. We use their inferred brightness to place general constraints on any model for this emission within the magnetar framework. We find that neutron-star surface emission cannot account for the observations and that the emission must be magnetospheric in origin. We propose a model for the optical/IR emission in which a distribution of energetic electrons in the neutron-star magnetosphere emits synchrotron radiation. This model can naturally reproduce the observed brightness and the rising spectra of AXPs as well as the observed pulsations at the stellar spin frequency and the correlation of the IR flux with their bursting activity. ",
        "introduction": "Magnetars are a class of neutron stars powered by the decay of their ultrastrong magnetic fields (Duncan \\& Thompson 1992). If they are isolated, magnetars are believed to appear as persistent pulsars in the X-rays, with occasional episodes of bursting activity in the hard X-rays/soft $\\gamma$-rays. Anomalous X-ray Pulsars (AXPs) and Soft Gamma-ray Repeaters (SGRs) are thought to be the observational manifestations of magnetars: their steady spin-down, quasi-thermal X-ray spectra, bursts, and the absence of binary companions all lend support to this identification (see, e.g., Kouveliotou et al.\\ 1998; Mereghetti, Israel, \\& Stella 1998; Woods et al.\\ 1999; \\\"Ozel, Psaltis, \\& Kaspi 2001; Gavriil \\& Kaspi 2002; Gavriil, Kaspi, \\& Woods 2003). The absence, until recently, of detectable emission from AXPs and SGRs in longer wavelengths constrained all such studies to the X-rays and $\\gamma$-rays.  In the last few years, faint counterparts of four AXPs have been detected in IR and optical wavelengths (Hulleman, van Kerkwijk, \\& Kulkarni 2000, 2004; Hulleman et al.\\ 2001; Israel et al.\\ 2002, 2003; Wang \\& Chakrabarty 2002) but no counterparts have yet been observed for SGRs, due at least in part to the high extinction along the line of sight to these sources (Corbel et al.\\ 1997, 1999; Vrba et al.\\ 2000; Kaplan et al.\\ 2001, 2002; Eikenberry et al.\\ 2001; Wachter et al.\\ 2004). The AXP counterparts were found to be variable, with fluxes possibly correlated with the bursting activity (Kaspi et al.\\ 2003), had rising spectra in $\\nu F_\\nu$ (Hulleman et al.\\ 2004; see Fig.~1 below), and showed pulsations at the stellar spin frequency (Kern \\& Martin 2002). These long-wavelength detections have been used to argue strongly against the presence of accretion disks and companion stars around AXPs (Hulleman et al.\\ 2000; Perna, Hernquist, \\& Narayan 2000). However, no predictions have been made so far for the expected long-wavelength emission from a magnetar, which can be tested against observations (see Eichler, Gedalin, \\& Lyubarsky 2002 for a preliminary discussion).  In this {\\it Letter}, we discuss the strong constraints imposed by the observed IR and optical brightness of AXPs on any magnetar emission model. In particular, we argue that the long-wavelength emission cannot originate from the surface of the neutron star. We suggest that synchrotron emission from energetic particles in the magnetosphere at $\\gtrsim 50$ stellar radii is responsible for this emission. ",
        "conclusions": "In this {\\em Letter}, we described a model of the optical/IR emission of AXPs in which a distribution of energetic electrons in the neutron star magnetosphere emits synchrotron radiation. This model can naturally reproduce the observed brightness and the rising spectra of AXPs.  In the calculations reported above, we have made some simplifying assumptions that may affect the quantitative predictions of the model. First, we did not take into account the angular dependence of the magnetic field and of the emitted radiation. Relaxing this assumption will result in a brightness and spectrum that depends on the relative orientation of the magnetic axis and the observer. Moreover, it will naturally give rise to pulsations of the long-wavelength emission at the stellar spin frequency. Second, we have assumed a mono-energetic distribution of electrons throughout the magnetosphere. Allowing for a radial profile of the electron Lorentz factor $\\gamma$ will affect the slope of the optical/IR spectrum and may give rise to spectral breaks at frequencies characteristic of the acceleration and cooling mechanisms in the magnetosphere.  The source of energy and the mechanism of particle acceleration in the magnetosphere of a magnetar are open questions. The long-wavelength observations of AXPs may provide within this model clues towards distinguishing between various possibilities. If the source of this radiation is dissipation of the rotational energy of the neutron star, then equation~(2) provides an upper limit on the flux of magnetospheric synchrotron photons (see the dashed line in Fig.~2). This, in return, sets an upper limit on the frequency of synchrotron radiation at $\\simeq 10^{15}$~Hz for the parameters of 4U~0142$+$61. Combining this limit with equation~(6) and the best-fit value of the parameter $B_{\\rm NS} \\gamma^4$ (see Fig.~2) results in a lower limit on the emission radius of $r/R_{\\rm NS} \\gtrsim 80 (B_{\\rm NS}/10^{14}$~G$)^{5/12}$. The observations of 4U~0142$+$61, in this framework, require that either the particles do not radiate at smaller radii or that particle acceleration takes place at large distances from the neutron star surface, possibly due to the formation of bound positronium states (Leinson \\& P\\'erez 2000) at the stronger magnetic fields close to the star.  The inferred IR spectrum of 4U~0142$+$61 is flatter than its optical spectrum (Hulleman et al.\\ 2004). This can be achieved for a radial distribution of Lorentz factors that rises locally with increasing radius. Such an energy distribution would strongly suggest that the particles are accelerated at $\\gtrsim 100 R_{\\rm NS}$ and lose energy via synchrotron radiation as they follow field lines back toward the star. If the low-frequency radiation is powered by magnetic energy, such a configuration may naturally arise from the dissipation of Alfv\\'en waves in the magnetospheres of magnetars (Thompson \\& Duncan 1996). Because of the much larger reservoir of magnetic energy, the optical spectrum may extend to higher frequencies than in the case of rotation power. Observations of AXPs and SGRs in UV wavelengths will be important in distinguishing between different mechanisms. Note also that the optical and IR observations of 4U~0142$+$61 were not carried out simultaneously and may point to a change in the magnetic field configuration in time. Simultaneous and/or repeat observations may also provide additional clues to the nature of the emission.  This magnetospheric model can also account for a number of other observed properties of the optical/IR emission of AXPs. As mentioned above, the magnetic field geometry and the angular dependence of synchrotron emission gives rise to pulsations at the stellar spin frequency as observed (Kern \\& Martin 2002). The optical/IR flux is also likely to be polarized due to the strong beaming of synchrotron emission but a quantitative prediction for the degree of polarization requires including in the calculations the angle dependences discussed above. Observations that search for polarization of the optical emission may help further constrain the magnetospheric model. In addition, the optical/IR emission will naturally respond to the changes in the magnetic field strength, configuration, or spin-down rate following an SGR-like burst. The timescale for this flux enhancement will be dictated by the rearrangement of magnetic field lines which are anchored in the neutron star crust and thus will proceed at a slower pace than any characteristic timescale in the magnetosphere. Such correlated changes have been observed following the burst of 1E~2259$+$586 (Kaspi et al.\\ 2003).   \\begin{figure}[t] \\centerline{ \\psfig{file=Fig3.eps,width=11truecm} } \\figcaption[]{The ratio of the unabsorbed IR flux in the K band to the X-ray flux at $5 \\times 10^{17}$~Hz as a function of the rotational luminosity for four persistent AXPs. Normalizing to the X-ray flux minimizes the differences in distance uncertainties and absorption column between the sources. The positive correlation supports the idea of magnetospheric origin for the long-wavelength emission. \\label{Fig:rot}} \\end{figure}  Finally, because of its proposed magnetospheric origin, the low-energy emission of AXPs is expected to correlate with the rate of rotational or magnetic energy losses. The absolute luminosities of AXPs in the IR are very hard to infer observationally because of the unknown distances and the effects of the large interstellar absorption to these sources. To minimize these uncertainties, we have normalized the unabsorbed IR flux in the K band to the unabsorbed persistent X-ray flux at $5 \\times 10^{17}$~Hz. The intrinsic X-ray luminosities of AXPs are expected to be similar between the sources and not to depend strongly on the magnetic field strength. Even though such a luminosity clustering is difficult to infer from the data, their quasi-thermal spectra with similar color temperatures and the fact that the emitting areas should all be comparable to the surface area of a neutron star point to such a theoretical expectation. The inferred flux ratios are plotted against the observed rate of rotational energy loss in Figure~3 (see Israel et al.\\ 2003 and references therein for the data). The overall correlation provides additional support to the model of magnetospheric optical/IR emission of AXPs discussed here."
    },
    "0404/astro-ph0404308_arXiv.txt": {
        "abstract": "{We selected a sample of eight bright unobscured (at least at the iron line energy) Seyfert Galaxies observed simultaneously by XMM-\\textit{Newton} and BeppoSAX, taking advantage of the complementary characteristics of the two missions. The main results of our analysis can be summarized as follows: narrow neutral iron lines are confirmed to be an ubiquitous component in Seyfert spectra; none of the analyzed sources shows unambiguously a broad relativistic iron line; all the sources of our sample (with a single exception) show the presence of a Compton reflection component; emission lines from ionized iron are observed in some sources; peculiar weak features around 5-6 keV (possibly arising from rotating spots on the accretion disk) are detected in two sources. The scenario emerging from these results strongly requires some corrections for the classical model of reprocessing from the accretion disk. As for materials farther away from the Black Hole, our results represent a positive test for the Unification Model, suggesting the presence of the torus in (almost) all sources, even if unobscured. ",
        "introduction": "The X-ray spectrum of Seyfert Galaxies consists of several components, whose relative importance changes with energy. It is therefore important to observe them in an energy range as wide as possible. In this respect, simultaneous XMM-\\textit{Newton}/BeppoSAX observations represent a unique opportunity to analyze the relations between different parts of the spectrum. Indeed, the combination of XMM-\\textit{Newton} superior effective area and BeppoSAX broad band coverage makes this kind of observations a powerful tool to disentangle the various components and then put them in the context of the overall spectrum of Seyfert Galaxies.  In particular, simultaneous observations are ideally suited for the study of the origin of the iron lines. On one hand, with XMM-\\textit{Newton} one can study in detail the line profile, which provides valuable hints on the distance of the emitting matter from the Black Hole (BH) by dynamic arguments. If the line is produced in the innermost regions of the accretion disk, the resulting profile has a characteristic double-peaked shape, due to Doppler and gravitational effects \\citep[see][for a review]{fab00}. If, instead, the line is produced much farther away from the BH, as for instance in the `torus' invoked in Unification Models \\citep{antonucci93}, a symmetric profile is preserved, broadening by kinetic motion of the emitting gas being likely below spectral resolution of present detectors. Finally, if the line is produced in the Broad Line Region (BLR), its width should be consistent with that of the broad optical/UV lines. On the other hand, BeppoSAX has proven so far to be the best X-ray mission to measure the amount of the Compton reflection component, which can be crucial to discern between an origin of the line from Compton-thick or Compton-thin matter \\citep{mgm03}.  Simultaneous XMM-\\textit{Newton}/BeppoSAX observations of Seyfert Galaxies have already been analyzed in the past: NGC~5506 \\citep{matt01,bianchi03}, NGC~7213 \\citep{bianchi03b}, NGC~5548 \\citep{pounds03}, IC~4329A \\citep{gond01}, MKN~841 \\citep{petr02} and MCG-6-30-15 \\citep{fab02}. In this paper, we selected a sample of eight sources with simultaneous observations, with the aim to understand the origin of their iron lines. ",
        "conclusions": "Conclusions}  We selected a sample of eight sources observed simultaneously by XMM-\\textit{Newton} and BeppoSAX, taking advantage of the complementary characteristics of the two satellites, which make them one of the most effective ways to investigate the origin of iron lines. The main results of our analysis can be summarized as follows.  \\begin{itemize}  \\item \\textbf{Narrow neutral iron lines are confirmed to be an ubiquitous component in Seyfert spectra.} Their width is generally unresolved, both at the EPIC pn and the \\textit{Chandra} gratings resolutions, thus requiring a production in a material much farther from the BH than the accretion disk. Even when marginally resolved, the measured widths are typically consistent either with an origin from the torus or the BLR. This result is in agreement with the analysis of larger XMM-\\textit{Newton} and \\textit{Chandra} samples \\citep[see e.g.][and references therein]{page04,yaq04}.  \\item \\textbf{None of the analyzed sources shows a broad relativistic iron line.} However, a very broad and gravitationally shifted iron line (as expected from the inner radii of a rapidly rotating BH) is not formally excluded in any of the sources.  \\item \\textbf{All the sources of our sample (with a single exception) show the presence of a Compton reflection component.} It is very appealing to associate the observed Compton reflection with the narrow iron line, since both features are produced by a (mostly) neutral gas and the values of R and the EW are compatible in all sources with a common origin in the same Compton-thick material. Then the most obvious identification for this matter would be  the torus, which has all the necessary properties. Moreover, its presence also in unobscured sources is in complete agreement with Unification Models.  \\item \\textbf{Our sample includes the only Seyfert Galaxy observed by BeppoSAX without a Compton reflection component.} In this source, namely NGC~7213, the detected iron line must originate in a Compton-thin material, such as the BLR or a Compton-thin torus.  \\item \\textbf{Emission lines from ionized iron are detected in three sources: NGC~5506, NGC~7213 and IC~4329A.} These features, at 6.68 and 6.97 keV (only the latter is present in the \\textit{Chandra} spectrum of IC~4329A) and with EWs of a few tens eV, are readily identified with emission from \\ion{Fe}{xxv} and \\ion{Fe}{xxvi}. A possible evidence for a \\ion{Fe}{xxvi} line was found also for MCG-02-58-022. Such lines are often observed with large EWs in Seyfert 2s, produced by an ionized, Compton-thin material. If these emitting regions are also present in the environment of Seyfert 1s, as expected in unification schemes, we should in principle be able to observe these lines as well, although with a much smaller EW due to dilution by the direct nuclear continuum. \\citet{bm02} found that the resulting EWs can be as large as a few tens eV, as observed in our sources. The limited sensitivity and energy resolution of past X-ray missions prevented, until a couple of years ago, the unambiguous detection of these features in unobscured sources. The large sensitivity of XMM-\\textit{Newton} allows now to detect for the first time these lines.  \\item \\textbf{Emission lines around 5-6 keV are detected in two sources: ESO198-G024 and NGC~3516.} These lines, even if sometimes detected with modest significance, appear to be variable on short (ESO198-G024) and long timescales (NGC~3516). A similar feature was also possibly found in the \\textit{Chandra} spectrum of IC~4329A. An interpretation for these lines consists of a 6.4 keV (rest-frame) iron line produced in a well-confined region of the accretion disk. If this is the case, the line photons suffer gravitational redshift and (relativistic) Doppler effects, as for the classical relativistic profile, but the observed line width would be in many cases narrow, because only the blue horn can be visible in moderate signal-to-noise ratio data. Such a possibility would arise if the illumination of the disk is provided in some very anisotropic way, differently from the standard corona model. An example would be a very localized `hot spot' just above the disk, possibly produced by a magnetic flare. Thus, the illumination would interest just a small region of the disk, and iron emission would be produced only there. We defer the reader to \\citet{dov04} for details on this model, which, if correctly explains the features, can provide a precise tool for measuring the BH mass.  \\end{itemize}  The picture emerging from the results presented above strongly requires some corrections for the classical model of reprocessing from the accretion disk. On one hand, the lack of broad, relativistic lines suggests that the physical properties of the disk are likely different from those generally invoked. In particular, it is possible that the disk is truncated at radii larger than the last stable orbit, or has a degree of ionization such as to prevent iron line emission. However, it is also possible that the lines are so broad (as expected, for example, if emitted from the innermost radii of an accretion disk around a rapidly rotating BH) that their profile cannot be easily disentangled from the underlying continuum in data with present statistical quality. On the other hand, the narrow features detected around 5-6 keV, if interpreted in terms of emission from hot orbiting spots above the disk, imply that the disk indeed extends down to the last stable orbit, emitting iron lines. These contradictory results may be telling us that the illuminating properties are very complex, and still not understood.  A powerful test for all the scenarios proposed in this paper would probably be represented by high resolution time-resolved analysis, which adds another valuable piece of information to the data. However, this opportunity has to await the next generation of X-ray satellites, such as \\textit{Constellation-X} and \\textit{XEUS}, to be fully exploited.  \\acknowledgement  We would like to thank the anonymous referee for useful suggestions. This paper is based partly on observations obtained with XMM-\\textit{Newton}, an ESA science mission with instruments and contributions directly funded by ESA Member States and the USA (NASA). SB, GM, IB and GCP acknowledge ASI and MIUR (under grant \\textsc{cofin-03-02-23}) for financial support."
    },
    "0404/astro-ph0404287_arXiv.txt": {
        "abstract": "We report on the first X-ray observations of the neutron star soft X-ray transient (SXT) XTE J2123--058 in quiescence, made by the {\\em Chandra X-ray Observatory} and {\\em BeppoSAX\\/}, as well as contemporaneous optical observations.  In 2002, the {\\em Chandra} spectrum of XTE J2123--058 is consistent with a power-law model, or the combination of a blackbody plus a power-law, but it is not well-described by a pure blackbody.  Using the interstellar value of $N_{\\rm H}$, the power-law fit gives $\\Gamma = 3.1^{+0.7}_{-0.6}$ and indicates a 0.3--8 keV unabsorbed luminosity of $(9^{+4}_{-3})\\times 10^{31}$ ($d$/8.5 kpc)$^{2}$ ergs~s$^{-1}$ (90\\% confidence errors).  Fits with models consisting of thermal plus power-law components indicate that the upper limit on the temperature of a 1.4\\Msun, 10~km radius neutron star with a hydrogen atmosphere is $kT_{\\rm eff} < 66$~eV, and the upper limit on the unabsorbed, bolometric luminosity is $L_{\\infty} < 1.4\\times 10^{32}$ ergs~s$^{-1}$, assuming $d = 8.5$~kpc.  Of the neutron star SXTs that exhibit short ($<$1 year) outbursts, including Aql X-1, 4U 1608--522, Cen X-4, and SAX J1810.8--2609, the lowest temperatures and luminosities are found for XTE J2123--058 and SAX J1810.8--2609.  From the {\\em BeppoSAX\\/} observation of XTE J2123--058 in 2000, we obtained an upper limit on the 1--10 keV unabsorbed luminosity of $9\\times 10^{32}$ ergs~s$^{-1}$.  Although this upper limit allows that the X-ray luminosity may have decreased between 2000 and 2002, that possibility is not supported by our contemporaneous $R$-band observations, which indicate that the optical flux increased significantly.  Motivated by the theory of deep crustal heating by Brown and co-workers, we characterize the outburst histories of the five SXTs.  The low quiescent luminosity for XTE J2123--058 is consistent with the theory of deep crustal heating without requiring enhanced neutron star cooling if the outburst recurrence time is $\\gsim$70 years. ",
        "introduction": "Accreting neutron stars can be found in high-mass (HMXB) or low-mass (LMXB) X-ray binary systems.  The majority of HMXBs have transient X-ray emission. Their outburst spectra are relatively hard and X-ray pulsations from these highly magnetized ($B\\sim 10^{12}$~G) neutron stars are typically detected. A wide variety of X-ray behaviors are seen for neutron star LMXBs, but, in general, the lack of X-ray pulsations from most (but not all) of these systems, and the emission of type I X-ray bursts from some, suggest that they harbor neutron stars with relatively low magnetic field strengths ($B\\sim 10^{8-9}$~G). During $\\approx 33$ years of X-ray observations, some sources (e.g., Sco X-1, Cyg X-2) have maintained luminosities approaching the Eddington limit of $\\approx 10^{38}$ ergs~s$^{-1}$, while others are able to maintain persistent luminosities several orders of magnitude lower \\citep{wilson03}.  In addition, there is a class of transient neutron star LMXBs for which the luminosity varies from a substantial fraction of Eddington to quiescent levels typically near $10^{32-33}$ ergs~s$^{-1}$.  In outburst, these systems have relatively soft spectra compared to the HMXBs, and are commonly grouped with black hole transients as soft X-ray transients (SXTs).  In quiescence, most neutron star SXTs exhibit X-ray energy spectra with a component that is typically fitted well by a blackbody, suggesting that the origin of this component is thermal emission from the surface of a cooling neutron star.  Although a pure blackbody often provides a good fit to the spectrum, unphysical neutron star radii near 1 km are inferred unless an atmosphere is modeled \\citep{rutledge99}.  In addition to the thermal component, the energy spectra often contain a second component that has a power-law shape.  The brightest and best studied systems in this class, Cen X-4 and Aql X-1, usually display both components \\citep{rutledge01,rutledge02a,cs03,campana04}.  However, other systems may be dominated by the thermal component, such as MXB 1659--29 \\citep{wijnands03b} and sources X-5 and X-7 in the globular cluster 47 Tucanae \\citep{heinke03}, or by the power-law component, such as SAX J1808.4--3658 \\citep{campana02} and EXO 1745--248 \\citep{wijnands03a}.  Although theories for the thermal component, such as the deep crustal heating model of \\cite{bbr98}, are relatively well-developed and are being tested with observations, the origin of the power-law component is not understood beyond suggestions that it may be related to accretion onto the neutron star magnetosphere \\citep{campana98} or a putative pulsar wind colliding with infalling matter from the companion star \\citep{tavani91}.  In addition to our lack of understanding of the power-law component, questions remain about the mass accretion rate in quiescence, the origin of rapid (100--10,000~s) variability \\citep{rutledge02a,campana04}, and the origin of variability in the thermal component on longer time scales \\citep{rutledge02a}.  Another important question is if quiescent observational properties correlate with other known differences between neutron star SXTs, such as whether they are millisecond X-ray pulsars (during outbursts) or not, whether the systems are in the field or in globular clusters, and whether their X-ray outbursts are long (years to decades) or short (weeks to months).  Here, we report on X-ray and optical observations of the field neutron star SXT XTE J2123--058 taken during quiescence.  XTE J2123--058 had its only detected X-ray outburst in 1998 June-August \\citep{lss98,tomsick99}, and we focus on observations made with {\\em Chandra}, {\\em BeppoSAX\\/}, and optical telescopes 2--4 years after the outburst.  During the outburst, the {\\em Rossi X-ray Timing Explorer (RXTE\\/)} detected type I X-ray bursts and a pair of kHz quasi-periodic oscillations \\citep{homan99,tomsick99}, indicating that the system contains a rapidly rotating neutron star. However, coherent X-ray pulsations were not found.  The 6~hr binary orbital period and the fact that the binary inclination of the system is relatively high were established from optical modulation and the presence of partial eclipses in the optical light curve \\citep{tomsick99,swg99,zurita00,shahbaz03}.  The high Galactic latitude ($b = -36^{\\circ}$) and low extinction have allowed for detailed optical studies of XTE J2123--058 in quiescence even though the source is rather faint at its relatively large distance ($8.5\\pm 2.5$ kpc).  The optical observations show that XTE J2123--058 consists of a K7~V star on or close to the main sequence and a neutron star for which mass determinations of $1.5\\pm 0.3$~\\Msun~\\citep{tomsick01,casares02,tomsick02} and 1.04--1.56\\Msun~\\citep{shahbaz03} have been obtained.  The focus of this paper is the first X-ray study of XTE J2123--058 in quiescence. ",
        "conclusions": ""
    },
    "0404/astro-ph0404078_arXiv.txt": {
        "abstract": "{ We present XMM-{\\it Newton} observations of the young ($\\sim 2-5$ Myr) cluster around the hot (O9.5V) star $\\sigma$~Orionis~AB, aimed at obtaining a high resolution RGS spectrum of the hot star as well as EPIC imaging data for the whole field. We show that the RGS spectrum of $\\sigma$~Ori~AB may be contaminated by weaker nearby sources which required the development of a suitable procedure to extract a clean RGS spectrum and to determine the thermal structure and wind properties of the hot star. We also report on the detection of a flare from the B2Vp star $\\sigma$~Ori~E and we discuss whether the flare originated from the hot star itself or rather from an unseen late-type companion. Other results of this observation include: the detection of 174 X-ray sources in the field of $\\sigma$~Ori of which 75 identified as cluster members, including very low-mass stars down to the substellar limit; the discovery of rotational modulation in a late-type star near $\\sigma$~Ori~AB; no detectable line broadenings and shifts ($\\la 800$~km~s$^{-1}$) in the spectrum of $\\sigma$~Ori~AB together with a remarkable low value of the \\ion{O}{vii} forbidden to intercombination line ratio and unusually high coronal abundances of CNO elements. ",
        "introduction": "The $\\sigma$~Ori cluster, discovered by {\\it ROSAT} \\citep{wolk96, walter97} around the O9.5V star $\\sigma$~Ori~AB, belongs to the OB1b association and is located at a distance of $352_{-85}^{+166}$ pc \\citep[from {\\it Hipparcos},][]{hippa}. In addition to several hot stars, it is known to contain $\\sim 100$ likely pre-main sequence late-type stars within $30^\\prime$ of $\\sigma$~Ori, as well as some brown dwarfs and planetary-mass objects \\citep{bejar99,zapat00,bejar01}. The estimated age of the cluster is $2-5$ Myr.  We have obtained an XMM-{\\it Newton} observation of the $\\sigma$~Ori cluster, centered on the hot star $\\sigma$~Ori~AB, with the purpose of obtaining: i) a high-resolution RGS spectrum of the central source; ii) imaging data as well as low-resolution spectra over the whole field, including both a few early-type stars and a large number of late-type stars down to the substellar limit. Given the high sensitivity of XMM-{\\it Newton}, and the good combination of low- and high-resolution spectroscopic instruments on board, these observations were expected to shed light on the coronal and/or wind properties of stars in a very young cluster.  X-ray emission from O and B stars is usually explained with the presence of winds. X-ray observations of the hot stars $\\zeta$~Pup (O4If) and $\\zeta$~Ori (O9Ib) have shown the presence of such winds, with velocity widths of the order of $600-1500$~km~s$^{-1}$ and blueshifted centroids \\citep{wald01,cas01,kahn}. However, some conflicting results have also been found, like high-densities close to the stellar surface in $\\zeta$~Ori \\citep{wald01}, where the velocity is too small to produce the shocks required for the X-ray emission, or a temperature structure in the Orion Trapezium hot stars that is similar to that of cool active stars, where the emission originates from magnetically confined coronal structures \\citep{sch03}. This has raised the question of whether coronal loops might be present in some hot stars partially contributing to their X-ray emission. High-resolution spectroscopic observations of the hot star $\\sigma$~Ori with XMM-{\\it Newton} can clarify some of these issues.  As it will be described below, there are several X-ray sources, both hot and cool stars, in our XMM-{\\it Newton} field close enough to the central source to potentially contaminate its high-resolution RGS spectrum. Although their X-ray intensity lies well below the level of $\\sigma$~Ori~AB, they produce lines that can contribute significantly at certain wavelengths. These lines are shifted in wavelength with respect to those emitted by $\\sigma$~Ori~AB, because of the different locations of these sources within the RGS field of view (FOV). These spurious lines must be identified in order to correctly analyze the spectrum of the central source. To this aim, the capability of XMM-{\\it Newton} of obtaining simultaneous high-resolution spectra of the central source with RGS and low-resolution spectra of the other sources in the field with EPIC is of great advantage. By using the information derived from the EPIC spectra, it is possible in fact to model the expected contributions of nearby sources to the RGS spectra. This, together with the wavelengths shifts in the RGS spectra caused by the different offsets of the sources in the RGS FOV, allows an accurate correction of the $\\sigma$~Ori~AB spectrum for the contribution of nearby sources.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[clip]{sori_epic_bw.ps}} \\caption{Composite EPIC (MOS1+MOS2+PN) image of the $\\sigma$~Ori cluster} \\label{fig:sori_ima} \\end{figure}  This paper is organized as follows. In Sect.~\\ref{sec:observ} we present the analysis of the XMM-{\\it Newton} observation, discussing first the imaging data obtained with the EPIC PN and MOS detectors, and then the high-resolution spectroscopic data obtained with the RGS. Since the observed RGS spectra of the central source might in principle be contaminated by up to three nearby sources in addition to $\\sigma$~Ori~AB itself, we present in this section light curves and low-resolution EPIC spectra of these four sources, showing that one of them (the hot star $\\sigma$~Ori~E) is undergoing a flaring episode, another one (a K star) is showing evidence of rotational modulation, while $\\sigma$~Ori~AB and the fourth source (another K star) are either quiescent or of low variability. With regard to the EPIC spectra also presented in this section, $\\sigma$~Ori~AB appears much softer than the other three sources, consistently with typical X-ray spectra of hot stars. In Sect.~\\ref{sec:results}, we present the results of our analysis, first for $\\sigma$~Ori~AB and then for the flare on $\\sigma$~Ori~E. In particular, we derive, for the former star, the ``cleaned'' differential emission measure distribution, and we put constraints on wind velocities and shifts, chemical abundances and densities. For $\\sigma$~Ori~E we discuss the flare properties and we present evidence for circumstellar absorption. In Sect.~\\ref{sec:disc}, we discuss the implication of our results, both for current models of shocked winds in early-type stars (as in $\\sigma$~Ori~AB), and for the possible occurence of flares in hot stars (as opposite to flares originating from unseen late-type companions). Finally in Sect.~\\ref{sec:conclusions} we summarize our conclusions. ",
        "conclusions": ""
    },
    "0404/astro-ph0404552_arXiv.txt": {
        "abstract": "An experimental method to determine the (n,$\\gamma$) cross section of short-living $s$-process branching points using data of the inverse ($\\gamma$,n) reaction is presented. The method was used to observe the branching point nucleus $^{95}$Zr because the elemental abundance patterns corresponding to this branching point cannot be reproduced by full stellar models and a possible error source is the neutron capture cross section of $^{95}$Zr. The analysis of the experiment is still under progress, we will outline the current status in this manuscript. ",
        "introduction": "\\label{sec:intro} The nucleosynthesis of the elements heavier than iron is today mainly explained by three processes: $s$-, $r$-, and $p$-process. $s$- and $r$-process are based on neutron-capture reactions with adjacent $\\beta$-decays while the $p$-process is governed by photodisintegration reactions such as ($\\gamma$,n), ($\\gamma$,p), or ($\\gamma$,$\\alpha$).  The distinction between $s$- and $r$-process is due to the different neutron densities $n_{\\rm n}$ involved. During $s$-process nucleosynthesis $n_{\\rm n}$ is in the order of $10^8\\ {\\rm neutrons}/{\\rm cm}^3$ while typical $r$-process sites deal with $n_{\\rm n} > 10^{20}\\ {\\rm neutrons}/{\\rm cm}^3$ \\cite{beer00}. Hence, the neutron capture rates $\\lambda_{({\\rm n},\\gamma)}$ are larger than typical $\\beta$-decay rates $\\lambda_\\beta$ during $r$-process nucleosynthesis ($\\lambda_{({\\rm n},\\gamma)} \\gg \\lambda_\\beta$, $r$: rapid). In the $s$-process the situation is the other way round ($\\lambda_{({\\rm n},\\gamma)} \\ll \\lambda_\\beta$, $s$: slow) and the involved nuclei are close to the valley of $\\beta$-stability.%  However even in a $s$-prozess scenario, if the half-life $T_{1/2}$ is long enough and the neutron capture cross section high enough another neutron capture might take place and the path ``branches'' out. Therefore, these nuclides are called branching points of the $s$-process. The branching ratio -- i.e. how often each of the paths is taken -- determines the corresponding elemental abundance patterns. Using a model for $s$-process nucleosynthesis and knowing precisely the nuclear physics input (half-life $T_{1/2}$ and Maxwellian averaged capture cross section (MACS) at typical $s$-process temperatures $kT = 30\\ {\\rm keV}$) it is possible to determine the astrophysical parameters temperature $T$ and neutron density $n_{\\rm n}$.  In the so-called classical approach temperature $T$ and neutron density $n_{\\rm n}$ are considered to be constant. Thus, three different components produced with different astrophysical parameters are needed to reproduce the observed abundances of $s$-only nuclei: the weak component corresponds to mass numbers $A < 90$ while the strong component only describes the termination of the $s$-process path at lead and bismuth. In between these borders the isotopes belong to the so-called main $s$-process \\cite{kaep99}. The abundances of $s$-only isotopes that are not affected by a branching can be reproduced with a mean square deviation of about 3\\% using this simple model \\cite{kaep90}.  A more realistic approach is a full stellar model, e.g. the AGB star model described in Ref.~\\cite{luga03a}. These models need very precise nuclear physics input data to reproduce the $s$-only abundances. Especially the MACS of several branching points hamper the reliability of the predictions: due to the lack of experimental data the theoretical predictions of the MACS sometimes show a broad spread (see \\cite{bao00}).  A direct measurement of the MACS of the branching points is only possible if their half-lives are in the order of years (e.g. $^{147}$Pm with $T_{1/2} = 2.62\\ {\\rm yr}$ \\cite{reif03}). In the case of short-living branching points with half-lives of about a dozen days or even less, adequate samples are not available. Thus, direct measurements are possible.  To solve this problem we have investigated an experimental method using the data of the inverse ($\\gamma$,n) reaction as described in Section~\\ref{sec:expmeth}. Section~\\ref{sec:prop} presents % $^{95}$Zr focusing on astrophysical aspects as well as on the constraints of our experimental method for this nuclide. The present status of the analysis is summarized in Section~\\ref{subsec:status}. In Section~\\ref{sec:summ} we describe the next step in the data evaluation. ",
        "conclusions": "\\label{sec:summ} The precise knowledge of the (n,$\\gamma$) cross section of the $s$-process branching point $^{95}$Zr is a crucial input parameter to test the reliability of the AGB star model. Due to the very wide range of the theoretically predicted values a measurement observing the inverse reaction with the photoactivation technique was carried out at the S--DALINAC in June 2002.  During the analysis of the data several problems occured: The energy dependence of the ($\\gamma$,n) cross section of $^{96}$Zr extracted from the data differs significantly from the theoretical prediction. In our opinion, the deviation is due to resonances that are not included in the used theory. Therefore, the experimental method described in Section~\\ref{sec:expmeth} determines an energy dependent correction factor $f(E_\\gamma)$ that cannot be used to correct the predicted MACS of $^{95}$Zr without further studies.  Hence, the experiment was extended in May 2003 by measuring the ($\\gamma$,n) cross section of $^{96}$Zr closer to the reaction threshold. Therefore, a new calibration standard with lower neutron separation energy was measured during the same beam-time. The analysis of the new standard $^{187}$Re is still under progress, preliminary results like the spectrum shown in Fig.~\\ref{fig:cali} are encouraging that it will work as well as our first standard $^{197}$Au.  Once a description of the ($\\gamma$,n) cross section of $^{187}$Re near threshold is available the new data point of $^{96}$Zr taken at $E_{\\rm max} = 8100\\ {\\rm keV}$ has to be added to the existing analysis. Therewith, the discrepancy between the experimental result and the theoretical prediction will hopefully disappear or at least become explainable. The outcome of the whole analysis that is the experimentally confirmed MACS of $^{95}$Zr can then be used in the AGB star model and might solve the problems in explaining the observed zirconium abundance patterns in SiC grains.  In future, we will take more data points at different energies $E_{\\rm max}$ in order to reduce difficulties due to resonances on top of the cross section. If the distance between two data points $\\Delta E_{\\rm max}$ is smaller, the determination of the location of a resonance is possible with higher precision.   \\begin{theacknowledgments} The authors would like to thank R.\\ Gallino (Torino, Italy) and F.\\ K\\\"appeler (Karlsruhe, Germany) for fruitful discussions concerning $s$-process modelling. This work was supported by the Deutsche Forschungsgemeinschaft (contracts Zi 510/2-2 and SFB 634) and Swiss NSF (grants 2124-055832.98, 2000-061822.00, 2024-067428.01). \\end{theacknowledgments}"
    },
    "0404/astro-ph0404546_arXiv.txt": {
        "abstract": "The basic workings of inflationary models are summarized, along with the arguments that strongly suggest that our universe is the product of inflation.  I describe the quantum origin of density perturbations, giving a heuristic derivation of the scale invariance of the spectrum and the leading corrections to scale invariance.  The mechanisms that lead to eternal inflation in both new and chaotic models are described.  Although the infinity of pocket universes produced by eternal inflation are unobservable, it is argued that eternal inflation has real consequences in terms of the way that predictions are extracted from theoretical models.  Although inflation is generically eternal into the future, it is not eternal into the past: it can be proven under reasonable assumptions that the inflating region must be incomplete in past directions, so some physics other than inflation is needed to describe the past boundary of the inflating region. The ambiguities in defining probabilities in eternally inflating spacetimes are reviewed, with emphasis on the youngness paradox that results from a synchronous gauge regularization technique. ",
        "introduction": "I will begin by summarizing the basics of inflation, including a discussion of how inflation works, and why many of us believe that our universe almost certainly evolved through some form of inflation.  This material is mostly not new, although the observational evidence in support of inflation has recently become much stronger.  Since observations of the cosmic microwave background (CMB) power spectrum have become so important, I will elaborate a bit on how it is determined by inflationary models. Then I will move on to discuss eternal inflation, showing how once inflation starts, it generically continues forever, creating an infinite number of ``pocket'' universes.  If inflation is eternal into the future, it is natural to ask if it can also be eternal into the past.  I will describe a theorem by Borde, Vilenkin, and me (Borde, Guth, \\& Vilenkin 2003) which shows under mild assumptions that inflation cannot be eternal into the past, and thus some new physics will be necessary to explain the ultimate origin of the universe. ",
        "conclusions": "In this paper I have summarized the workings of inflation, and the arguments that strongly suggest that our universe is the product of inflation.  I argued that inflation can explain the size, the Hubble expansion, the homogeneity, the isotropy, and the flatness of our universe, as well as the absence of magnetic monopoles, and even the characteristics of the nonuniformities. The detailed observations of the cosmic background radiation anisotropies continue to fall in line with inflationary expectations, and the evidence for an accelerating universe fits beautifully with the inflationary preference for a flat universe. Our current picture of the universe seems strange, with 95\\% of the energy in forms of matter that we do not understand, but nonetheless the picture fits together very well.  Next I turned to the question of eternal inflation, claiming that essentially all inflationary models are eternal. In my opinion this makes inflation very robust: if it starts anywhere, at any time in all of eternity, it produces an infinite number of pocket universes.  Eternal inflation has the very attractive feature, from my point of view, that it offers the possibility of allowing unique (or possibly only constrained) predictions even if the underlying string theory does not have a unique vacuum.  I discussed the past of eternally inflating models, concluding that under mild assumptions the inflating region must have a past boundary, and that new physics (other than inflation) is needed to describe what happens at this boundary.  I have also described, however, that our picture of eternal inflation is not complete.  In particular, we still do not understand how to define probabilities in an eternally inflating spacetime.  The bottom line, however, is that observations in the past few years have vastly improved our knowledge of the early universe, and that these new observations have been generally consistent with the simplest inflationary models."
    },
    "0404/astro-ph0404293_arXiv.txt": {
        "abstract": "Explosions of type Ic supernovae (SNe Ic) are investigated using a relativistic hydrodynamic code to study roles of their outermost layers of the ejecta in light element nucleosynthesis through spallation reactions as a possible mechanism of the \"primary\" process. We have confirmed that the energy distribution of the outermost layers with a mass fraction of only $0.001$ \\% follows the empirical formula proposed by previous work when the explosion is furious.  In such explosions, a significant fraction of the ejecta ( $>$0.1 \\% in mass ) have the energy greater than the threshold energy for spallation reactions. On the other hand, it is found that the outermost layers of ejecta become more energetic than the empirical formula would predict when the explosion energy per unit ejecta mass is smaller than $\\sim 1.3\\times 10^{51}\\mbox{ ergs/}\\Msun$. As a consequence, it is necessary to numerically calculate explosions to estimate light element yields from SNe Ic. The usage of the empirical formula would overestimate the yields by a factor of $\\gtsim 3$ for energetic explosions such as SN 1998bw and underestimate the yields by a similar factor for less energetic explosions like SN 1994I. The yields of light elements Li, Be, and B (LiBeB) from SNe Ic are estimated by solving the transfer equation of cosmic rays originated from ejecta of SNe Ic and compared with observations. The abundance ratios Be/O and B/O produced by each of our SNe Ic models are consistent with those of metal-poor stars. The total amounts of these elements estimated from observations indicate that energetic SNe Ic like SN 1998bw could be candidates for a production site of Be and B in the Galactic halo only when the fraction of this type out of all the SNe was a factor of $>100$ higher than the value estimated from current observational data.  This primary mechanism would predict that there are stars significantly deficient in light elements which were formed from the ISM not affected by SNe Ic. Since this has no support from current observations, other primary mechanisms such as the light element formation in superbubbles are needed for other types of SNe. The observed abundance pattern of all elements including heavy elements in metal-poor stars suggests that these two mechanisms should have supplied  similar amounts of Be and B. Our calculations show that SNe Ic can not produce an appreciable amount of Li. ",
        "introduction": "The amounts of the light elements Li, Be, and B (LiBeB) at present are thought to be the sum of the products of the big bang nucleosynthesis and the subsequent cosmic-ray spallation reactions. The contribution from the big bang nucleosynthesis to $^7$Li is thought to be most pronounced among these three elements. Though some fractions of Li have been depleted inside some cool stars, its abundances on the surfaces of metal-poor stars are the only information that we can use to  identify the contribution from the big bang nucleosynthesis \\citep[e.g.,][]{Ryan_99}. Thus it is important to know the contribution from the cosmic-ray spallation reactions to Li in the early stages of the Galaxy evolution during which metal-poor stars were formed. With this knowledge, we can constrain the cosmological parameters to synthesize light elements in the big bang nucleosynthesis \\citep{Ryan_00}.  In addition, the investigation of the evolution of these elements in the early  stages of the Galaxy enables us to understand the origin of  cosmic-rays responsible for the light element nucleosynthesis. Recent observations of extremely metal-poor stars  apparently suggest that the primary (not secondary) spallation process dominated the light element nucleosynthesis in the early Galaxy \\citep{Duncan_92}. In other words, cosmic-rays composed of C/O interacting with protons and/or He nuclei in the interstellar medium (ISM) have predominantly produced LiBeB.  A primary mechanism, in which LiBeB were produced in supernova ejecta-enriched superbubbles, was suggested to explain the Be evolution in the early Galaxy \\citep{Higdon_98}. Since then, two-component models in which light elements are produced by standard Galactic cosmic-rays and metal-enriched particles in superbubbles have been investigated by several authors \\citep{Ramaty_00, Fields_00, Suzuki_01}. These studies concluded that a primary mechanism is needed to explain the observed abundance trends of Be and B with O and Fe. Recently,  \\citet{Suzuki_01} investigated the chemical evolution of LiBeB in the early epoch of the Galaxy using an inhomogeneous chemical evolution model developed by \\citet{Tsujimoto_99} in which star formation is assumed to be induced by supernova (SN) explosions. Their model succeeded in reproducing not only the observed metallicity distribution of Galactic halo stars but also the observed abundance correlations of heavy elements. \\citet{Suzuki_01} also considered two origins of Galactic cosmic-rays that synthesize light elements. One is from the ISM accelerated by SN shocks, and the other from SN ejecta accelerated by the SN shock. The authors claimed that $\\sim$2 \\% of Galactic cosmic-rays must be originated from the SN ejecta to reproduce the behavior of the abundance of Be at the metal-poor ends. However, their model suffers from shortage of energy supply from each SN to cosmic-rays. Moreover, the energy distributions of the cosmic-rays from these two origins might be different, though \\citet{Suzuki_01} assumed they are the same.  \\citet{Fields_96} found from a numerical model of type Ic supernovae (SNe Ic) \\citep{Nomoto_90} that the outermost C/O layers of  an SN could attain energies beyond the threshold value to produce LiBeB ($\\sim$30 MeV per nucleon for O$+$H$\\rightarrow$Be, which corresponds to the Lorentz factor of $\\sim$1.03). Energetic explosions of massive stars with  stripped H-rich and He layers might be able to produce cosmic-rays enriched with C/O in the outermost layers of ejecta.  In the model of \\citet{Nomoto_90}, the outermost layers of the ejecta do not become so relativistic. It may, however, be attributed to their coarse zoning in the outermost layers, which has to be as accurate as possible for this purpose.  Furthermore, their numerical calculations were performed with a hydrodynamic code that does not take into account relativistic effects. Thus it was impossible to accurately estimate the contribution from SNe Ic to the light element nucleosynthesis with their results.  The phenomena taking place when the supernova blast wave hits the surface of a relatively compact star were investigated in a more sophisticated fashion by \\citet{Ensman_92} and \\citet{Blinnikov_00} with their radiation-hydrodynamic codes. Their codes allow the radiation and the gas to go out of equilibrium. Both of their codes were non-relativistic except that they included light-travel-time corrections. They were concerned with the shock breakout of SN 1987A, because the most detailed observations immediately after the shock breakout were available for this SN. One of their main objectives was a detailed modelling of the UV burst immediately after the shock emergence. In addition to these numerical approaches, there have been semi-analytic approaches to the shock emergence in the plane-parallel medium in which the flow is assumed to be self-similar \\citep{Gandel'man_56, Sakurai_60}. \\citet{Kazhdan_92}  took into account the sphericity of the flow in a neighborhood of the surface. These semi-analytic approaches are also limited to the non-relativistic flow.  After the emergence of a very bright type Ic supernova, SN 1998bw, was found to be associated with a $\\gamma$-ray burst GRB 980425 \\citep{Galama_98} and \\citet{Kulkarni_98} inferred from their radio observations that the radio shock front of SN 1998bw was moving at relativistic speeds with the Lorentz factor between 1.6 and 2, the relativistic motion from supernova explosions has been investigated. \\citet{Matzner_99} estimated how much mass of relativistic ejecta could be obtained from the explosion of a massive star and derive an empirical formula giving the mass of relativistic ejecta from explosions of stars with simple polytropic density structures. Later, \\citet{Tan_01} revised the empirical formula using their relativistic hydrodynamic code. Their formula would give $\\sim 10^{-6}\\,M_\\odot$ of relativistic ejecta (Lorentz factor $>2$) from the explosion of a star with a mass of $10\\,M_\\odot$ and an energy of $10^{52}$ ergs.  Using the empirical formula of \\citet{Matzner_99} for the energy distribution of SN ejecta,  \\citet{Fields_02} concluded that SNe Ic, especially energetic events like SN 1998bw significantly contribute to light element nucleosynthesis through spallation reactions. The energy distribution of particles in SN ejecta used in \\citet{Fields_02} follow a power law with a power index of $-4.6$. This power index is quite different from that of the energy distribution of the ISM accelerated by SN remnant shock \\citep[e.g.][]{Meneguzzi_71}. Then \\citet{Fields_02}  adopted a ``thick target''  approximation instead of solving the cosmic-ray transfer equation to derive the yields of LiBeB from the energy distribution of SN ejecta.  There still remain some problems to be addressed in their work;  the model for stellar structures \\citet{Matzner_99} assumed is so simple that it involves suspicions of inaccurate estimates. In addition, they set the adiabatic index $\\gamma$=constant over the whole stages of explosions, which may vary according to the physical conditions and affect the resultant energy distribution of the ejecta.  Our objective is to construct a realistic model for SN ejecta moving at relativistic speeds and to investigate their contribution to light element syntheses. To improve the above situations, we calculate the energy distributions of ejecta as a result of SN explosions using a relativistic hydrodynamic (RHD) code with realistic numerical models for massive stars as the initial conditions. The atmospheres in radiative equilibrium are attached to the original models for massive stars (see references in Table \\ref{tbl-model}). This is essential to investigate the energy distribution of ejecta at highest energies after the shock breakout. We also verify the validity of $\\gamma$=constant by using more realistic equation of states that incorporates radiation and ideal gas in thermodynamic equilibrium. To investigate the change of the energy distribution of relativistic ejecta transferring in the ISM, we solve the transfer equation that takes into account the energy loss due to  ionization and spallation reactions. From these calculations, the amounts of synthesized LiBeB through primary spallation reactions are obtained.  In the next two sections, we describe our RHD code in  Lagrangian coordinates (\\S2) and initial conditions (\\S3). Then, we show the results in \\S4,  compare them with that of other authors, and discuss the differences. In \\S5, we estimate the yields of LiBeB using our explosion models together with the leaky box model. ",
        "conclusions": "We have performed numerical calculations for SNe Ic explosions using  a relativistic hydrodynamic code to investigate how much mass of ejecta is accelerated beyond the threshold energy for spallation reactions to synthesize light elements. In our calculations, realistic massive star models are used as the initial conditions of SNe and the EOS takes into account the thermal radiation and ideal gas. We have compared the resultant energy distributions of ejecta with the empirical formula derived in \\citet{Fields_02} for some SN explosions including furious and normal  SNe and found that the energy distributions of ejecta from the numerical calculations and the empirical formula agree only in the high energy tail when the explosion energy per unit ejecta mass significantly exceeds $1.3\\times 10^{51}\\mbox{ ergs/}\\Msun$. Otherwise, the empirical formula overestimates or underestimates the ejecta mass at around the threshold energy for spallation reactions. Therefore it is necessary to numerically calculate SN explosions to obtain a correct energy distribution of ejecta for estimations of the yield of light elements.  To obtain the yields of light elements from the calculated SN ejecta, we have numerically solved the transfer equation taking into account the energy loss due to ionization of the ISM and spallation reactions with the ISM. The results suggest that light elements synthesized from energetic SNe Ic like SN 1998bw by this mechanism can explain the enrichment of Be and B observed in the Galactic halo stars if the fraction of SN 1998bw like SNe in the early Galaxy was a factor of $>100$ higher than current observational data  suggest. This mechanism must not be the only primary mechanism that worked in the Galactic halo. Other primary mechanisms like superbubbles that supply light elements regardless of SN type are required to reproduce the observed abundance ratios such as Be/Fe. These SNe Ic can produce Li with more than one order of magnitude smaller amounts than indicated by observations. However, lack of information on the Li abundances in the metallicity range of $-2\\ltsim$[Fe/H]$\\ltsim -1$ prevents us from deducing a firm conclusion.  SNe are suggested to be associated with aspherical explosions. The deviation from spherical symmetry will be able to increase the mass of ejecta with enough energies for spallation reactions for a given $E_{\\rm ex}/M_{\\rm ej}$ because $M(>\\epsilon)$ is proportional to the $\\sim3.4$---3.6 power of this value. To illustrate this effect, a simplified situation will be considered. Suppose the energy injected in the direction with a solid angle of $\\pi$ steradian is enhanced by a factor of two and the energy in the other directions is reduced by a factor of $1.5$, then the empirical formula for the mass $M(>\\epsilon)$ indicates that this mass will increase by a factor of $\\sim 2^{3.4}\\times 1/4 + (2/3)^{3.4}\\times3/4\\sim 2.8$ while the total energy will be unchanged. Applying the empirical formula obtained from the spherically symmetric calculations to this situation might lead to an erroneous result. Thus to further explore SNe Ic as a production site for light elements, we need to perform multi-dimensional relativistic hydrodynamic calculations for SNe Ic that can trace the motion of the outermost ejecta with a sufficient accuracy such as the calculations presented here."
    },
    "0404/astro-ph0404400_arXiv.txt": {
        "abstract": "We compute the covariance expected between the spherical harmonic coefficients $a_{\\ell m}$ of the cosmic microwave temperature anisotropy if the universe had a compact topology. For fundamental cell size smaller than the distance to the decoupling surface, off-diagonal components carry more information than the diagonal components (the power spectrum). We use a maximum likelihood analysis to compare the Wilkinson Microwave Anisotropy Probe first-year data to models with a cubic topology. The data are compatible with finite flat topologies with fundamental domain $L > 1.2$ times the distance to the decoupling surface at 95\\% confidence. The WMAP data show reduced power at the quadrupole and octopole, but do not show the correlations expected for a compact topology and are indistinguishable from infinite models. ",
        "introduction": "The simplest model for the universe is a spatially homogeneous, isotropic spacetime with a Euclidian (flat) geometry. This simple model is consistent with observations, but leaves unaddressed the question of topology or the connectedness of spacetime. Schwartzschild \\citep{schwartzschild:1900} first noted the possiblitity of non-trivial topology for the Universe even before Einstein's discovery of his field equations. Almost immediately after Einstein's discovery, de Sitter \\citep{desitter:1917} pointed out the the field equations did not constrain the topology. Since general relativity provides no theoretical guidance, we turn to observations to constrain topology.  Observational tests of topology all rely on multiple imaging of distant objects. If the universe is multiply-connected with cell size smaller than the distance to some object, photons from that object can reach the observer via multiple paths. Simply searching the sky for multiply-imaged point sources, e.g. quasars, is problematic: since the travel time to each image is different, each image shows the same object at a different time. If source evolution is important, the multiple images may no longer be recognizable as such. The ideal source for topological tests would fill the whole sky with a pattern centered on the observer and emitted at a single time. The cosmic microwave background (CMB) is an excellent approximation to this ideal source.  A number of authors have used the CMB to constrain the topology of the universe. These tests fall into two general categories. A compact topology can not support spatial structure with wavelength longer than the cell size. The CMB temperature anisotropy will thus be suppressed on angular scales larger than the (projected) cell size. The first category tests use the CMB power spectrum (or its Legendre transform, the 2-point correlation function) to test for non-trivial topology.  The CMB in fact shows significantly less power in the quadrupole and octopole than would be expected for a model based on higher-order moments. The discrepancy was first detected by the Cosmic Background Explorer \\citep{bennett/etal:1996} and verified at much high signal to noise ratio by the Wilkinson Microwave Anisotropy Probe (WMAP) \\citep{bennett/etal:2003}. Figure \\ref{fig1-powerspectrum} shows the angular power spectrum of the WMAP first-year data compared to the best-fit $\\Lambda$CDM model \\citep{spergel/etal:2003}. Models with compact topology $L \\sim 1$ provide a good match to the observed power spectrum, motivating tests of finite-universe models \\citep{oliveira/smoot:1995, luminet/etal:2003}.  The suppression of power on large angular scales is a necessary but not sufficient condition for the existence of a compact topology. The power spectrum is rotationally invariant, averaging over any phase information in the pattern of CMB anisotropy. Such phase information must exist for compact topologies, and forms the basis for a second class of tests. A ``circles on the sky'' search \\citep{cornish/etal:1998} provides a more stringent test for compact topologies. The CMB decoupling surface is a sphere centered on the observer. If the cell size is smaller than the distance to the decoupling surface, the multiple images of this sphere induced by a compact topology will intersect to produce patterns that match along circles. Such circles are not observed, limiting the cell size $L > 1.7$ for a wide class of models \\citep{cornish/etal:2003}.  Additional tests are possible. Compact topologies will not produce circles on the sky if the cell size is larger than the distance to the source, since the resulting images will not intersect. Compact topologies with $L > 2$ may still be distinguished using phase information. In this paper, we describe the correlations imposed on the microwave background by the topology. We use this formalism to compare the WMAP first-year data to a model with cubic topology and derive constraints on the cell size $L$. ",
        "conclusions": "A compact topology imposes a specific pattern of correlations $\\left<{a_{\\ell m}}\\, a_{\\ell' m'}\\right>$ between the spherical harmonic expansion of the CMB temperature. We compute the expected correlations for the simplest non-trivial topology, the cubic torus, and compare a range of cell size $L$ to the WMAP first-year data using a maximum-likelihood algorithm. The covariance matrix explicitly includes the contribution from the integrated Sachs-Wolfe effect on large angular scales and the acoustic peaks at small scales. We separate the covariance into a piece dependent on the topology and a piece dependent on the cosmology. Although the transfer functions $\\Delta^{(S)}_{T\\ell}(k)$ for the cosmology assume isotropy in $k$-space, which is no longer exact for compact models, the effect is predominantly in the cosmology with negligible effect on the topology.  The algorithm is sensitive both to the power spectrum of the data (diagonal elements of the covariance matrix for different ${a_{\\ell m}}$) as well as the phase information contained in the off-diagonal elements. For cell size $L < 2$ the off-diagonal elements are larger than the diagonal elements. A comparison of the data to topological models that utilizes only the power spectrum can produce false positives by ignoring the additional information in the off-diagonal elements. We demonstrate this using Monte Carlo simulations. The power spectrum is rotationally invariant and does not specify orientation. We may thus modify Eq. \\ref{chisq_def} to use the power spectrum $C_\\ell$ and its covariance (Eqs. \\ref{power_spec_def} and \\ref{eq-Cl-covar}) in place of the spherical harmonic coefficients $a_{\\ell m}$. When only considering the power spectrum, the maximum likelihood for the WMAP data occurs at $L=1.1{\\Delta\\tau}$; this is the ``finite'' model power spectrum plotted in Fig. \\ref{fig1-powerspectrum}. Does this imply a positive detection of finite topology? To test this, we generate 1000 Monte Carlo realizations drawn from a parent population with $L=1.1$ and generate the likelihood for each realization using the full covariance matrix (Eqs. \\ref{eq-alm-covar} and \\ref{bcut_matrix_eq}). For such a small topology scale, almost all realizations have their likelihood peak at $L=1.1$. This scale is small enough that the bias from maximizing over orientation is not important. For each realization, we also generate a ``companion'' realization with exactly the same power spectrum, but with completely uncorrelated ${a_{\\ell m}}$'s. The two realizations by construction must give the same results for an analysis based solely on the power spectrum. When we analyze the ``companion'' realizations using the full $a_{\\ell m}$ covariance matrix, we obtain results similar to the infinite models displayed in Fig. \\ref{fig3-maxprob}. A likelihood analysis using the full $a_{\\ell m}$ covariance matrix successfully distinguishes models with compact topology from models with identical power spectra but without the correlations between different $a_{\\ell m}$ required by the topology. Suppression of power in the quadrupole and octopole moments is a necessary but not sufficient condition for a compact topology. The WMAP data show reduced power at $\\ell = 2$ and 3, but do not show the correlations expected for a compact topology and are indistinguishable from infinite models.  For cell size comparable to the distance to the decoupling surface, the correlations become weaker. Maximizing the likelihood over orientation then allows chance alignments to introduce a bias in the likelihood estimator. We quantify this using Monte Carlo simulations. The WMAP first-year data are consistent with input models drawn from parent populations with infinite fundamental domain. We establish 95\\% confidence limit $L > 17$ Gps for the cell size of a cubic topology, in agreement with the result of 24 Gpc obtained by \\cite{cornish/etal:2003}."
    },
    "0404/astro-ph0404150_arXiv.txt": {
        "abstract": "A parallax method to determine transverse velocity in a gravitationally lensed system is described. Using the annual motion of the Earth around the Sun allows us to probe the local structure of the magnification map that, under certain assumptions, can be used to infer the effective transverse velocity. The method is applied to OGLE data for QSO2237+0305 and the velocity value is estimated to be about $15\\pm 10\\,\\rmn{km}\\,\\rmn{s}^{-1}$ if attributed to the lensing galaxy or about $420\\pm 300\\,\\rmn{km}\\,\\rmn{s}^{-1}$ if attributed to the quasar. We find this estimate unreasonably small and conclude that we have not measured a parallax effect. We give a short list of properties that a system should possess to allow a successful implementation of this method. ",
        "introduction": "\\label{firstpage}  Among the various parameters needed to specify a microlensing model of a lensed system one of the most important is the overall transverse velocity. This is required in order to determine physical parameters from the observed temporal properties. Besides gravitational lensing, measuring the transverse velocities of galaxies is of considerable interest for studies of large-scale structure and kinematics \\citep{dekelreview, bernardeaureview}.  There have been a number of successful attempts to use the annual motion of the Earth around the Sun for measuring the transverse velocities of microlenses in our own Galaxy. The microlensing optical depth remains very low within the Local Group and, wherever double lenses are not involved, can be well  described by a Schwarzschild lens model. \\citet{gould92} developed a convenient formalism to describe the `annual parallax' effect in this case, that allows one to extract information on physical parameters of the lensing configuration, breaking some of the degeneracies inherent to the classical Paczy\\'nski light curve \\citep{paczynski}. The first microlensing event which showed a strong annual parallax signal was detected by the MACHO collaboration in the first-year data of the group's Galactic bulge program \\citep{alcock95} and six more were found among the longest events over the next seven years; two of these are currently among the best stellar mass black hole candidates \\citep{bennett02}. Events with an annual parallax signal were also found by other collaborations -- EROS/PLANET \\citep{an}, MOA \\citep{bond} and OGLE \\citep{ogle99, ogle01, ogledb}.  Perhaps the most spectacular results related to this effect are the detection of the first multi-peaked microlensing event by OGLE collaboration \\citep{oglemultipeak}, and direct observations of a lens based on parameters determined via this method \\citep{alcock01}.  Much effort has also been put into determination of the degeneracies present in the annual parallax effect description itself. It has been found that constant acceleration of the lens or the source can sometimes mimic the parallax signal \\citep{smp}, and there also exists a discrete degeneracy between jerk and parallax parameters that becomes a continuous one when the accuracy of observations decreases \\citep{gould04, park}. These degeneracies arise, basically, from the symmetry of the Schwarzschild lens that mixes different kinematic effects into a single geometric parameter -- the distance between the source and the lens expressed in the units of the Einstein-Chwolson radius (e.g., \\citealt{schneiderbook}). Some of these degeneracies are broken when the lens is a double object \\citep{an}, and as lens becomes more complex the above mentioned radius loses its unique status in the description of the lensing event. The single Schwarzschild lens approximation is, however, applicable to most of the microlensing events in the Galaxy due to the low optical depth to microlensing in our neighbourhood.  In contrast, microlensing in the images of strongly lensed quasars necessarily takes place in regions of high optical depth where the structure of the magnification map is complex and unknown, effectively rendering it impossible to model the light curve in detail. In addition, the intrinsic variability that quasars are expected to possess may be another contaminant of any microlensing signal present in the images' light curves.  For the particular multiply imaged quasar QSO2237+0305 \\citep{huchra} in which the microlensing phenomenon was observed for the first time \\citep{irwin}, a number of attempts have been made to determine the transverse velocity of the system. The small time delay differences between the four images -- less than a day \\citep{schneiderondelays} -- means that the observed fluctuations, uncorrelated between the images, are dominated by microlensing, with negligible intrinsic variations.  \\citet{wyithe} introduced a method to determine the transverse velocity based on the statistics of time derivatives of the microlensing-induced flux variations. Although their method seems to be the only one that makes it feasible to take into account proper motion of the microlenses, a number of parameters must be specified in order to make the velocity estimate, including the microlensing mass spectrum, which, in effect, translates temporal quantities into spatial ones. Under the assumptions made in their paper they find the transverse velocity to be less than $500\\,\\rmn{km}\\,\\rmn{s}^{-1}$ (at 95\\% confidence level) favouring, depending on the model, values in the range $60\\,\\rmn{km}\\,\\rmn{s}^{-1}$ to $400\\,\\rmn{km}\\,\\rmn{s}^{-1}$ -- a tighter constraint compared to previously assumed value of $\\sim 600\\,\\rmn{km}\\,\\rmn{s}^{-1}$ \\citep{wittmao}.  Another approach to estimating the effective velocity is to compare the spatial extent of `quiescent' regions of the magnification map models with the temporal extent of the periods of steady rise and fall in the actual light curves of the quasar images -- mostly, image D \\citep{quiescent}. However, all these methods are seriously dependent on the assumed microlensing parameters and deal with statistical properties of the light curves, which are difficult to establish with the currently available data.  In this paper we implement another, rather simple approach, underlain by a few natural assumptions. The essence of the method is the following: the light curve of an image of the lensed quasar is composed of the values of magnification on the observer's path through the magnification map (observer's plane). When the region of the map considered is small enough and is far from caustic curves, it is natural to expand the magnification as a function of the observer's position into a Taylor series and restrict ourselves to its linear terms. Where this approximation holds, the well-known motion of the Earth around the Sun can be used to obtain the local values of magnitude and direction of the gradient on this magnification map. Combined with the measured time derivative, they can be used to estimate the velocity of the Sun with respect to magnification map. We apply this method to the data available for QSO2237+0305.  In this analysis it is hardly possible to incorporate proper motions of the lenses and we will also neglect the intrinsic variability of the quasar; therefore it is unreasonable to expect a full accounting of all the observed features. However, as a (nearly) model-independent estimate of the velocity, this approach seems to be an interesting application of the microlensing phenomenon which could be extended to different microlensed systems if their properties are favourable.  In the next section we describe the method we use for probing the structure of the magnification map and obtaining the velocity estimate in greater detail, in section 3 the method is applied to the observational data for QSO2237+0305 obtained by the OGLE collaboration \\citep{oglehuchra} and the results are discussed: the value for transverse velocity obtained seems to us to be too low and we conclude that the method has failed in this case. Discussion of desirable properties of lensed systems, which may permit successful application of this method, concludes this study in section 4. ",
        "conclusions": ""
    },
    "0404/astro-ph0404199_arXiv.txt": {
        "abstract": "We present the results of U-band and multi-color photometry during the 2001 superoutburst of WZ Sagittae. Our 10 nights of U-band photometry span the time interval from HJD 2452118 to 2452197 while our multi-color observations range from HJD 2452115 to 2452197. The U-band light curves are generally in agreement with other datasets obtained during the superoutburst showing highly modulated light early on, rebrightenings, and superhumps of similar shape and period (except during the rebrightening peak). One of our multi-color datasets fortuitously covers the first rebrightening and allows determination of the accretion disk color temperature before, during, and after the event. It is seen that the rebrightening is a change from a neutral disk (T$\\sim$7000K) to an ionized disk (T$\\sim$10000K) and back again. We develop a simple limit cycle model for this behavior which approximately predicts the semi-periodic timescale observed for the rebrightenings. We discuss our results in relation to accretion disk structure during superoutburst. ",
        "introduction": "WZ Sagittae is often considered the quintessential short period dwarf nova, especially for the class of objects called TOADS (Tremendous Outburst Amplitude Dwarf novae) which only have infrequent (yearly to decadal) superoutbursts (see Howell et al., 1995). WZ Sagittae is the brightest TOAD at minimum light (V=15.0), has the longest intra-outburst timescale (20-30 years), is the closest cataclysmic variable at 43.5 pc (Harrison et al., 2004), and has one of the shortest known orbital periods (P=81.6 min).  Study of dwarf novae during superoutburst provide astronomers with information ranging from the binary orbital period to a mass estimate for the secondary star via detailed study of the period, shape, and time scale of so-called superhumps visible in the light curves (e.g., Patterson et al., 2002 and references therein). Superhumps are periodic hump-like modulations observed in the photometric signature of the star during superoutburst with periods a few percent different from that of the binary orbital period. Additionally, the morphology of the superoutburst itself, often combined with multi-wavelength observations, can provide a direct measurement of the properties of the accretion disk and its behavior during the outburst (e.g., Howell et al., 1999).  Superoutbursts are often observed via world-wide observer networks with the majority of observations being made by small telescopes and amateur astronomers. Their telescopes are generally equipped with CCDs and observations are routinely obtained in ``white light\". These unfiltered CCD images, when convolved with quantum efficiencies typical of the CCDs used, produce ``pink light\" observations. Only one TOAD, WX Cet, has multi-color photometric superoutburst observations (Howell et al., 2002) from which the authors concluded that the superhump period was grey in the optical and the observed period agreed with that determined solely by white light observations.  During August 2001, WZ Sagittae erupted in superoutburst 11 years prior to prediction and was observed by essentially every professional telescope available (both on the ground and in orbit) in addition to hundreds of amateurs around the world. A summary of the massive ground-based observational campaign obtaining CCD photometric observations and a detailed interpretation of their meaning is presented in Patterson et al. (2002).  During this same time period, we collected multi-color photometric measurements throughout the superoutburst as well as the first set of U-band time series observations for WZ Sagittae during a superoutburst. We present these data and interpret them in relation to other photometric observations of WZ Sagittae during superoutburst as well as relating the observations to accretion disk structure and the formation and cause of the semi-periodic rebrightenings. ",
        "conclusions": "We find that our U-band time series photometry generally agrees in light curve shape and period (both orbital modulations and superhumps) with the more detailed pink light results presented in Patterson et al. (2002). Our multi-color results show that the superoutburst has a blue nature to it early on (related to the high temperatures produced) and returns to a blue color in U-B and B-V near the end due to the emergence of the hot white dwarf. Additionally, we find that the rebrightenings appear to be a transition in a local disk volume from mostly neutral conditions to fully ionized gas and back. A toy model based on a cooling wave-temperature limit cycle provides a time scale consistent with observations.  Our limit cycle model is a subset of the typical dwarf nova accretion disk outburst limit cycle. Perhaps the hydrogen ionization limit cycle operates within accretion disks on various physical scale lengths leading to many different observed phenomena such as the semi-periodic modulations which we call disk rebrightenings. This type of cycle may also be the cause for the observed rapid outbursts in the ER UMa stars and similar type oscillations at late times in classical novae outbursts."
    },
    "0404/astro-ph0404366_arXiv.txt": {
        "abstract": "In the past few years, small scale anisotropy has become a primary focus in the search for source of Ultra-High Energy Cosmic Rays (UHECRs).  The Akeno Giant Air Shower Array (AGASA) has reported the presence of clusters of event arrival directions in their highest energy data set. The High Resolution Fly's Eye (HiRes) has accumulated an exposure in one of its monocular eyes at energies above $10^{19.5}$~eV comparable to that of AGASA.  However, monocular events observed with an air fluorescence detector are characterized by highly asymmetric angular resolution.  A method is developed for measuring autocorrelation with asymmetric angular resolution. It is concluded that HiRes-I observations are consistent with no autocorrelation and that the sensitivity to clustering of the HiRes-I detector is comparable to that of the reported AGASA data set.  Furthermore, we state with a 90\\% confidence level that no more than 13\\% of the observed HiRes-I events above $10^{19.5}$~eV could be sharing common arrival directions. However, because a measure of autocorrelation makes no assumption of the underlying astrophysical mechanism that results in clustering phenomena, we cannot claim that the HiRes monocular analysis and the AGASA analysis are inconsistent beyond a specified confidence level. ",
        "introduction": "Over the past decade, the search for sources of Ultra-High Energy Cosmic Rays (UHECRs) has begun to focus upon small scale anisotropy in event arrival directions. This refers to statistically significant excesses occurring at the scale of $\\leq2.5^\\circ$.  The interest in this sort of anisotropy has largely been fueled by the observations of the Akeno Giant Air Shower Array (AGASA).  In 1999 \\cite{Takeda:1999sg} and again in 2001 \\cite{Takeda:2001}, the AGASA collaboration reported observing what eventually became seven clusters (six ``doublets'' and one ``triplet'') with estimated energies above $\\sim3.8\\times10^{19}$~eV. Several attempts that have been made to ascertain the significance of these clusters returned chance probabilities of $4\\times10^{-6}$ \\cite{Tinyakov:2001ic} to 0.08 \\cite{Finley:2003on}.  By contrast, the monocular (and stereo) analyses that have been presented by the High Resolution Fly's Eye (HiRes) demonstrate that the level of autocorrelation observed in our sample is completely consistent with that expected from background coincidences \\cite{bellido,bellido2,apj}. Any analysis of HiRes monocular data needs to take into account that the angular resolution in monocular mode is highly asymmetric.  It is very difficult to compare the results of the HiRes monocular and AGASA analyses.  They are very different in the way that they measure autocorrelation. Differences in the published energy spectra of the two experiments suggest an energy scale difference of 30\\% \\cite{prl,Takeda:1998ps}. Additionally, the two experiments observe UHECRs in very different ways.  The HiRes experiment has an energy-dependent aperture and an exposure with a seasonal variability \\cite{prl}.  These differences make it very difficult get an intuitive grasp of what HiRes should see if the AGASA claim of autocorrelation is justified.   In order to develop this sort of intuition, we apply the same analysis to both AGASA and HiRes data. ",
        "conclusions": "We conclude that the HiRes-I monocular detector sees no evidence of clustering in its highest energy events.  Furthermore, the HiRes-I monocular data has an intrinsic sensitivity to global autocorrelation such that we can claim at the 90\\% confidence level that there can be no more than 3.5 doublets above that which would be expected by background coincidence in the HiRes-I monocular data set above $10^{19.5}$~eV.  From this result, we can then derive, with a 90\\% confidence level, that no more than 13\\% of the observed HiRes-I events could be sharing common arrival directions. This data set is comparable to the sensitivity of the reported AGASA data set if one assumes that there is indeed a 30\\% energy scale difference between the two experiments.  It should be emphasized that this conclusion pertains only to point sources of the sort claimed by the AGASA collaboration.  Furthermore, because a measure of autocorrelation makes no assumption of the underlying astrophysical mechanism that results in clustering phenomena, we cannot claim that the HiRes monocular analysis and the AGASA analysis are inconsistent beyond a specified confidence level."
    },
    "0404/astro-ph0404016_arXiv.txt": {
        "abstract": "{ Previous studies of young stellar objects (YSOs) have uncovered a number of associated parsec\\,-\\,scale optical outflows, the majority of which are driven by low\\,-\\,mass, embedded Class I sources. Here we examine more evolved Classical T Tauri stars (CTTSs), i.e. Class II sources, to determine whether these are also capable of driving parsec\\,-\\,scale outflows. Five such sources are presented here - CW Tau, DG Tau, DO Tau, HV Tau C and RW Aur, all of which show optical evidence for outflows of the order of 1pc (24$'$ at the distance of Taurus\\,-\\,Auriga). These sources were previously known only to drive ``micro\\,-\\,jets'' or small\\,-\\,scale outflows $\\la$ 1\\arcmin\\ in length. A parsec\\,-\\,scale outflow from a less evolved source (DG Tau B) which was noted in the course of this work is also included here. Examination of the five newly discovered large\\,-\\,scale outflows from CTTSs shows that they have comparable morphologies, apparent dynamical timescales and degrees of collimation to those from less evolved sources. There is also strong evidence that these outflows have blown out of their parent molecular clouds. Finally we note that the ``fossil record'' provided by these outflows suggests their sources could have undergone FU Orionis\\,-\\,type outbursts in the past. ",
        "introduction": "Herbig\\,-\\,Haro (HH) objects are the optically visible tracers of mass outflow from YSOs and are therefore ultimately powered by accretion \\citep{Cabrit90,Hartigan95}. Over the years many bipolar HH outflows have been observed and most were found to be driven by embedded, low\\,-\\,mass sources of $\\la$ 1\\Msun$\\!\\!$. Initially it was assumed that their lengths were only a fraction of a parsec, however the discovery of a $\\sim$ 1.4pc long outflow from RNO\\,43 \\citep{Ray87} hinted that this may not always be the case. In the mid 1990's it was realised that many of these outflows can have projected lengths much greater than 1pc \\citep{Bally96,Eisloffel97,Reipurth97}, reaching up to $\\sim$ 11pc. Some well known examples are the HH\\,34 outflow consisting of HH\\,33, HH\\,40, HH\\,85, HH\\,126, HH\\,34N, HH\\,34, HH\\,34X, HH\\,173, HH\\,86, HH\\,87 and HH\\,88 which is 3pc in projected length \\citep{Bally94}; the 5.9pc long HH\\,401, HH\\,1, HH\\,2 and HH\\,402 outflow \\citep{Ogura95}; and the HH\\,113\\,/\\,HH\\,111\\,/\\,HH\\,311 outflow at 7.7pc in length \\citep{Reipurth97}.  It is not surprising that such outflows can attain these lengths when we consider that they have tangential velocities of between 50 -- 200 kms$^{-1}$ \\citep{Devine97,Reipurth97} and the outflow phase, for even (low\\,-\\,mass) Class I sources, lasts at least 10$^5$ years. In reality, it should be {\\em expected} that most will attain parsec\\,-\\,scale lengths. The main observational hindrance in the past to observing them was the relatively small fields of view offered by most CCD cameras. With the advent of large CCD mosaics more and more parsec\\,-\\,scale outflows have been discovered.  The morphology of parsec\\,-\\,scale outflows yields valuable information about their driving sources. They are, in effect, a fossil record of the mass\\,-\\,loss history of their source over their dynamical timescales ($\\sim$10$^3$ to $\\sim$10$^5$ yr). They suggest, for example, quiescent phases between periods of violent mass ejection that give rise to the large HH complexes we see today. The morphology of an outflow can also indicate whether it is precessing and, if so, the rate of precession.  As mentioned earlier many of the parsec\\,-\\,scale outflows that have been observed to date are driven by young, Class I, low\\,-\\,mass YSOs. The classification scheme used here is based on the shape of the spectral energy distribution (SED) of the YSO from 10$\\mu$m to 100$\\mu$m \\citep{Lada84,Lada87}. The SED of a YSO can be modelled as an approximate blackbody with an infrared excess longwards of 2$\\mu$m due to circumstellar dust and gas. The infrared excess is very strong in the young, embedded Class I sources and is almost non\\,-\\,existent in the most evolved Class III sources. Here we observed a number of Class II low\\,-\\,mass sources -- CTTSs. These CTTSs were not previously associated with parsec\\,-\\,scale outflows; in fact many were only known to drive ``micro\\,-\\,jets'' of the order of $\\sim$ 5\\arcsec\\ to 40\\arcsec\\ ($\\la$ 0.03pc at a distance of 140pc to the Taurus\\,-\\,Auriga Cloud). Although outflows from these more evolved sources are not nearly as spectacular as those from Class I YSOs, their sources are no longer surrounded by significant amounts of dust and so their outflows can be traced right back to the origin. These Class II sources are still actively accreting and ejecting matter, albeit at rates 10\\,-\\,100 times smaller than Class I sources \\citep{Hartigan95}. In this paper, we present a number of parsec\\,-\\,scale outflows from CTTSs and we investigate whether these CTTSs show evidence for having undergone FU Orionis-like outbursts, based on the fossil record of their outflows.  Details about the observations are given in Section 2. In Section~\\ref{sec-resultsI} we report the discovery of parsec\\,-\\,scale outflows from five CTTSs in the Taurus\\,-\\,Auriga Cloud, at a distance of 140pc \\citep{Elias78,Wichmann98}. We also include serendipitous observations of a parsec\\,-\\,scale outflow from a less evolved Class I source. These results are discussed in Section~\\ref{sec-discussion} and our conclusions are presented in Section~\\ref{sec-conclusions}. ",
        "conclusions": "\\label{sec-conclusions} We have shown here that a number of (Class II) CTTSs, DG Tau, CW Tau, DO Tau, HV Tau C and RW Aur, which were previously known to drive ``micro\\,-\\,jets'' or short outflows of  $\\la$ 1$'$ ($\\la$ 0.04pc at the distance of the Taurus\\,-\\,Auriga cloud) in length, are actually capable of driving outflows with lengths of the order of 1pc. The serendipitous discovery of a 0.5pc long outflow from the Class I source DG Tau B is also reported here.  The morphological trends observed in the CTTSs outflows are comparable to those noted in younger sources i.e. increasing distance between successive HH objects coupled with increased size and complexity with distance from the source. In a few cases, small variations in the direction of propagation of the outflow have been found. The high degree of collimation of the five extended outflows from CTTSs compared well with that observed in the case of extended large\\,-\\,scale outflows from less evolved sources, suggesting that outflows remain focussed even as the source evolves from the Class I to the Class II stage. It is clear that the observed parsec\\,-\\,scale lengths of the CTTS outflows are minimum values and in reality they are much larger. These outflows all show evidence for having blown out of the parent cloud.  The apparent dynamical timescale of these extended outflows is typically a few times 10$^4$ years. This suggests a linkage between the major accretion events that give rise to the largest HH complexes and the FU Orionis phenomenon. As FU Orionis stars are CTTSs in quiescence, the extended outflows of the latter provide the best ``fossil record'' to test this linkage.  \\acknowledgement{ We thank the anonymous referee for helpful comments that clarified the presentation of these results. FMcG and TPR acknowledge support from Enterprise Ireland. This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France.}"
    },
    "0404/astro-ph0404246.txt": {
        "abstract": "{ We report a theoretical study of sulphur chemistry, as applied to hot cores, where S-bearing molecular ratios have been previously proposed and used as chemical clocks. As in previous models, we follow the S-bearing molecular composition after the injection of  grain mantle components into the gas phase. For this study, we developed a time-dependent chemical model with  up-to-date reaction rate coefficients. We ran several cases, using different realistic chemical compositions for the grain mantles and for the gas prior to mantle evaporation. The modeling shows that  S-bearing molecular ratios depend very critically on the gas temperature and density, the abundance of atomic oxygen, and, most importantly, on the form of sulphur injected in the gas phase,  which is very poorly known. Consequently, ratios of S-bearing molecules cannot be easily used as chemical clocks. However, detailed observations and careful modeling of both physical and chemical structure can give hints on the source age and constrain the mantle composition (i.e. the form of sulphur in cold molecular clouds) and, thus, help to solve the mystery of the sulphur depletion. We analyse in detail the cases of Orion and IRAS16293-2422. The comparison of the available observations with our model suggests that the majority of sulphur released from the mantles is mainly in, or soon converted into, atomic form. ",
        "introduction": "It is a long-standing dream to use relative abundances of different molecules as chemical clocks to measure the ages of astronomical objects. Studies of the ages of star formation regions have recently focused on S-bearing molecules. \\citet{1997ApJ...481..396C} and \\citet{1998A&A...338..713H} were the first to propose that the relative abundance ratios of SO, SO$_2$ and H$_2$S could be used to estimate the age of the hot cores of massive protostars. The underlying idea is that the main reservoir of sulphur is H$_2$S on grain mantles, and that when the hot core forms, the mantles evaporate, injecting the hydrogen sulphide into the gas phase. Endothermic reactions in the hot gas convert H$_{2}$S into atomic sulphur and SO from which more SO and, subsequently, SO$_{2}$ are formed, making the SO$_{2}$/SO and SO/H$_2$S ratios nice functions of time. These studies have triggered a variety of work, both observational and theoretical \\citep{1997ApJ...481..396C, 1998A&A...338..713H,2003A&A...399..567B}.  This line of research, however, has been challenged by recent ISO observations, which have cast doubt on the basic assumption that sulphur is mainly trapped in grain mantles as H$_2$S. The lack of an appropriate feature in the ISO spectra of high \\citep{2000ApJ...536..347G} and low \\citep{2000A&A...360..683B} mass protostars sets an upper limit on the mantle H$_{2}$S abundance which cannot exceed about $10^{-7}$ with respect to H$_{2}$ \\citep{1998ARA&A..36..317V}. Indeed, the identity of the major reservoir of sulphur in cold molecular clouds is a long standing and unresolved problem, for the sum of the detectable S-bearing molecules is only a very small fraction of the elemental S abundance \\citep{1994A&A...289..579T}. Since sulphur is known not to be depleted in the diffuse medium \\citep[e.g.][]{1994ApJ...430..650S}, it is usually assumed that sulphur in dense clouds is depleted onto the grain mantles rather than in refractory cores \\citep*[e.g. ][]{1994ApJ...421..206C},   but how this happens is a mystery. In a theoretical study, \\citet{1999MNRAS.306..691R} proposed that in collapsing translucent clouds sulphur is efficiently adsorbed onto grain mantles.  In fact, in these regions, most of the gas-phase sulphur is in the form of S$^{+}$, while grains are typically negatively charged, so that the collisional cross section for sulphur is enhanced compared with neutral species (e.g. O) and sulphur is removed from the gas phase more rapidly.  Another mystery is the form of sulphur on dust grains. The simplest possibility is that it consists of relatively isolated atoms, as would occur in a matrix, or perhaps as isolated pairs of atoms (S$_{2}$). Another possibility is that the sulphur is amorphous (or even crystalline), having formed islands of material from the initially adsorbed atoms.  Crystalline sulphur is known to come in two forms - rhombic and monoclinic - both of which consist of S$_{8}$ cyclic molecules.  Vaporization leads to a complex mixture of sulfur polymers through S$_{8}$ in complexity. If sulphur is elemental and amorphous, evaporation is also likely to lead to molecules of sulphur  through eight atoms in complexity. So far, the only S-bearing species firmly detected on granular surfaces is OCS, but with a relatively low fractional abundance of $10^{-7}$ \\citep{1997ApJ...479..839P}. Recently, \\citet{2002Natur.417..148K} claimed the detection of iron sulphide (FeS) grains in protoplanetary disks, but there is no evidence to suggest that solid FeS is the main form of sulphur in the parent collapsing environment. Actually, if the main form of solid sulphur is FeS, S should follow Fe depletion, which is not observed \\citep{1994ApJ...430..650S}. Even more recently, \\citet{2003MNRAS.341..657S} suggested that hydrated sulphuric acid (H$_{2}$SO$_{4}$ $\\cdot$ H$_{2}$O) is the main sulphur reservoir. In whatever form sulphur resides in the grain mantles, there is the possibility that the species, once evaporated, are very quickly destroyed to give atomic sulphur. In summary, although all the evidence is that sulphur is depleted onto grain mantles in cold clouds, its particular form is very uncertain.  Given the need for  chemical clock methods,  it is timely to reconsider the use of S-bearing molecules in this fashion. In this paper, we present a model with an up-to-date chemical network involving S-bearing molecules. We run several cases to cover a large, realistic parameter space for hot core sources, consistent with present observational constraints. Based on the results we obtain, we conclude that it is tricky to use abundance ratios of S-bearing molecules as chemical clocks in the absence of other constraints, for they depend more on the initial conditions, gas density, temperature, and the initial form of sulphur injected in the gas phase than on the age of the source.  The paper is organized as follows: we describe the model in \\S 2, the model results in \\S 3, and in \\S 4 we discuss the practical consequences of those results and apply the model to the specific cases of Orion and IRAS16293-2422. ",
        "conclusions": "%The fractional abundances and the abundance ratios of S-bearing molecules %have been used by several authors to estimate the age of hot cores %     \\citep[e.g. %][]{1997ApJ...481..396C,1998A&A...338..713H,2003A&A...399..567B}, %with models that contain the bulk of mantle sulphur in the form of %H$_2$S. %  Contrary to these studies, new ISO observations have shown that %solid H$_2$S is not %the main sulphur reservoir \\citep{1998ARA&A..36..317V}, %and no alternative major forms of sulphur %have been found, making the S--problem even more mysterious.  We have studied in detail the influence of the mantle form of sulphur on the post-evaporative gas-phase abundances of S-bearing molecules in hot star-forming regions, with the goal of understanding whether those molecules can be used to estimate ages. We considered four different reasonable mantle mixtures, from which gas-phase H$_2$S, OCS, S and S$_2$ emerge after a process of evaporation and, for the last two species,  possible rapid reaction, with different relative abundances, joining other species in the gas-phase prior to evaporation. We then followed the post-evaporative chemical evolution, with an emphasis on the abundance ratios of the main sulphur-bearing species for realistic physical conditions present in hot cores. Our results show that none of the ratios involving the four most abundant S-bearing molecules, namely H$_2$S, OCS, SO and SO$_2$, can be easily used by itself for estimating the age, because the ratios depend at least as strongly on the physical conditions and on the adopted grain mantle composition as on the time.  Also, the abundance of atomic oxygen in the gas phase, if not correctly accounted for, can seriously affect the chemistry. The situation, however, is not totally hopeless, because a careful comparison between observations and model predictions can give some useful hints on time estimates, and on the mantle composition. Such a careful analysis has to be done on each single source, however, for both the physical conditions and mantle composition can vary from source to source, so that the abundance ratios are not directly comparable. In practice, a careful derivation of the molecular abundances (which takes into account the source structure) coupled with a careful modeling of the chemistry at the right gas temperature and density is necessary.  We applied our model to two well studied hot cores: Orion KL and IRAS16293. For the S-bearing abundances towards Orion KL, we assumed that their emission arises from the hot core region (which is strongly debatable) and  is not beam-diluted. We were not able to reproduce all of the observed abundances ratios with any of our models. The agreement with Model 2 is satisfactory if we decrease the initial amout of atomic sulphur by a factor of 10. In that case, we derive a best age of $4\\times 10^3$~yr. However, the predicted abundance of CS is three orders of magnitude lower than the observed one. Contrary to the case of Orion KL, the sulphur-bearing abundances though the low mass hot core of IRAS16293-2422 have been carefully determined through a sophisticated model \\citep{2002A&A...390.1001S}, which takes into account the density and temperature structure of the source, as well as the abundance profile of each studied molecule. Using the standard value of cosmic ray rate, we found that Model 2, in which a large amount of atomic sulphur is initially present in the post-evaporative gas, best reproduces the observed abundance ratios. In that case, we derived an age of $\\sim 2\\times 10^3$~yr from the evaporation era to the current stage of this particular low mass hot core. If we decrease the initial amount of atomic sulphur in  Model 2 as for Orion KL, the agreement is still good and gives a similar age. This analysis favors the hypothesis that sulphur is mainly evaporated from the grains in the atomic form or in a form quickly converted into it. On the contrary, if a higher rate is used as suggested by the recent modelling of \\citet{2004doty}, best agreement occurs with Model 1, where no atomic sulphur can be found in the grain mantle and only H$_2$S and OCS are initially present. The strongest prediction of our atomic sulfur-rich model is the presence of large abundances of SO$_2$, derived from this form of sulfur, at late stages of hot cores. A futher systematic study of S-bearing-species towards older hot cores where the physical structure is well known would provide information to test this model. Moreover, the fact that not all of the sulfur need be initially in atomic form, given the reasonable agreement obtained using Model 2', suggests that a signficant portion of the granular elemental sulphur may be tied up in materials such as iron sulphide \\citep{2002Natur.417..148K}. %{\\bf Future investigations in other low mass protostars together with to constrain the cosmic ray ionization rate to"
    },
    "0404/astro-ph0404330_arXiv.txt": {
        "abstract": "We divide  some of the observed  diffuse interstellar bands (DIBs) into families which appear  to have spectral structures of single species. Three different methods are applied to separate such families, exploring the  best approach  for future investigations of this type. Starting with statistical treatment of the data, we found that statistical methods by themselves give insufficient results. Two other methods of data analysis (\"Averaging EWs\" and \"Investigating the figures with arranged spectrograms\") were  found to be more useful as tools for finding the spectroscopic families of DIBs. On the basis of these methods, we suggest some candidates as \"relatives\" of 5780 and 5797 bands. ",
        "introduction": "Diffuse interstellar bands (see e.g. Herbig  1995),  absorption structures  of  interstellar  origin   still   await explanation. The identification of the carrier(s) of  DIBs is one of the most difficult challenges for spectroscopists.  To the present day, huge amounts of  observational  data  on DIBs have been stored  by  astronomers  and  published  in  hundreds  of papers.  Unfortunately,  astronomical  data  still  do  not   meet sufficient   understanding   from  the point of view of laboratory spectroscopists.  One tries, in general, to solve the mystery of  the  carrier(s) of  DIBs  on  the   field   of   interdisciplinary   spectroscopic collaboration between molecular physicists, molecular chemists and astronomers. One expects that some progress on  this  field  will  be possible when all known DIBs (about  three  hundred  of  them have been discovered in visible light region) are divided  into  families  in such a way that only one carrier  is  responsible  for  all  bands belonging  to a given  family.  Such  families  of  bands  we  call 'spectroscopic' ones, to distinguish them from 'characterological' families isolated by the other authors (Chlewicki et.al. 1986, Kre{\\l}owski \\& Walker 1987, Josafatson \\& Snow 1987). All bands belonging to the spectroscopic family are, by  definition,  caused  by  the  same carrier. Bands belonging to the characterological family have some common characteristics (e.g. all are rather narrow), but  they  may  be  produced  by  different carriers.  To isolate spectroscopic families of bands, first of all one has to by-pass in some way the problem of the \"noisy correlation\". In  a  given  spectrogram  we  have  to  deal  with a complicated mixture of interstellar absorption structures. This is because the medium between the target star and the  observer  contains various species. For different directions (various  target  stars) we have to deal with different column  densities  of  interstellar matter giving contributions to the  spectra.  Intensities  of  all spectral lines (bands) measured in a spectrogram are well correlated with  the  column  density  of  relevant  matter and therefore also lines originated by different species are mutually correlated. Such correlation we call \"noisy correlation\". Of course not only differences in column densities of relevant interstellar matter may produce noisy correlation between DIBs. Other contribution to correlation of this kind may be given, for example, by mutually dependent astrochemical processes. Noisy correlation tells us nothing about spectroscopic families of bands.  The number of recognized DIBs has  grown  dramatically  in  recent  years, primarily due to better quality of  observational  material.  More and more weak DIBs (WDIBs)  seem  to  appear  every  time a  given spectral region is analysed carefully. One of the authors of  this paper (BW) spent a few years analysing spectra  for  new, very weak DIBs. It turned out to  be  of  great  importance  to reexamine the problem of DIB families. The aim of this paper is to explore further the properties of DIBs in the context of isolating families  of  the  structures.  We  first  describe  observational material which  revealed  the above mentioned  absorption  bands.  Then  we present the results  of the measuring  procedure  and  describe  DIB searching methods. In the last section of the paper we discuss the problem of some adequate method for separating spectroscopic  families of DIBs, and we pick out two probable \"relatives\" for 5780 band and four other ones for 5797 band. ",
        "conclusions": "As   mentioned  in the introductory  section,  the  main obstacle to isolating spectroscopic families of DIBs is the  noisy correlation. Due to a very high level (as the analysed  data reveals)  of noisy correlation in the considered data, the ability of statistical method to isolate spectroscopic families  are  very  limited.  Tight linear correlation, expected between members of the  same spectroscopic family, is efectively hidden by noisy  correlation and by measuring errors.  This  is evident  when  we  study results of statistical analysis  described  in the  previous  section. Looking at Table 5 only, we are not able to indicate which WDIB belongs to the same spectroscopic family as, for example, the  band  5797. On the other hand, when considering Table 4, we would be inclined to isolate, e.g, the family of strong lines: 6196, 6203,  6270,  6426, 6614, 6699, 5780, 5797; and this would be a mistake, since 5780  and  5797 belong to different families, as  mentioned in section 5.  Taking into account that statistical analysis requires  plenty of usable data and gives insufficient (for solving our problem by itself) results, it  is therefore not recommended as an appropriate tool for isolating spectroscopic families. [However, the statistical approach may be useful to distinguish linearly correlated bands  from these ones  which  are correlated  non-linearly.  Non-linearly  correlated  DIBs   should belong to different spectroscopic families. Also multidimensional statistical analysis could be useful in this case. Further study of this problem will be the subject of a separate paper (God\\l owski \\& Wszo\\l ek, in preparation)].  Methods described in subsections 6.2  and  6.3  are  much  more appropriate than the formal statistical aproach. They need a fewer  number of spectrograms and much  less  time  for  making  EW  (or  depth) measurements. These methods seem  also  to  give  valuable results. Using these methods we performed preliminary separation of two presumable spectroscopic families: (i) 5780, 5776  and  5795,  and (ii) 5793, 5797, 5819, 5829 and 5850.  Most probably, designated families are not complete yet. One cannot exclude also the possibility that we made wrong indications. In the case of almost constant ratios between column densities  of various DIB carriers in interstellar clouds, we have a chance to  get results quite similar those of the case when we have to deal with few spectral lines of the same  carrier. Further  investigation, based on better data samples and  involving  other  spectral ranges, is necessary to isolate  true  spectroscopic  families  of bands."
    },
    "0404/astro-ph0404040_arXiv.txt": {
        "abstract": "{We present a survey of all 3CR sources imaged with ISOCAM onboard the {\\it Infrared Space Observatory (ISO)}.  The sample consists mostly of radio--loud active galactic nuclei (AGN).  For each source, we present spatially integrated mid--infrared (MIR, $5 - 18\\mu$m) fluxes measured from newly calibrated ISOCAM images.  In total, we detected 68 objects of the 3CR catalogue, at redshifts $z \\le 2.5$, and obtained upper limits for 17 objects.  In addition, we detected 10 galaxies not listed in the 3CR catalogue.  The one with the highest redshift is 4C$+$72.26 at $z = 3.53$.  ISOCAM data are combined with other photometric measurements to construct the spectral energy distribution (SED) from optical to radio wavelengths.  The MIR emission may include synchrotron radiation of the AGN, stars of the host galaxy or dust.  Extrapolation of radio core fluxes to the MIR show that the synchrotron contribution is in most cases negligible. In order to describe dust emission we apply new radiative transfer models.  In the models the dust is heated by a central source which emits photons up to energies of 1keV. By varying three parameters, luminosity, effective size and extinction, we obtain a fit to the SED for our objects.  Our models contain also dust at large (several kpc) distance from the AGN.  Such a cold dust component was neglected in previous computations which therefore underestimated the AGN contribution to the far infrared (FIR).  In 53 cases ($\\sim 75$\\,\\% of our detected 3CR sources), the MIR emission can be attributed to dust. The {\\it hot dust} component is mainly due to small grains and PAHs. The modelling demonstrates that an AGN heating suffices to explain the ISO broad band data, starburst activity is not necessary.  In the models, a type 1 AGN is represented by a compact dust distribution, the dust is therefore very warm and emission of PAHs is weak because of photo--destruction.  In AGNs of type 2, the dust is relatively colder but PAH bands are strong. ",
        "introduction": "More than 20 years ago Kotanyi \\& Ekers (1979) pointed out that for radio--loud elliptical galaxies, the orientation of the (optical) dust lane appears preferentially perpendicular to the radio axis.  In unification theories dust plays an important role.  In all of them, a supermassive blackhole accretes gas from a disk and, in radio--loud objects, a jet is ejected parallel to the rotation axis.  The blackhole and the disk are surrounded by a torus of interstellar gas and dust which may block light from the central region on the way to the observer.  The diversity in the appearance of activity types is then explained foremost as a result of different viewing angles (Barthel 1989).     The obscuring dust torus is at the heart of unification theories.  It reaches towards the blackhole as close as dust can survive ($T \\simless 1500$\\,K).  As the torus is not directly seen on optical images, its dimensions are at most a few hundred parsec. Because of the immense bolometric luminosity emitted from the AGN, the dust in the torus, even at a distance of 100\\,pc, must be very warm ($ > 100$\\,K), and further in even {\\it hot}.  Such hot dust will emit preferentially in the mid infrared (MIR), at wavelengths covered by the ISOCAM filters.  Dust in radio galaxies is not restricted to the nuclear region, but is ubiquitous in the host galaxy.  It absorbs and scatters the light, but also radiates thermally as it is heated either by the nucleus or by stars.  Of course, the diffuse and spatially extended dust is much colder than the dust near the AGN.  Rapid star formation is also a potential heating source, besides the AGN. A detailed search for dust in the nuclear regions of 120 radio sources of the 3CR catalogue was carried out by de Koff et al.~(2000) using optical HST images at 0.7$\\mu$m center wavelength and of $0.1{''}$ resolution.  In one out of three galaxies, they found evidence for dust obscuration, with a large variety of morphologies such as disks, lanes or filaments.  Dust distributed smoothly on a large scale of several kpc is difficult to detect in absorption maps but reveals itself by infrared emission.  The dust shroud of an AGN is often not transparent and to model the emission spectrum one has to compute the radiative transfer.  Models for optically thick dusty nuclei and galaxies have been carried out in various approximations (Pier \\& Krolik 1993, Laor \\& Draine 1993 Kr\\\"ugel \\& Siebenmorgen 1994, Granato \\& Danese 1994, Efstathiou \\& Rowan--Robinson 1995, Siebenmorgen et al. 1997, Nenkova et al. 2002, Popescu et al. 2004). Previous AGN dust models underestimated the FIR as they did not include dust located at large distances from the center and thus missed the cold component.  To overcome the deficit emission in the FIR, these authors add an additional component to the models and attribute it to starburst activity (e.g.~Rowan--Robinson 2000).  In Section~4, we present AGN dust models where large scale dust emission is incorporated and where the geometry is radically simplified so that they can be described by only three parameters.  A cold dust component was discovered in 18 sources of the 3CR catalogue already by IRAS (Heckman et al. 1994) and more recently, in more objects, by ISOPHOT (Fanti et al.~2000, Meisenheimer et al.~2001, Van Bemmel et al.~2001, Andreani et al. 2002, Spinoglio et al.~2002, Haas et al. 2004).  These authors detected also a few sources in the MIR.  The hot dust responsible for it must be close to the nucleus and its detection serves as a further test for the unified model hypothesis.  This motivated us to survey the 3CR galaxies in the MIR by means of ISOCAM images available from the ISO archive.  The sensitivity of ISOCAM is $2 - 3$ orders of magnitude greater than IRAS at 12$\\mu$m.  In Section~2 and 3, we present the images, pertaining photometric measurements as well as results.  The spectral energy distributions (SED) are displayed in Section~4.  To interprete them, we apply radiative transfer models.  A discussion of generic SED properties in the infrared for different AGN types is given in Section~5.  The conclusions are given in Section~6.  In Appendix~{\\ref{galnotes.ap} we remark on individual galaxies (previous indications of dust, X-ray properties, jets, presence of high excitation lines, source morphology, companions).  In Appendix~\\ref{fits.ap}, we make for each galaxy a note on the radiative transfer models of Section~4.  Details of the ISOCAM observational set up is summarised in Tab.~\\ref{tab.obs}.  The fully reduced and astrometrically corrected ISOCAM images are shown as overlays on optical images in Fig.~\\ref{images}.    \\section {Survey data}  We present the data of our survey, describe which ISOCAM fields were extracted from the ISO archive, briefly overview the ISOCAM observing modes, and outline data reduction, source identification and photometric procedures.    \\subsection{Source selection} \\label{ident.sec}  The parent catalogue for our sample is the third update of the revised third Cambridge catalogue (3CR) by Spinrad et al.~(1985).  It contains 298 sources of which 195 are classified as radio galaxies and 53 as radio quasars.  Most of the objects in the 3CR catalogue are well observed at optical and near infrared wavelengths. We searched in the ISO archive for ISOCAM (Cesarsky et al. 1996) observations made with the long wavelength array (LW).  All ISOCAM images containing, potentially, at least, one 3CR source have been extracted. Altogether, we found 146 such ISOCAM observations. Often the same source was observed in multiple filter band passes. On the fully reduced ISOCAM images, we determined and improved the astrometric pointing by identifying objects with accurate coordinates listed in the SIMBAD Astronomical Database or the NASA/IPAC Extragalactic Database (NED). The images are presented in Fig.~\\ref{images}. We also made 10 serendipitous detections of galaxies which are not members of the 3CR catalogue. They are eye-ball detections at the position of NED coordinates.     \\subsection{Observations}  Most of the observations were done in the so called mini-raster mode where different ISOCAM detector pixels saw the same part of the sky (see ISO Handbook for a description of observing modes, Leech \\& Pollock, 2001). In addition, there are nine staring and one beam switch observation.  Three observations were performed in the circular variable filter (CVF) scan mode and one in the polarimetric mode.  Depending on the lens used in the experiment, the pixel scale of the 32 $\\times$ 32 element detector was 1.5$''$, 3$''$ and 6$''$ and the total field of view $48 \\times 48$, $96 \\times 96$ and $180 \\times 180$\\,arcsec$^2$, respectively. The most frequently applied broad band filter was LW10 with a bandwidth from 8--15$\\mu$m and central wavelength at 12$\\mu$m. In addition there are observations with filter: LW1 (4--5, 4.5$\\mu$m), LW2 (5--8.5, 6.7$\\mu$m), LW3 (12--18, 14.3$\\mu$m), LW7 (8.5--11, 9.6$\\mu$m) and CVF scans (5--17$\\mu$m) at spectral resolution of $\\lambda/\\Delta \\lambda \\approx 50$. In a single observing template often more than one filter was selected.  \\begin{table} \\caption{ISOCAM photometry of 3CR sources} \\label{tab.flux} \\begin{tabular}{lrrrr} \\hline \\hline & & & & \\\\ Name  & TDT & Band & Flux & RMS \\\\ &  &$\\mu$m  & mJy  & mJy \\\\ & & & & \\\\ \\hline 1 &2 &3 &4 &5  \\\\ \\hline & & & & \\\\ 3C006.1  & 70101081 & 12.0 & 0.9  & 0.2  \\\\ 3C013    & 80801283 & 12.0 & 0.7  & 0.2 \\\\ 3C017    & 57502102 & 12.0 & 4.0  & 0.4   \\\\ 3C018    & 61901003 & 12.0 & 9.1  & 0.6 \\\\ 3C020    & 59702305 &  12.0 & 5.3 & 0.4   \\\\ 3C022    & 78500882 &  12.0 & 7.1 & 0.4   \\\\ 3C031    & 58703801 &   4.5 &30 & 2   \\\\ & 40201422 &   6.7 &36 & 2   \\\\ & 58703793 &  12.0 &25 & 1  \\\\ & 58703801 &  14.3 &30 & 2   \\\\ 3C033.1 & 59702607 &  12.0 &18 & 1   \\\\ 3C048    & 43901804 &  14.3 &59 & 4   \\\\ 3C061.1  & 56201411 &  12.0 & 4.3 & 0.3 \\\\ 3C66B   & 45304902 &4.5 &14 & 1   \\\\ & 43502724 &6.7 &13 & 1   \\\\ & 45304902 &  14.3 & 3.8 & 1.6\\\\ 3C071 \t & 63301602 & 4.5 &  6860 & 340 \\\\ (NGC1068)& 63301902 & 6.8 &  14030 & 700   \\\\ & 63302202 & 14.9 & 50830 & 2540   \\\\ 3C076.1 & 46601603 &4.5 & 5.1 & 1.0   \\\\ & 46601603 &6.7 & 3.0 & 0.8   \\\\ & 46601603 &  14.3 & 1.0 & 1.5  \\\\ 3C079    & 61503807 &  12.0 &21 & 2 \\\\ 3C083.1 & 65901304 &4.5 &31 & 2   \\\\ & 65901032 &6.7 &47 & 3   \\\\ & 65901304 &  14.3 & 9.7 & 1.5\\\\ 3C084    & 61503617 &6.7 &  230 &20 \\\\ & & 9.6  &  400   &   70 \\\\ & &14.3  &  1000   & 100 \\\\ 3C098    & 63302405 &4.5 & 6.5 & 1.0  \\\\ & 63302405 &6.7 & 7.7 & 0.9   \\\\ & 63302405 &  14.3 &24 & 2   \\\\ 3C231    & 12300106  & 6.7  & 25000 & 1250 \\\\ (M82)    & 12300106  & 12.0 & 63000 & 3150   \\\\ & 12300106  & 14.9 & 68000 & 3400   \\\\ 3C249.1  & 21305001 & 12.0  & 19 & 1 \\\\ 3C265    & 22400201 &4.5 & 1.3 & 0.5   \\\\ & 22201802 &6.7 & 2.4 & 0.2   \\\\ & 22400303 &  12.0 & 6.0 & 1.6  \\\\ 3C270    & 22801205 &4.5 &  160 & 8  \\\\ & 22801205 &6.7 &95 & 5   \\\\ & 22801205 &  14.3 &38 & 3  \\\\ 3C272.1 & 23502406 &4.5 &  900 &  46 \\\\ & 23100414 &6.7 &  240 & 12   \\\\ & 23100414 &9.6 &  150 & 8   \\\\ & 23100414 &  14.3 &53 & 3   \\\\ 3C273    & 24100504 & 14.3 & 290 & 15 \\\\ & 24100504 &  6.7 & 190 & 10 \\\\ 3C274    & 23800308 &4.5 &  740 & 38 \\\\ (M87)    & 23800308 &6.7 &  260 & 13   \\\\ & 23901834 &  12.0 &51& 3   \\\\ 3C277.3  & 24500106 &  6.7 & 3.6 & 0.3  \\\\ & 24500106 &  14.3 & 8.5 & 0.5   \\\\ &  & & & \\\\ \\hline \\end{tabular} \\end{table}   \\setcounter{table}{0} \\begin{table} \\caption{ - continued.} \\begin{tabular}{lrrrr} \\hline \\hline & & & & \\\\ Name  & TDT & Band & Flux & RMS \\\\ &  &$\\mu$m  & mJy  & mJy \\\\ & & & & \\\\ \\hline 1 &2 &3 &4 &5  \\\\ \\hline & & & & \\\\ 3C286\t& 38800808 & 12.0 & 3.6& 0.3  \\\\ 3C288.1 & 24407233 &  12.0 & 2.3 & 0.6   \\\\ 3C293   & 61701307 &4.5 &16 & 1   \\\\ & 61701307 &6.7 &14 &  1   \\\\ & 61701413 &  12.0 &19 & 2  \\\\ & 61701307 &  14.3 &23 & 1   \\\\ 3C295   & 18001405 &6.7 & 1.4 & 0.3 \\\\ 3C296   & 27000606 &4.5 &30 & 2    \\\\ & 27000606 &6.7 &16 & 1   \\\\ & 27000606 &  14.3 &13 & 1  \\\\ 3C303.1 & 52900115 &  12.0 & 1.9 & 0.2 \\\\ 3C305   & 46300408 &4.5 &12 & 1   \\\\ & 12300205 &6.7 & 7.1 & 2.7   \\\\ & 51400760 &  12.0 &21 & 1   \\\\ & 12300205 &  14.3 &27 & 5   \\\\ 3C305.1 & 71702771 &  12.0 & 1.5 & 0.2   \\\\ 3C309.1 & 60001472 &  12.0 & 8.2 & 0.5   \\\\ & 60092200 &  6.7 & 4.4 & 0.5   \\\\ 3C319   & 54100619 &  12.0 & 1.3 & 0.2   \\\\ 3C321\t& 65800208 &  6.7 & 12 &  1    \\\\ & 65800208 & 12.0 & 27 &  1    \\\\ & 65800208 & 14.4 & 51 &  3    \\\\ 3C324   & 30300612 &6.7 & 0.1 & 0.2 \\\\ & 30300413 &  12.0 & 1.7 & 0.2   \\\\ 3C330   & 33600323 &6.7 & 1.3 & 0.2   \\\\ & 60001373 &  12.0 & 1.7 & 0.2   \\\\ & 33600323 &  14.3 & 2.6 & 0.5   \\\\ 3C332   & 60201725 &  12.0 & 9.2 & 0.5 \\\\ 3C336   & 30400442 &  12.0 & 2.4 & 0.5   \\\\ & 30492200 &  6.7 & 0.9 & 0.2   \\\\ 3C338  & 10601408 &4.5 &16 & 1  \\\\ & 10601408 &6.7 & 9.4 & 0.8   \\\\ & 51100561 &  12.0 & 8.8 & 0.5   \\\\ & 10601408 &  14.3 & 7.3 & 1.4   \\\\ 3C341   & 60201626 &  12.0 & 8.0 & 0.5   \\\\ 3C343  & 61900565 &  12.0 & 1.4 & 0.2   \\\\ 3C345 & 51100679 &  12.0 &22 & 1   \\\\ 3C346   & 64001327 &  12.0 & 4.2 & 0.3  \\\\ 3C349   & 52500429 &  12.0 & 3.4 & 0.3   \\\\ 3C351   & 52500130 &  12.0 &32 & 2   \\\\ 3C356   & 52100668 &  12.0 &0.6 & 0.2  \\\\ 3C371   & 52901232 &  12.0 &66 & 4   \\\\ & 55500476 &  14.3 & 91 & 5   \\\\ 3C379.1 & 52901433 &  12.0 & 3.6 & 0.3   \\\\ 3C380   & 54100969 &  12.0 &14 & 1   \\\\ 3C381   & 50304334 &  12.0 &19 & 1   \\\\ 3C382   & 52500562 &  12.0 &85 & 5   \\\\ 3C386   & 47100710 &4.5 &13 & 1   \\\\ & 47100710 &6.7 & 6.3 & 0.8   \\\\ & 51301035 &  12.0 & 5.6 & 0.3  \\\\ & 47100710 &  14.3 & 3.3 & 1.4   \\\\ 3C388   & 54100836 &  12.0 & 1.2 & 0.2 \\\\ 3C390.3 & 52901537 &  12.0 &94 & 5    \\\\ & &  & & \\\\ \\hline \\end{tabular} \\end{table}    \\setcounter{table}{0} \\begin{table} \\caption{ - continued.} \\begin{tabular}{lrrrr} \\hline \\hline & & & & \\\\ Name  & TDT & Band & Flux & RMS \\\\ &  &$\\mu$m  & mJy  & mJy \\\\ & & & & \\\\ \\hline 1 &2 &3 &4 &5  \\\\ \\hline & & & & \\\\ 3C402   & 49604639 &  12.0 & 3.4 & 0.5   \\\\ 3C410   & 53503041 &  12.0 &22 & 1   \\\\ 3C418   & 54101870 &  12.0 & 3.6 & 0.3 \\\\ 3C430   & 54301344 &  12.0 & 1.7 & 0.3   \\\\ 3C433 & 52201845 &  12.0 & 32.0 & 2 \\\\ 3C445 & 54500518 &6.7 &93 &5\\\\ & 54500518 &  14.3 &  210 & 11   \\\\ 3C449 & 54801112 &4.5 &16 & 1 \\\\ & 36701710 & 6.7 &10 & 1   \\\\ & 37403041 &  12.0 & 7.6 & 0.9 \\\\ & 54801252 &  12.0 & 8.7 & 0.5 \\\\ & 36701710 &  14.3 &11 & 2   \\\\ 3C452 & 54801053 &  12.0 &23 & 1 \\\\ 3C456 & 54700854 &  12.0 & 9.1 & 0.6    \\\\ 3C459 & 37500303 & 4.5  & 0.1 & 0.5 \\\\ &          & 6.0  & 5.5 & 0.9 \\\\ &          & 7.7  & 6.8 & 1.1 \\\\ &          & 14.9  & 29.5 & 4.0 \\\\ 3C465 & 39501112 &4.5 &22 & 1   \\\\ & 39501112 &6.7 & 8.1 & 0.8   \\\\ 3C469.1 &56201064 &  12.0 & 1.8 & 0.2 \\\\ & & & & \\\\ \\hline \\end{tabular} \\end{table}    \\begin{table} \\caption{ISOCAM $3\\sigma$ upper limits of 3CR sources} \\label{tab.upper} \\begin{tabular}{lrrr} \\hline \\hline &  &   & \\\\ Name  & TDT & Band & Flux \\\\ &     & $\\mu$m& mJy \\\\ &  &  &   \\\\ \\hline 1 &2 &3 &4  \\\\ \\hline &  &  &   \\\\ 3C002   & 41907203 &  12.0 &  $<$1.7 \\\\ 3C313   & 08800387 &  12.0 &  $<$6.2 \\\\ 3C314.1 & 52900417 &  12.0 &  $<$0.6 \\\\ 3C320   & 59501620 &  12.0 &  $<$0.5 \\\\ 3C343.1 & 39901248 &  12.0 &  $<$3.3 \\\\ 3C352   & 51800667 &  12.0 &  $<$0.6 \\\\ 3C357   & 51800831 &  12.0 &  $<$0.6 \\\\ 3C368   & 71200283 &  6.7  &  $<$1.5  \\\\ & 10200421 &  6.7  &  $<$8.1 \\\\ & 71200283 &  14.3 &  $<$2.1 \\\\ 3C401   & 53200838 &  12.0 &  $<$0.6 \\\\ 3C427.1 & 52902663 &  12.0 &  $<$0.6 \\\\ 3C434   & 53503146 &  12.0 &  $<$1.0 \\\\ 3C436   & 52201948 &  12.0 &  $<$1.0 \\\\ 3C437   & 53503288 &  12.0 &  $<$1.4 \\\\ 3C438   & 54000949 &  12.0 &  $<$0.5 \\\\ 3C442   & 53702911 &  14.3 &  $<$4.5 \\\\ 3C454.1 & 56001074 &  12.0 &  $<$0.6 \\\\ 3C460   & 56501357 &  12.0 &  $<$0.4 \\\\ &  &  &  \\\\ \\hline \\end{tabular} \\end{table}  Parameters of the observing templates are detailed in Tab.~\\ref{tab.obs}: target dedicated time (TDT) by which the observations are identified (column~1), celestial coordinate of the map center (column~2 and 3), observing date (column~4), band pass filter (column~5), lens (column~6), number of raster points $M$ (column~7) and raster lines $N$ (column~8), step between raster points $dM$ (column~9) and raster lines $dN$ (column~10), both in arcsec, and the median number of exposures taken on each raster position (column~11) together with the integration time per read--out (column~12).    \\subsection{Data reduction}   The data were reduced with the ISOCAM Interactive Analysis (CIA, Ott et al.~1996). We used the default data reduction steps of CIA: dark current subtraction, initial removal of cosmic ray hits (glitches), detector transient fitting, exposure coaddition, flat fielding and flux conversion to astronomical units. Except for staring modes, the coadded images at each raster position were sky projected and corrected for field distortion.  The dark current depends on the orbital position of the ISO space-craft and the temperature of the ISOCAM detector.  The applied correction is based on a model described by Roman \\& Ott (1998). The deglitching is done by following the temporal signal variation of a pixel using a multi--resolution wavelet transform algorithm (Starck et al.~1997).  The response of the detector pixels depends on the previous observations and there may be long term hysteresis effects for each detector element after changes of the photon flux level.  We applied the detector flux transient fitting method for ISOCAM data (Coulais \\& Abergel 2000).  After application of the default deglitcher some residuals of cosmic ray impacts were still visible in the data.  Therefore we applied after the detector flux transient correction a second cosmic ray rejection method which is basically a multi-sigma clipping of the temporal signal (Ott et al.~2000).  For raster observations, flat field estimation can be improved by an iterative method exploiting the fact that the same sky area is seen by different detector pixels (Starck et al.~1999).  The method works best for highly redundant data where a large number of pixels see the same sky.  In case of low redundancy rasters and for staring observations, we flat--fielded according to the standard calibration library.   \\subsection{Source photometry} \\label{photo.sec}  To determine the source flux, we perform a multi--aperture photometry. For each aperture centered on the brightest pixel of the source we determine the background as the mean flux derived in a 4 pixel wide annulus which is put 2 pixels away from the greatest ($\\sim 18''$) aperture. In this way, we obtain aperture fluxes that flatten with increasing aperture radius and approach an asymptotic value. For point sources the same procedure is repeated on a theoretical and normalised point source image (Okumura 1998) to find the correction factor of the multi--aperture analysis. The correction factor is typically 5\\%. The determination if a source is a point source or not is done by eye. The procedure will introduce an error of less than the correction factor for slightly extended sources misjudged to be a point source. In order to derive the statistical flux uncertainty, for the source aperture we quadratically add the RMS image and the 1$\\sigma$ uncertainty of the background estimate.  The systematic error of ISOCAM photometry is typically $\\simless 5\\%$ and color corrections of order 10\\% (see ISOCAM Handbook, Blommaert et al.~2001). ",
        "conclusions": "ISOCAM detected 80\\% (71 out of 88) 3CR sources.  We find evidence of hot dust in 53 radio galaxies or 75\\% of the sources detected in our survey.  For each detected source, we compiled photometry from optical to radio wavelengths.  The extrapolation of radio core fluxes to ISOCAM wavelengths shows that synchrotron emission from the core is usually negligible in the MIR, exceptions being flat spectrum radio sources.  The contribution of the host galaxy to the MIR is also generally small, here exceptions are a few sources of type FRI.  Thus for most objects the origin of the MIR can be attributed mainly to dust.  This dust is hot and heated by the central engine.  The emission comes from large grains with temperatures of a few hundred Kelvin and small grains and PAHs undergoing temperature fluctuations.  This picture is supported by simple radiative transfer calculations.  Most SEDs can be successfully fit between 1 and 1000$\\mu$m by a three parameter model where AGN luminosity, size of the galaxy and extinction are varied. Our modelling shows that the broad band data are consistent with AGN only heating of the dust.  For similar luminosities, we compare AGNs of type 1 and 2.  We find that the model parameters depend strongly on the AGN type and a dichotomy of the infrared SEDs is derived.  The IR fluxes of BLRGs and QSOs peak around 40$\\mu$m while the maximum of the dust emission for NLRGs is reached at $\\simgreat 100\\mu$m.  Therefore, the dust in type 1 AGNs is warmer than in type 2.  The models also predict much weaker PAH emission in type 1 AGNs than in type 2, as a result of evaporation.  Unfortunately, PAH bands cannot be resolved by the broad--band ISOCAM observations.  AGNs and starbursts are accompanied by tremendous IR luminosities. Although the possible coexistence of both energy sources and their relative contributions to the dust heating is difficult to assess, starbursts tend to favour PAH emission and this may serve as a possible discriminator of activity."
    },
    "0404/astro-ph0404276_arXiv.txt": {
        "abstract": "We present an analysis of the relative bias between early- and late-type galaxies in the Two-degree Field Galaxy Redshift Survey (2dFGRS) -- as defined by the \\etapar\\ parameter of \\citet{Madgwick_2002}, which quantifies the spectral type of galaxies in the survey. Our analysis examines the joint counts in cells between early- and late-type galaxies, using approximately cubical cells with sides ranging from 7$h^{-1}$Mpc to 42$h^{-1}$Mpc. We measure the variance of the counts in cells using the method of Efstathiou et al.~(1990), which we find requires a correction for a finite volume effect equivalent to the integral constraint bias of the autocorrelation function. Using a maximum likelihood technique we fit lognormal models to the one-point density distribution, and develop methods of dealing with biases in the recovered variances resulting from this technique. We use a modified $\\chi^{2}$ technique to determine to what extent the relative bias is consistent with a simple linear bias relation; this analysis results in a significant detection of nonlinearity/stochasticity even on large scales. We directly fit deterministic models for the joint density distribution function, $f(\\delta_{E},\\delta_{L})$, to the joint counts in cells using a maximum likelihood technique. Our results are consistent with a scale invariant relative bias factor on all scales studied. Linear bias is ruled out on scales less than $\\ell=28h^{-1}$Mpc. A power-law bias model is a significantly better fit to the data on all but the largest scales studied; the relative goodness of fit of this model as compared to that of the linear bias model suggests that any nonlinearity is negligible for $\\ell\\gtrsim40h^{-1}$Mpc, consistent with the expectation from theory that the bias should become linear on large scales. ",
        "introduction": "Measurements of large-scale structure from galaxy redshift surveys obviously measure the distribution of luminous matter only; the total mass distribution will be dominated by dark matter. The question of how the galaxies trace the total matter density field is therefore extremely pertinent, both to the estimation of cosmological parameters, and also as a probe of the physics of galaxy formation. A common assumption is `linear biasing', which can be expressed $\\delta_{g}=b\\delta_{m}$, where $\\delta_{g}$, $\\delta_{m}$ are the fractional overdensities relative to the mean in galaxies and mass respectively. This assumption becomes unphysical when $b > 1$ since, by definition, $\\delta_{g} \\ge -1$, but we can still define a bias parameter, $b(r)$, by e.g. $\\xi_{gg}(r)=b(r)^{2}\\xi_{mm}(r)$. Many of the constraints on cosmological parameters derived from large-scale structure measurements rely on an understanding of galaxy bias. Both the 2dFGRS power spectrum analysis (Percival \\etal~2001) and the constraints obtained for the neutrino mass (Elgar\\o y \\& Lahav 2003) assume scale independent bias. Joint constraints obtained by combining the 2dFGRS results with measurements of the CMB power spectrum (Percival et al.\\ 2002; Efstathiou et al.\\ 2002; Verde et al.\\ 2003) also require a model for galaxy bias. Dekel \\& Lahav (1999) show that nonlinearity and stochasticity in the bias relation can explain discrepancies between different methods of measuring parameters which assume a linear bias factor, such as measurements of $\\beta=\\Omega_{m}^{0.6}/b$ (Peacock et al.\\ 2001; Hawkins et al.\\ 2003).  In fact both theoretical approaches (Mo \\& White 1996) and simulations predict that bias may be non-linear and scale dependent, at least on some (small) scales. Kauffmann et al.~(1997) find only weak scale dependence on large scales and a bias relation consistent with linear bias. Benson et al.~(2000) find that semi-analytic galaxies in a LCDM model could reproduce the APM correlation function given a scale dependent bias taking the form of an antibias of galaxies relative to matter on small scales. Somerville et al.~(2001) also use semi-analytic modelling to demonstrate that the physics of galaxy formation introduces a small scatter in the galaxy--mass relation; they find the mean bias to have only a weak dependence on scale for $r\\lesssim 12h^{-1}$Mpc, (where the Hubble constant, $H_{0}=100$\\,$h$\\,km\\,s$^{-1}$).  In principle the true mass distribution can be directly measured from measurements of galaxy peculiar velocities using e.g.~POTENT reconstruction (Dekel, Bertschinger \\& Faber 1990). In practice accuracy is hard to achieve by such methods; the technique requires heavy smoothing since the error bars per galaxy are large and the volumes surveyed up to the present are relatively local. A useful probe is instead to compare the clustering of different types of galaxy: if these cluster differently, at least one type cannot exactly follow the mass distribution.  It has been known for some considerable time that galaxies of different morphological type have different clustering properties. Early-type galaxies, such as ellipticals or S0s, are highly clustered, accounting for almost 90\\% of galaxies in the cores of rich clusters. This fraction drops off steeply, however, with distance from the cluster cores and in the field 70\\% of galaxies are late-type galaxies: spirals and irregulars (Dressler 1980; Postman \\& Geller 1984). The level of fluctuations in each of the early- and late-type density fields can also be compared using the correlation functions or power spectra for the two sub-populations. This kind of study is optimised for small separations ($\\lesssim 10h^{-1}$Mpc) and has generally revealed that the clustering amplitude of ellipticals is greater than that of spirals by a factor of 1.3--1.5 (e.g.~Loveday et al.~1995; Norberg et al.~2002a; Madgwick et al.~2003b). If both density fields were perfectly correlated with the matter density field this factor would be equivalent to the ratio between linear bias parameters $(b_{E}(r)/b_{L}(r))^{2}$. There is also evidence that the relative bias between sub-populations of galaxies is more complex than the global galaxy bias. Measurements of the 2dFGRS bispectrum (Verde \\etal~2002) found no evidence for nonlinearity in the bias for 2dFGRS galaxies. More recently however, Kayo \\etal~(2004) find evidence for relative bias being complex on weakly non-linear to non-linear scales from a measurement of the redshift-space three-point correlation function, as a function of galaxy colour and morphology, in the Sloan Digital Sky Survey. Wild et al. (2004) have carried out a counts-in-cells analysis using volume limited samples from the 2dFGRS, and find evidence for non-linearity and stochastic effects.  A detailed framework for dealing with possible nonlinearities and stochasticity in the bias relation is given by \\citet{Dekel_1999}, based on the joint probability distribution of the galaxy and mass densities $f(\\delta_{g},\\delta_{m})$. In an analogous manner we will consider the joint probability distribution of the early- and late-type galaxy density fields for magnitude limited samples in the 2dFGRS. This approach in large part follows the methods described in Blanton (2000) for the Las Campanas Redshift Survey (LCRS), although the geometry of the 2dFGRS is considerably more amenable to this kind of study than that of the LCRS and allows us, for example, to examine a large range of scales.  This paper is organised into sections as follows. In Section~\\ref{survey} we summarize details of the 2dFGRS, the PCA-$\\eta$ parameter and the division into cells. We present a measurement of the variances of the counts in cells using the method of Efstathiou et al.~(1990), which we have corrected for integral constraint bias, in Section~\\ref{variances}. In Section~\\ref{dist_fns} we present an analysis of the one-point distribution of the counts in cells based on fits to a lognormal distribution function. In Section~\\ref{rel_bias} we discuss the relative bias. We present the results of applying the `modified $\\chi^{2}$' statistic of \\citet{Tegmark_1999} to the joint counts in cells and then move on to describe the application of the maximum-likelihood technique of \\citet{Blanton_2000} to constrain the relative bias between spectral types. We summarize our conclusions in Section~\\ref{concl}. ",
        "conclusions": "\\label{concl} In this paper we have presented a number of measurements of the relative bias between early- and late-type galaxies in the 2dFGRS derived using the counts in cells and the joint counts in cells for the separate galaxy populations. The behaviour of individual estimators for the linear relative bias parameter as a function of scale, as well as the relationship between different estimators of the linear relative bias parameter as a function of scale both have important implications for the scale dependence, nonlinearity and stochasticity of the relative bias between early- and late-type density fields. We have also used a power-law bias model as the simplest model including non-linear effects and demonstrated the characteristics of the best fit power-law bias model as a function of scale. \\subsection{Variances and the one-point distribution function} We have presented the variance of the counts in cells using the maximum-likelihood technique of Efstathiou et al.~(1990), which we have shown is subject to a significant bias when dividing the data into  redshift shells of low volume. We have shown that the method can be corrected for this integral constraint bias using the approximation of Hui \\& Gazta\\~naga (1999).  The one-point distribution of the counts in cells for early- and late-type galaxies, and the distribution for all $\\eta$-typed galaxies, has been fit by lognormal models, using a maximum-likelihood technique. The variances found using this technique are significantly biased on small scales when empty cells are included in the analysis, and we have been able to measure reliable variances only by fitting to counts in cells with empty cells removed. We have corrected our results on small scales to compensate for the inevitable bias resulting from the removal of empty cells. We find that the lognormal model is in general an adequate fit to the distribution functions, as measured by a Kolmogorov--Smirnov test. However the values for the variance implied by the best fit model parameters are slightly high in comparison with both predictions from the correlation functions and relative to the direct counts-in-cells variance measurements presented in this paper. The fact that this bias can be corrected by introducing a weighting scheme giving more weight to regions of higher density contrast suggests that the lognormal model is a relatively crude approximation to the true distribution. It is likely that a generalized lognormal model, such as the `skewed' lognormal model (SLNDFk) (Colombi 1994; Ueda \\& Yokoyama 1996), would be a better approximation. Unfortunately, the SLNDFk cannot be used in our maximum likelihood approach since it is not positive definite, and therefore is not strictly speaking a distribution function.  \\subsection{Comparison of relative bias parameters} \\label{last_bit}  \\begin{table} \\caption{Average bias parameters over all scales from $\\ell>14h^{-1}$Mpc for all of the measurements presented in the paper. Error bars are derived assuming measurements for adjacent bins in $\\ell$ scales are correlated.}  \\begin{center} \\begin{tabular}{ccc} \\hline Bias measurement & NGP & SGP \\\\ \\hline $1/b_{\\rm var}$ (from Efstathiou $\\sigma(\\ell)$) & 1.24$\\pm$0.06 & 1.26$\\pm$0.04 \\\\ $1/b_{\\rm var}$ (from $\\sigma_{\\rm LN}$ fits) & 1.28$\\pm$0.05 & 1.27$\\pm$0.04 \\\\ $1/b_{1,{\\rm lin}}$ (maximum likelihood) & 1.27$\\pm$0.04 & 1.17$\\pm$0.04 \\\\ $f\\approx1/b_{1,{\\rm lin}}$ (Tegmark test) & 1.28$\\pm$0.03 & 1.16$\\pm$0.03 \\\\ $1/b_{1,{\\rm PL}}$ & 1.36$\\pm$0.05 & 1.29$\\pm$0.04 \\\\ \\hline \\label{av_bias} \\end{tabular} \\end{center} \\end{table}  We present in Table~\\ref{av_bias} a comparison of the relative bias parameters from all of the measurements presented in the paper. We have averaged the bias measurements for all scales with $\\ell>14h^{-1}$Mpc which we expect to be unaffected by biases from empty cells. The error bars on each average are obtained assuming that the measurements in adjacent bins in $\\ell$ are perfectly correlated, which is a better approximation than assuming the measurements on separate scales are independent. Where relevant we have also used the more realistic error bars obtained from our Rayleigh-L\\'evy flight models.  As previously noted, the results for $1/b_{\\rm var}$ are consistent between regions and also consistent between measurements from direct variance estimation and fitting lognormal models to the one-point distribution.  Comparing the two estimates of the linear relative bias parameter $1/b_{1,{\\rm lin}}$, from the maximum likelihood method and from the Tegmark test, we find in both cases a significant discrepancy between NGP and SGP regions. The magnitude of this discrepancy is around 2-$\\sigma$. On the other hand the power-law bias measurements are approximately consistent between regions at a value of $b_{1,{\\rm PL}}$ which is further from unity. As we noted in Section~\\ref{bias} the assumption of linear bias when fitting to joint counts in cells which contain a significant degree of nonlinearity pushes the best fit relative bias closer to unity. This effect was also noted by Wild \\etal\\ (in preparation). It is likely that the apparent discrepancy between NGP and SGP linear bias parameters is also partly an artefact produced when non-linear joint distributions are fit with a linear bias model.  \\subsection{Scale dependence of the relative bias} In general, the relative bias is expected to be scale dependent on small scales ($r\\lesssim r_{0}$). The scale at which the bias relation becomes scale independent depends on the scales over which the biasing mechanism(s) operates. Non-local bias models (Bower et al.~1993, Matsubara 1999) are those on which the physical processes acting to produce the bias act on scales larger than those defined by the movement of massive particles, for example those models where radiation from QSOs has a significant effect. Local bias models (e.g.~Narayanan, Berlind \\& Weinberg, 2000) are those which are defined by some property of the local matter field, for example its density.  Narayanan, Berlind \\& Weinberg (2000) determine the variation with scale of a number of local and non-local bias models applied to N-body simulations. A general conclusion of this work is that local bias models are generically unable to influence the biasing relation on scales greater than $r=8h^{-1}$Mpc, which corresponds to $\\ell\\approx12h^{-1}$Mpc in this work. Although there does appear to be some variation of the best fit linear bias parameter on scales $\\ell\\gtrsim15h^{-1}$, when we factor in the larger error bars derived from models including selection function variation across cells in a more realistic way, the significance of any variation becomes negligible. Even if the variation in the linear bias parameter were significant it would not necessarily imply scale dependence of the bias since there is no significant scale dependence of the best fit power-law bias parameter. This illustrates the interdependence of non-linear, non-local and stochastic biasing effects. We conclude that any non-local contribution to the relative bias cannot be a dominant effect on large scales.  A special case of local relative bias which was considered in detail by Narayanan, Berlind \\& Weinberg (2000) is a local morphology density relation, of the type measured in the local environment of clusters and groups by e.g.~Postman \\& Geller (1984). A general conclusion for the relative bias produced by such a local effect is that the constant bias factor to which the scale dependent bias asymptotes on large scales is not equal to unity; assigning galaxy types based on local density produces a difference in clustering strength of the different galaxy types on all scales. Our results are fully consistent with this picture, which leads on to the question of whether the relatively well studied morphology-density relation can be held solely responsible for the relative bias measured in the 2dFGRS."
    },
    "0404/astro-ph0404510_arXiv.txt": {
        "abstract": "The physical limits to computation have been under active scrutiny over the past decade or two, as theoretical investigations of the possible impact of quantum mechanical processes on computing have begun to make contact with realizable experimental configurations. We demonstrate here that the observed acceleration of the Universe can produce a universal limit on the total amount of information that can be stored and processed in the future, putting an ultimate limit on future technology for any civilization, including a time-limit on Moore's Law. The limits we derive are stringent, and include the possibilities that the computing performed is either distributed or local. A careful consideration of the effect of horizons on information processing is necessary for this analysis, which suggests that the total amount of information that can be processed by any observer is significantly less than the Hawking-Bekenstein entropy associated with the existence of an event horizon in an accelerating universe. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404456_arXiv.txt": {
        "abstract": "We report the discovery of the minor planet 2003 VB12 (popularly named Sedna), the most distant object ever seen in the solar system. Pre-discovery images from 2001, 2002, and 2003 have allowed us to refine the orbit sufficiently to conclude that 2003 VB12 is on a highly eccentric orbit which permanently resides well beyond the Kuiper belt with a semimajor axis of 480$\\pm$40 AU and a perihelion of 76$\\pm$4AU. Such an orbit is unexpected in our current understanding of the solar system, but could be the result of scattering by a yet-to-be-discovered planet, perturbation by an anomalously close stellar encounter, or formation of the solar system within a cluster of stars. In all of these cases a significant additional population is likely present, and in the two most likely cases 2003 VB12 is best considered a member of the inner Oort cloud, which then extends to much smaller semimajor axes than previously expected. Continued discovery and orbital characterization of objects in this inner Oort cloud will verify the genesis of this unexpected population. ",
        "introduction": "The planetary region of the solar system, defined as the region that includes nearly circular low inclination orbits, appears to end at a distance of about 50AU from the sun at the edge of the classical Kuiper belt (Allen et al. 2001, Trujillo and Brown 2001). Many high eccentricity bodies from the planetary region -- comets and scattered Kuiper belt objects -- cross this boundary, but all have perihelia well within the planetary region. Far beyond this edge lies the realm of comets, which are hypothesized to be stored at distances of $\\sim 10^4$ AU in the Oort cloud.  While many objects presumably reside in this Oort cloud indefinitely, perturbation by passing stars or galactic tides occasionally modifies the orbit of a small number of these Oort cloud objects, causing them to reenter the inner solar system where they are detected as dynamically new comets (Oort 1950, Duncan et al. 1987), allowing a dynamical glimpse into the distant region from which they came. Every known and expected object in the solar system has either a perihelion in the planetary region, an aphelion in the Oort cloud region, or both.  Since November 2001 we have been systematically surveying the sky in search of distant slowly moving objects using the Samuel Oschin 48-inch Schmidt Telescope at Palomar Observatory (Trujillo and Brown 2004) and the Palomar-Quest large-area CCD camera (Rabinowitz et al. 2003). This survey is designed to cover the majority of the sky visible from Palomar over the course of approximately 5 years and, when finished, it will be the largest survey for distant moving objects since that of Tombaugh (1961). The major goal of the survey is to discover rare large objects in the Kuiper belt which are missed in the smaller but deeper surveys which find the majority of the fainter Kuiper belt objects (i.e, Millis et al. 2001).  In the course of this survey we detected an object with an R magnitude of 20.7 on 14 November 2003 which moved 4.6 arcseconds over the course of 3 images separated by a total of 3.1 hours (Figure 1). Over such short time periods, the motion of an object near opposition in the outer solar system is dominated by the parallax caused by the Earth's motion, so we can estimate that $R \\approx 150/\\Delta$, where $R$ is the heliocentric distance of the object in AU and $\\Delta$ is the speed in arcseconds per hour. From this estimate we can immediately conclude that the detected object is at a distance of $\\sim$100AU, significantly beyond the 50 AU planetary region, and more distant than any object yet seen in the solar system. The object has been temporarily designated minor planet 2003 VB12.  Followup observations from the Tenagra IV telescope, the Keck Observatory, and the 1.3-m SMARTS telescope at Cerro Tololo between 20 November 2003 and 31 December 2003 \\footnote{see http://cfa-www.harvard.edu/mpec/K04/K04E45.html for a table of astrometric positions.} allow us to compute a preliminary orbit for the object  using both the method of Bernstein and Khushalani (2000; hereafter BK2000), which is optimized for distant objects in the solar system, and a full least-squares method which makes no a priori assumptions about the orbit\\footnote{see http://www.projectpluto.com/find\\_orb.html}. Both methods suggest a distant eccentric orbit with the object currently near perihelion, but derived values for the semimajor axis and eccentricity are very different, showing the limitations of fitting an orbit for a slowly moving object with such a small orbital arc. For such objects a time baseline of several years is generally required before an accurate orbit can be determined. ",
        "conclusions": "Each of the plausible scenarios for the origin of the distant object predicts a specific dynamical population beyond the Kuiper belt. With only a single object, though, little dynamical evidence exists for preferring any one scenario. With any new discoveries in this region, however, evidence should quickly mount.  We can make a simple order of magnitude estimate of the ease of future discovery of objects in this population. We find a single distant object in our survey while we have found 40 Kuiper belt objects discovered to date in the survey. Assuming the size distribution of the distant population is the same as that of the Kuiper belt, other surveys should find similar proportions, assuming they are equally sensitive to slow motions. As of 15 March 2004, 831 minor planets have been detected beyond Neptune, we thus expect to have seen $\\sim$20 similar objects from other surveys. Even with this rough estimate, the lack of previous detection appears significant, suggesting either than most surveys have not been sensitive to motions as slow as $\\sim$1.5 arcseconds per hour or that there is an overabundance of comparatively bright objects in the distant population. In either case, it appears likely that new objects in this population should be detected reasonably soon.  The most plausible scenario for the origin of our object appears to be the dynamical effect of the creation of the solar system within a dense stellar cluster. In this scenario the Oort cloud extends from its expected location at $\\sim$100000 AU all the way in to the location of 2003 VB12. If this scenario is indeed correct the total mass of the Oort cloud must be many times higher than previously suspected. The expected population of large objects like the one discovered here is large. Our survey could only have detected this object during $\\sim$1\\% of its orbit, suggesting a population of $\\sim$100 objects on similar orbits. Moreover, if the population is nearly isotropic, $\\sim$5 more such objects must be observable in the current sky, with a total population of 500. Assuming a size distribution similar to the Kuiper belt, the total mass of this population is $\\sim$5 earth masses. The unseen population with ever more distant perihelia are likely even more numerous. With only the single object known in this population, extrapolation of a precise mass is not possible, nonetheless the existence of a nearby massive previously unsuspected inner Oort cloud appears likely. Even in the other origins scenarios a significant new mass must likely be present. At these distances and, in particular for isotropic distributions, current dynamical methods are unable to rule out any reasonable populations (Hogg et al. 1991). If the distant populations are sufficiently large, however, they may be detectable in future occulatation surveys.  While the genesis of 2003 VB12 is currently uncertain, continued discovery and orbital characterization of similar high perihelion objects should allow a unique and straightforward interpretation of this population. Each hypothesized formation mechanism leads to the prediction of a different dynamically distinct population in the outer solar system. Study of these populations will lead to a new knowledge of the earliest history of formation of the solar system."
    },
    "0404/astro-ph0404383_arXiv.txt": {
        "abstract": "When stars form within small groups (with $\\nstar \\approx 100 - 500$ members), their circumstellar disks are exposed to relatively little EUV ($h\\nu > 13.6$ eV) radiation but a great deal of FUV (6 eV $< h\\nu <$ 13.6 eV) radiation ($\\sim10^3$ times the local interstellar FUV field) from the most massive stars in the group. This paper calculates the mass loss rates and evaporation time scales for circumstellar disks exposed to external FUV radiation.  Previous work treated large disks and/or intense radiation fields in which the disk radius $r_d$ exceeds the critical radius $r_g$ where the sound speed in the FUV heated surface layer exceeds the escape speed; it has often been assumed that photoevaporation occurs for $r_d > r_g$ and is negligible for $r_d < r_g$.  Since $r_g \\gta 100$ AU for FUV heating, this would imply little mass loss from the planet-forming regions of a disk. In this paper, we focus on systems in which photoevaporation is suppressed because $r_d < r_g$ and show that significant mass loss still takes place as long as $r_d/r_g \\gta 0.1 - 0.2$. Some of the gas extends beyond the disk edge (or above the disk surface) to larger distances where the temperature is higher, the escape speed is lower, and an outflow develops. The resulting evaporation rate is a sensitive function of the central stellar mass and disk radius, which determine the escape speed, and the external FUV flux, which determines the temperature structure of the surfaces layers and outflowing gas. Disks around red dwarfs, low mass stars with $M_\\ast \\lta$ 0.5 $M_\\odot$, are evaporated and shrink to disk radii $r_d \\lta 15$ AU on short time scales $t \\lta 10$ Myr when exposed to moderate FUV fields with $G_0$ = 3000 (where $G_0$ = 1.7 for the local interstellar FUV field).  The disks around solar type stars are more durable. For intense FUV radiation fields with $G_0$ = 30,000, however, even these disks shrink to $r_d \\lta 15$ AU on time scales $t \\sim10$ Myr. Such fields exist within about 0.7 pc of the center of a cluster with $N_\\star \\approx 4000$ stars. If our solar system formed in the presence of such strong FUV radiation fields, this mechanism could explain why Neptune and Uranus in our solar system are gas poor, whereas Jupiter and Saturn are relatively gas rich. This mechanism for photoevaporation can also limit the production of Kuiper belt objects and can suppress giant planet formation in sufficiently large clusters, such as the Hyades, especially for disks associated with low mass stars. ",
        "introduction": "The collapse of molecular cloud cores leads to the formation of stars with orbiting accretion disks.  The dust in these disks can settle, coagulate, and form solid objects ranging in sizes from pebbles to planetesimals to planets.  However, a number of mechanisms act to disperse gas from these disks, either driving the gas back out into the interstellar medium, or spiraling it into the central star. During this dispersal, the gas can, in turn, drag small dust particles (with radii $b \\lta 1$ cm) along in the flow.  Gas dispersal thus disrupts planet formation in at least two important ways: (i) If the gas is dispersed before the disk dust particles have coagulated to sizes sufficient to decouple from the gas flow ($b \\gta 1$ cm), then the formation of planetesimals, Kuiper Belt Objects, and rocky planets will be curtailed because all the orbiting solid material in the gas flow region is removed before it has a chance to grow.  (ii) If the gas is dispersed before large ($m_P \\gta 5-15$ earth masses) rocky planets are formed, and if giant planets form by the gravitational accretion of gas onto these large rocky cores (Bodenheimer \\& Pollack 1986; Lissauer 1993), then the formation of gas giant planets like Jupiter and Saturn will be suppressed.  The dispersal of gas and small dust particles has other important effects on the formation of planetary systems. The presence of a moderately massive gas disk leads to planetary migration (Lin \\& Papaloizou 1986, Ward 1997). The presence of even small amounts of gas at time scales $t \\sim 10 - 100$ Myr after disk formation influences the dynamics and evolution of orbiting objects in solar systems. For example, such gas can affect the orbital eccentricities of both planets and any remaining planetesimals (e.g., Tanaka \\& Ida 1997; Kominami \\& Ida 2002; Chiang et al. 2002). Reducing the eccentricities can, in turn, alter the time required for the collisional agglomeration of large planets.  Observations of disk systems of various ages suggest that the small ($b \\lta 1$ mm) dust particles disappear on time scales of roughly 3 -- 10 Myr (Haisch et al. 2001). Near infrared continuum observations probe dust orbiting in the central regions of disks, $r \\lta 0.1$ AU, whereas submillimeter and millimeter wavelength continuum observations probe dust in the outer disks, $r \\gta 30$ AU.  The small dust grains in both regions disappear on roughly the same time scale. Presumably, after the disappearance time scale, some of the small dust particles have coagulated to form pebbles and larger objects that are no longer detectable at IR or millimeter wavelengths.  Several authors cite observational evidence for coagulated dust of size $b \\sim 1$ cm in young disks (d'Alessio et al. 1999, Throop et al. 2001).  However, the present observations do not specify how much of the small dust has been dispersed, and how much has coagulated into larger objects and disappeared from view.  Observations of gas in disks indicate that gas can also be dispersed in a relatively short time of only $t \\lta 10$ Myr (Zuckerman et al. 1995). Generally, the gas is traced by millimeter and submillimeter observations of the trace species $^{12}$CO in the J=1-0, 2-1, or 3-2 transitions, which are sensitive to low masses of gas as long as the gas disk is extended.  Because these lines are optically thick, and beam dilution reduces the observed intensity for these small ($\\lta 1\"$) disks, the current surveys are insensitive to gas of any mass at $r \\lta 50$ AU for the nearby ($d \\sim 100$ pc) young star/disk systems.  Nevertheless, significant gas masses have been detected via CO observations in disks as old as 10 Myr (cf. Carpenter 2002). In short, from an observational point of view, it appears that the bulk of the gas is dispersed from the outer disks in time scales $t \\lta 10$ Myr, but the evolution of the gas in the inner, planet-forming region of the disk is uncertain.  Hollenbach, Yorke, \\& Johnstone (2000) reviewed theoretical models for dispersing the gas and small dust from disks.  Observational evidence from our solar system and other planetary systems indicates that more gas and dust is accreted onto the central star and dispersed back into the ISM than forms planets or other solid orbiting objects. This dispersal is dominated by photoevaporation in the outer regions of disks and viscous evolution (accretion onto the star coupled with protostellar outflows) in the inner parts of disks. The boundary between these two regimes -- viscous evolution and photoevaporation -- remains uncertain.  We need to develop a better understanding of viscous accretion and develop better photoevaporation models that accurately track mass loss at moderate radii from the central star. This paper addresses the latter problem for the case of external irradiation.  Stellar winds may play a significant, but probably not dominant, role in dispersing gas at moderate radii near the boundary of the inner viscosity-dominated region and the outer photoevaporating region.  Nearby stellar encounters, even for star/disks born in dense clusters like the Trapezium cluster, only affect the outermost regions ($r\\gta 100$ AU) of the largest disks, and, even there, the photoevaporation of disks in these same clusters is likely to dominate the dispersal of the outer regions (e.g., Scally \\& Clarke 2001, Clarke 2002, Adams \\& Laughlin 2001).  Photoevaporation occurs when energetic photons heat the surface of the disk to elevated temperatures. The radiation of interest includes FUV photons in the energy range 6eV -- 13.6eV, EUV photons in the energy range 13.6eV -- 100eV, and X-rays in the energy range 100eV -- 10keV. If EUV photons can penetrate the outward flow and reach the disk surface, they will ionize and heat the surface to $T \\approx 10^4$ K, whereas the FUV and/or X-ray photons tend to heat the neutral gas to lower temperatures, typically in the range 100 K $< T < $3000 K. The thermal pressures in these heated regions drive the gas outward and create a flow into the interstellar medium. An important critical radius $r_g$ can be defined -- this fiducial length scale is the radius where the sound speed of the gas (hydrogen atoms) equals the escape speed from the gravitationally bound system, i.e., \\be r_g = {G M_\\ast \\muc \\over k T} \\approx 100 {\\rm AU} \\, \\Bigl( {T \\over 1000 {\\rm K} } \\Bigr)^{-1} \\, \\Bigl( {M_\\ast \\over 1 M_\\odot } \\Bigr) \\, , \\label{eq:rcrit} \\ee where $M_\\ast$ is the mass of the central star and $\\muc$ is the average mass of the gas particles. Most previous work on photoevaporation (Hollenbach et al. 1994, Johnstone et al.  1998, St\\\"orzer \\& Hollenbach 1999) assumed that photoevaporation flow is only active for $r > r_g$, and that the disk is static (with a warm surface corona held in orbit by the stellar gravity) for $r < r_g$. In this paper, we generalize this picture to include a more proper treatment of the flow hydrodynamics and show that significant photoevaporation can take place for smaller radii, $r \\gta 0.2 r_g$ (see Figure 1). In any event, disk photoevaporation can be considered like a slow ($v \\sim 1-5$ km s$^{-1}$) thermal Parker wind originating from the outer portion of the disk ($r \\sim 3 - 100$ AU).  Other authors have discussed the possibility of significant flow from $r<r_g$ for disks surrounding compact objects (e.g., Begelman, McKee \\& Shields 1983; Woods et al. 1996). This work showed that significant photoevaporation can take place outwards from $r \\sim 0.2 r_g$ for the case of X-ray heated disks around black holes. Recently, Liffman (2003) presented an analytic argument for photoevaporative flow inside of $r_g$. Flow inside $r_g$ is important for protoplanetary disks.  Often the heating raises the disk surfaces to $T \\lta 1000$ K, so that $r_g \\gta 100$ AU. If photoevaporative flows are still significant at $r \\sim 0.2 r_g$, then photoevaporation can effectively remove gas and dust from the region near 20 AU and thereby affect the formation of Uranus, Neptune, and Kuiper Belt Objects in our solar system.  Photoevaporation can be initiated by the energetic photons from the central star or from a nearby, more massive, and luminous star in the stellar birth cluster.  Hollenbach et al. (1994) originally modeled the evaporation caused by EUV photons from the central star. Johnstone, Hollenbach \\& Bally (1998) and St\\\"orzer \\& Hollenbach (1999) presented the first semi-analytic models of disks (with $r_d > r_g$) around low mass stars being photoevaporated by the FUV and EUV fluxes from a nearby OB star.  These models were successfully applied to the PROtoPLanetarY DiskS, or ``proplyds'', observed in the cluster of low mass stars around the Trapezium in Orion (e.g., O'Dell 1998, Bally et al. 1998, Churchwell et al. 1987). A complementary set of models (Richling \\& Yorke 1997, 1998, 2000; Yorke \\& Richling 2002) studied the hydrodynamical flow for disks subjected to both radiation from their central stars and external radiation.  This previous work produced two results of interest here: (i) In the case of external illumination, the FUV photons often initiate the mass loss and the incident EUV flux is absorbed at an ionization front in the neutral flow which is several disk radii away from the disk surface. (ii) The externally-illuminated disks evaporate from outside in, whereas the bulk of the mass loss for internally-EUV-illuminated disks occurs at $r \\sim r_g$. In other words, in the former (external) case, a disk with outer radius $r_d$ shrinks from $r_d > r_g$ to $r_d \\lta r_g$ as evaporation proceeds.  In the latter case, in the absence of turbulent viscosity to drive radial flow and replenish material at $r_g$, the disk evaporates at $r_g$ until a gap is formed there, and then the photoevaporation proceeds from $r_g$ outward to $r_d$.  These early models effectively assumed large disks with $r_d > r_g$.  This paper presents a more in depth treatment for the case of small disks (with $r_d < r_g$) that are externally illuminated by FUV radiation. However, this work also has important implications for photoevaporation at $r<r_g$ for the internally-illuminated disks.  Several recent papers combine these early photoevaporation models with models of viscous accretion in attempts to model the time evolution of the dispersal of the entire protoplanetary disk. Clarke et al. (2001) treat the EUV photoevaporation by the central star coupled with viscous accretion and evolution to explain why disks are observed to rapidly disperse at the end of their lives on a time scale that is a small fraction of the disk lifetime. Matsuyama et al. (2003ab) model disks with both internal EUV and external FUV and EUV illumination and with viscous evolution.  Such models will need modification in light of the results of this paper.  Stars often form in groups or clusters. If these stellar aggregates are large enough ($\\nstar \\gta 200$ stars), the system has a good chance of containing at least one O or early B star. In such systems, the low mass stars in the cluster are subject to significant photoevaporation by the FUV flux from other, larger stellar members. The ultimate goal for the work presented in this paper is to calculate the probability that a given low-mass star/disk system in the Galaxy, like the early solar nebula, experienced sufficient external illumination and consequent photoevaporation to affect the formation of gas-giant planets and/or the formation of planetesimals and planets in the Kuiper Belt region.  Such a calculation would require the knowledge of the probability of being born in a cluster of given size $\\nstar$ (e.g., see Lada \\& Lada 2003, Porras et al. 2003, Adams \\& Myers 2001, Carpenter 2000), the probability that one or several high mass stars are members of this cluster (given by the stellar initial mass function), the delay time between low mass star formation and high mass star formation, and the fraction of time that the low mass stars lie at a given distance from the unembedded OB star. Armitage (2000) presents a first attempt at such a model, which assumes that EUV-induced photoevaporation operates only for $r_d > r_g$.  This paper presents the corresponding analysis for the case of FUV-induced photoevaporation that occurs for $r_d < r_g$; this latter process often dominates the mass loss for typical disks in typical star formation environments.  This paper is organized as follows. In \\S 2, we discuss the physical mechanisms in photoevaporating disks, including heating processes, dust properties and attenuation, cooling mechanisms, thermal balance, and chemistry. We then summarize in more detail (in \\S 3) the previous results for the photoevaporation of ``large'' (supercritical) disks with $r_d > r_g$, since these analytic results will be useful for generalization to the case of smaller disks. In \\S 4, we calculate the photoevaporative mass loss rates and time scales for subcritical disks (with $r_d < r_g$) due to external FUV illumination. In general, photoevaporation takes place on both the disk surface, creating an initially vertical flow, and from the disk edge at $r_d$, creating a radial flow.  Although the disk edge has less area, the radial flow tends to dominate the mass loss because the material here is bound more weakly.  In a previous paper (Hollenbach \\& Adams 2003), we presented the isothermal case, where we can obtain analytic approximations which provide physical insight; here we develop the more complicated (but more realistic) non-isothermal case where the temperature is determined from the heating and cooling of the gas in the flow. We determine how the mass loss rate depends on the incident FUV flux, the size $r_d$ of the disk, and the mass of the central star.  We apply these results to the possible evaporation of the early solar nebula (\\S 5), the formation of Kuiper Belt objects and debris dust (\\S 6), the suppression of giant planet formation in large clusters like the Hyades (\\S 7), and the evaporation of disks around low mass stars (\\S 7).  We conclude, in \\S 8, with a summary and discussion of our results. ",
        "conclusions": "In this paper, we have studied the photoevaporation of small circumstellar disks ($r_d < r_g \\sim 100$ AU) due to the heating by FUV radiation from the stellar birth environment. Because this work applies to small disk radii, we can determine the effects of photoevaporation on inhibiting planet or planetesimal formation in the disk region where $r = 10 - 100$ AU.  This work complements previous studies, which have considered the evaporation of circumstellar disks due to EUV radiation from their parental stars (e.g., Shu et al. 1993) and the evaporation of large disks ($r_d \\gta 100$ AU) due to UV radiation in large clusters like the Trapezium (e.g., SH99).  We show that FUV photoevaporation is likely to dominate EUV evaporation both in large clusters (e.g., $N_\\star \\approx 4000$, $G_0 \\approx 30,000$) and in more moderate sized groups (e.g., $N_\\star \\approx 300$, $G_0 \\approx 3000$), until the disks shrink to sizes $r_d \\lta 10$ AU. By the time disks evaporate to such small radii (on time scales $t \\gta 30$ Myr), the major episodes of planet formation are expected to be over, so that EUV photoevaporation does not generally play an important role in affecting planet formation.  [1] For solar type stars, with $M_\\ast \\approx 1 M_\\odot$, relatively intense FUV radiation fields are required for significant photoevaporation to take place. In particular, FUV radiation with $G_0$ = 30,000 will efficiently evaporate disks with radii down to $r_d \\sim 20$ AU on time scales of $\\sim10$ Myr.  The outer parts of these circumstellar disks can be effectively evaporated through the action of this level of FUV radiation, which is expected to be present in the cores of dense stellar clusters (e.g., $d \\lta 0.7$ pc with $N_\\star \\approx 4000$).  [2] In our own solar system, the relative paucity of gas in Neptune and Uranus can be understood if the outer solar nebula ($r \\gta 20$ AU) is stripped of its gas before the planets complete their formation (\\S 5). The action of FUV radiation can remove enough gas on a sufficiently rapid time scale if the early solar system is exposed to FUV radiation fields with intensity $G_0 \\ge 30,000$. We expect such strong FUV radiation fields to be somewhat rare.  [3] FUV radiation fields can affect the formation of Kuiper belt objects and other rocky bodies in the outer portion of our solar system, and others. In these systems, dust grains coagulate as they settle and eventually grow too large to be removed from the disks. This process competes against evaporation, which acts to remove gas and dust from the disk. We find that dust coagulation tends to take place more rapidly (than mass loss) for radii less than a cutoff radius $r_c \\approx 100$ AU, even in relatively harsh stellar birth environments (\\S 6). As a result, Kuiper belt objects, and the debris dust that they generate later on, can be formed (out to $r\\sim100$ AU) around most stars. However, we cannot completely rule out the possibility that photoevaporation in the solar nebula could have produced the observed cutoff in Kuiper Belt objects at $r_c \\sim 50$ AU.  [4] Relatively large clusters contain B stars (and even O stars) with high probability. Sufficiently rich clusters thus provide a hostile environment for giant planet formation because the FUV radiation from the background cluster is effective at removing gas from nebular disks.  Applying this result to known clusters, such as the Hyades (\\S 7), we find that giant planet formation can be compromised in such environments.  [5] We have calculated (numerically) mass loss rates $\\dot M$ as a function of stellar mass $M_\\ast$, disk radius $r_d$, and FUV radiation field $G_0$. We also provide a simple analytic solution that approximately shows the scaling of the mass loss rate with these parameters. However, the analytic results are presented in terms of the column density $N_C$ of the heated surface gas, which is assumed to be isothermal with sound speed $a_s$. Comparison to PDR codes is required to determine $N_C$ and $a_s$ for a given radiation field $G_0$.  [6] The mass loss rate is significant for disk radii much smaller than the critical radius, in particular for $r_d/r_g \\gta 0.15$.  Previous work assumed negligible mass loss for $r_d < r_g$, so this finding increases the range of viable parameter space for mass loss. However, the mass loss rate drops exponentially for $r_d \\lta 0.15 r_g$, scaling roughly as ${\\dot M} \\propto \\exp[-r_g/2r_d]$.  [7] If a disk has enough viscosity, then viscous spreading of the outer disk edge can affect photoevaporation. As a disk becomes smaller in radius, its photoevaporation time increases whereas its viscous spreading time decreases. As a result, disks will shrink down to the size at which the two time scales are in balance (see Figure 8).  This process tends to enhance the effectiveness of photoevaporation by feeding new material into the outer disk where it can be efficiently removed by the outflow.  [8] Photoevaporation is most effective for disks surrounding stars of low mass (\\S 7). For example, a disk around an M dwarf with $M_\\ast$ = 0.25 $M_\\odot$ can be evaporated down to 10 AU in only 12 Myr when exposed to a modest FUV radiation field with $G_0$ = 3000. Such radiation intensities occur readily in moderately sized stellar groups, those with $N_\\star \\sim 300$, which represent a common star forming environment (e.g., Lada \\& Lada 2003, Porras et al. 2003).  A intriguing result emerges from this consideration of disk evaporation and the corresponding loss of planet forming potential for stars with varying mass. High mass stars are efficient at evaporating their own circumstellar disks and are thus not expected to harbor planets. At the other end of the mass spectrum, red dwarfs easily lose their disks due to photoevaporation in the presence of modest external FUV radiation fields (e.g., $G_0$ = 3000), which are expected in common star forming units. As a result, solar type stars (loosely speaking, stars with masses within a factor of two of 1.0 $M_\\odot$) are the preferred locations for giant planet formation.  \\bigskip \\centerline{\\bf Acknowledgments} \\medskip  We would like to thank M. Kaufman and A. Parravano for useful discussions. We also thank an anonymous referee for comments that clarified the paper. This work was supported by a grant from the NASA Origins of the Solar System Program, the NASA Astrophysics Theory Program, and by the Michigan Center for Theoretical Physics.  \\newpage \\centerline{\\bf APPENDIX: ANALYTIC APPROXIMATION} \\centerline{\\bf FOR THE PHOTOEVAPORATION OF SMALL DISKS ($r_d \\ll r_g$)} \\bigskip  In this Appendix, we derive simple analytic results for the scaling of the photoevaporative mass loss rate $\\dot M$ as a function of stellar mass $M_\\ast$, disk radius $r_d$, and (implicitly) the strength of the FUV radiation field $G_0$. Specifically, we make the following simplifying assumptions:  [i] The gas is essentially static in the inner region where $r < r_s$, with thermal pressure balancing gravity. This assumption is equivalent to neglecting the $v dv/dr$ term in equation (\\ref{eq:forcezero}) and solving the remaining equation for the density structure $n(r)$.  [ii] The outflow velocity $v$ is constant in the outer region where $r > r_s$, with $v=a_s$, the sound speed at the sonic point. This assumption implies that the density profile in the outer region has the form $n(r) = n_s (r_s/r)^2$, where $n_s$ is the number density at $r_s$.  [iii] The FUV field $G_0$ heats a column $N_C$ of surface gas to a constant temperature $T_s$ (i.e., the surface layer is isothermal). We thus obtain results that depend on $T_s$ ($a_s$) and $N_C$, but these parameters are actually surrogates for the radiation field $G_0$.  We can relate $a_s$ and $N_C$ to $G_0$ using the results from the PDR code as shown in Figure 2.  Figure 1 shows a schematic representation of the photoevaporation process for subcritical disks. Here, we work in the limit $r_d \\ll r_g$ and assume that the critical radius $r_g$ and the sonic radius $r_s$ are comparable ($r_s \\sim r_g$). With this set of approximations, the density profile for the subsonic region takes the form $$ n(r) = n_d \\exp\\Bigl[ - {r_g \\over 2 r_d} (1 - r_d/r)^2 \\Bigr] \\, , \\eqno({\\rm A}1) $$ where $r_g$ is the critical radius (given by eq. [\\ref{eq:rcrit}]) with $T=T_s$. Assuming that $r_s, r_g \\gg r_d$, we thus obtain $$ n_s \\approx n_d \\exp[-r_g/2r_d] \\, . \\eqno({\\rm A}2) $$ The mass loss rate from the disk edge (at $r_d$) is given by the continuity equation and takes the form $$ {\\dot M} = \\muc n_s a_s {\\cal A}_s \\, , \\eqno({\\rm A}3) $$ where $\\muc$ is the mass per particle and where ${\\cal A}_s$ is the area subtended by the flow at $r_s$. This area can be written $$ {\\cal A}_s = 2 \\pi r_d H_d (r_s/r_d)^2 \\equiv 2 \\pi \\alpha r_g (r_d r_g)^{1/2} \\, . \\eqno({\\rm A}4) $$ In the second equality, we have evaluated the disk scale height $H_d$ = $r_g (r_d/r_g)^{3/2}$ and have defined a dimensionless constant $\\alpha \\equiv (r_s/r_g)^2$, which is of order unity. Finally, we apply the condition that the external FUV flux $G_0$ heats a column density $N_C$ given by the integral $$ N_C = \\int_{r_d}^\\infty n(r) dr \\, . \\eqno({\\rm A}5) $$ In the limit that $r_d \\ll r_g$, most of the support of this integral occurs for small $r$ where equation (A1) applies. This condition (A5) relates the column density $N_C$ to $n_d$, and, to leading order, this relation takes the form $$ n_d \\approx \\Bigl( {2 \\over \\pi} \\Bigr)^{1/2} \\Bigl( {r_g \\over r_d} \\Bigr)^{1/2} {N_C \\over r_d} \\, . \\eqno({\\rm A}6) $$ Collecting all of the results given above, we obtain the following expression for the mass loss rate $$ {\\dot M} = C_0 N_C \\muc a_s r_g \\Bigl( {r_g \\over r_d} \\Bigr) {\\rm e}^{-r_g/2r_d} \\, , \\eqno({\\rm A}7) $$ where $C_0$ is a dimensionless constant of order unity.  Although the derivation of equation (A7) applies only in the limit $r_d \\ll r_s \\sim r_g$, the resulting function can be evaluated when $r_d \\approx r_s \\approx r_g$ and implies nearly the same result as the supercritical mass loss rate of \\S 3 (see eq. [\\ref{eq:mdotzero}]). Therefore, we can use equation (A7) as an analytic approximation to the mass loss rate (for a given radiation field $G_0$) as a function of $r_d/r_g$ (for $r_d/r_g \\le 1$).  This approximation should match onto the subcritical mass loss rates calculated in \\S 4 (where $r_d/r_g \\approx 0.125$) and should also match onto the supercritical mass loss rates of \\S 3 (for $r_d \\to r_g$). Notice that we are implicitly assuming that $a_s$ and $N_C$ do not change with $r_d/r_g$ for a given radiation field.  With these approximations, we can estimate the mass loss rates for systems that are intermediate between the subcritical regime of \\S 4 and the supercritical regime of \\S 3. We can also use the resulting form of $\\dot M$ to understand how the mass loss rate depends on the various parameters in the problem. As $r_d$ becomes comparable to $r_g$, the mass loss rate approach its supercritical value. As $r_d/r_g$ decreases, the mass loss rate decreases, but only slowly at first. The outflow rate $\\dot M$ is half its supercritical value when $r_d/r_g \\approx 0.17$ (significantly below unity). For even smaller values of $r_d/r_g$, however, the decaying exponential behavior wins and the mass loss rates drop dramatically.  Finally, we can also make an analytic estimate for the mass loss rate from the disk surface, i.e., for vertical flow off the top and bottom of the disk. This estimate can be compared to that for mass loss from the disk edges (see eq. [A7]). For vertical flow, we treat each increment of disk surface area $2 \\pi r dr$ with the same formulation used above for the disk edges, with one exception: We must replace the radius $r_d$ with $r \\le r_d$ and then integrate over $r$.  This procedure takes into account the fact that material at $r < r_d$ lives deeper in the gravitational potential well and is harder to extract from the system. The resulting mass loss rate from the disk surface is $$ {\\dot M}_{sur} = C_1 N_C \\muc a_s r_g \\Bigl( {r_g \\over r_d} \\Bigr)^{1/2} {\\rm e}^{-r_g/2r_d} \\, , \\eqno({\\rm A}8) $$ where all of the dimensionless quantities are collected into the constant $C_1$ (which is comparable to, but not quite the same as, the constant $C_0$ appearing in eq. [A7]).  Comparing the mass loss rates from the disk edge and the disk surface, we find that ${\\dot M}_{sur}/{\\dot M} \\approx (r_d/r_g)^{1/2}$. In the limit $r_d/r_g \\ll 1$, the mass loss rate from the disk edge dominates the mass loss rate from the disk surfaces.  \\newpage \\bigskip"
    },
    "0404/astro-ph0404397_arXiv.txt": {
        "abstract": "We present a model for the distribution of void sizes and its evolution within the context of hierarchical scenarios of gravitational structure formation. For a proper description of the hierarchical buildup of the system of voids in the matter distribution, not only the {\\it void-in-void} problem should be taken into account, but also that of the {\\it void-in-cloud} issue. Within the context of the excursion set formulation of an evolving void hierarchy is one involving a {\\it two-barrier} excursion problem, unlike the {\\it one-barrier} problem for the dark halo evolution. This leads to voids having a peaked size distribution at any cosmic epoch, centered on a characteristic void size that evolves self-similarly in time, in distinct contrast to the distribution of virialized halo masses in not having a small-scale cut-off. ",
        "introduction": "\\noindent Hierarchical scenarios of structure formation have been very succesfull in understanding the formation histories of gravitationally bound virialized haloes. Particularly compelling has been the formulation of a formalism in which the collapse and virialization of overdense dark matter halos within the context of hierarchical clustering can be treated on a fully analytical basis. This approach was originally proposed by Press \\& Schechter (1974), which found a particularly useful and versatile formulation and modification in the the {\\it excursion set formalism} (Bond et al.~1991).  It is based on the assumption that for a structure to reach a particular nonlinear evolutionary stage, such as complete gravitational collapse, the sole condition is that its {\\em linearly extrapolated primordial density} should attain a certain value. The canonic example is that of a spherical tophat overdensity collapsing once it reaches the collapse barrier $\\delta_c\\approx 1.69$. The successive contributions to the local density by perturbations on a (mass) resolution scale $S_m$ may be represented in terms of a density perturbation random walk, the cumulative of all density fluctuations at a resolution scale smaller than $S_m$. By identifying the largest scale at which the density passes through the barrier $\\delta_c$ it is possible (1) to infer at any cosmic epoch the mass spectrum of collapsed halos and (2) to reconstruct the merging history of each halo (see Fig.~3, lefthand). \\begin{figure} \\vskip -3.0truecm \\mbox{\\hskip -0.5truecm\\includegraphics[width=14.2cm]{weytorfig1.ps}} \\vskip -1.5truecm \\caption{Illustration of the two essential ``void hierarchy modes'': (top) the {\\em void-in-void} process (top), with a void growing through the merging of two or more subvoids; (bottom) the {\\em void-in-cloud} process: a void demolished through the gravitational collapse of embedding region.} \\vskip -0.2truecm \\label{voidproc} \\end{figure} In this study we demonstrate that also the formation and evolution of foamlike patterns as a result of the gravitational growth of primordial density perturbations is liable to a succesfull description by the excursion set analysis. A slight extension and elaboration on the original formulation enables us to frame an analytical theory explaining how the characteristic observed weblike Megaparsec scale galaxy distribution, and the equivalent frothy spatial matter distribution seen to form in computer simulations of cosmic structure formation, are natural products of a hierarchical process of gravitational clustering. This is accomplished by resorting to a complementary view of clustering evolution in which we focus on the evolution of the {\\em voids} in the Megaparsec galaxy and matter distribution, spatially {\\it th\\'e} dominant component (see e.g. Muller \\& Maulbatsch, these proceedings). \\begin{figure} \\vskip 0.5cm \\mbox{\\hskip 0.8truecm\\includegraphics[width=12.0cm]{weytorfig2a.ps}} \\vskip -2.0cm \\mbox{\\hskip 0.0truecm\\includegraphics[width=6.5cm]{weytorfig2b.ps}} \\vskip -6.5cm \\mbox{\\hskip 7.0truecm\\includegraphics[width=6.5cm]{weytorfig2c.ps}} \\vskip 0.0cm \\caption{Identification of void collapse sites: near the boundary of an expanding voids small voids get squashed and sheared. Zoom-in on central lower region of large void (top image) at 2 different timesteps, showing the void compression process.} \\vskip -1.0truecm \\end{figure} ",
        "conclusions": ""
    },
    "0404/astro-ph0404168_arXiv.txt": {
        "abstract": "We use high-resolution dissipationless simulations of the concordance flat $\\Lambda$CDM model to make predictions for the galaxy--mass correlations and compare them to the recent SDSS weak lensing measurements of \\cite{sheldon_etal04}.  The simulations resolve both isolated galaxy-size host halos and satellite halos (subhalos) within them.  We use a simple scheme based on the matching of the circular velocity function of halos to the galaxy luminosity function and on using the observed density-color correlation of the SDSS galaxies to assign luminosities and colors to the halos.  This allows us to closely match the selection criteria used to define observational samples.  The simulations reproduce the observed galaxy--mass correlation function and the observed dependence of its shape and amplitude on luminosity and color, if a reasonable amount of scatter between galaxy luminosity and circular velocity is assumed. We find that the luminosity dependence of the correlation function is primarily determined by the changing relative contribution of central and satellite galaxies at different luminosities. The color dependence of the galaxy--mass correlations reflects the difference in the typical environments of blue and red galaxies. We compare the cross-bias, $b_x\\equiv b/r$, measured in simulations and observations and find a good agreement at all probed scales.  We show that the galaxy--mass correlation coefficient, $r$, is close to unity on scales $\\gtrsim 1 \\hMpc$. This indicates that the cross bias measured in weak lensing observations should measure the actual bias $b$ of galaxy clustering on these scales.  In agreement with previous studies, we find that the aperture mass-to-light ratio is independent of galaxy color in the range of luminosities probed by observational samples. The aperture mass scales approximately linearly with luminosity at $L_r> 10^{10}h^{-2}\\ \\rm L_{\\odot}$, while at lower luminosities the scaling is shallower: $M_{\\Delta\\Sigma}\\propto L_r^{0.5}$.  We show that most of the luminous galaxies ($M_r<-21$) are near the centers of their halos and their galaxy--mass correlation function at $r\\lesssim 100h^{-1}\\ \\rm kpc$ can therefore be interpreted as the average dark matter density profile of these galaxies. Finally, we find that for galaxies in a given narrow luminosity range, there is a broad and possibly non-gaussian distribution of halo virial masses. Therefore, the average relation between mass and luminosity derived from the weak lensing analyses should be interpreted with caution. ",
        "introduction": "\\label{sec:intro}  Understanding the processes that shape the clustering of dark matter and galaxies is one of the main goals of observational cosmology. Modern large redshift surveys, such as the Two Degree Field Galaxy Redshift Survey ~\\citep[2dFGRS,][]{colless_etal01} and the Sloan Digital Sky Survey ~\\citep[SDSS,][]{york_etal00}, allow measurements of galaxy and galaxy--mass correlations and of their dependence on galaxy properties and environment with unprecedented accuracy.  Concurrently, cosmological $N$-body simulations have developed into a powerful tool for calculating the gravitational clustering of collisionless dark matter in hierarchical cosmologies with well-specified initial conditions. The main obstacle in direct comparisons between models and data is understanding the dependence of theoretical predictions on both the relatively straightforward physics of gravitational clustering and the more complex physics of galaxy formation. The former determines the distribution of dark matter, while the latter affects the relationship between the distribution of galaxies and mass, the ``bias''. The complexity of processes operating during galaxy formation on a very wide range of scales makes it difficult to include them directly in simulations, although efforts in this direction are ongoing \\citep[e.g.,][]{katz_etal99, white_etal01, yoshikawa_etal01, pearce_etal01, berlind_etal03,weinberg_etal04}.  Given the difficulty of large-scale galaxy formation simulations, two simpler approaches have recently been pursued to make theoretical predictions for galaxy properties.  In the first, the galaxies are identified with dark matter halos and subhalos in dissipationless cosmological simulations \\citep[e.g.,][]{colin_etal99, kravtsov_klypin99, neyrinck_etal03,kravtsov_etal04}. In the second hybrid approach, collisionless $N$-body simulations are combined with a semi-analytic treatment of galaxy formation \\citep[e.g.,][]{kauffmann_etal97, governato_etal98, kauffmann_etal99a, kauffmann_etal99b, kolatt_etal99, benson_etal00a, benson_etal00b, somerville_etal01, wechsler_etal01, berlind_etal03}.  Galaxy-galaxy lensing is a relatively new but important observational probe of the relation between galaxies and dark matter. It directly measures the galaxy--mass cross-correlation function around galaxies of different types and environments \\citep[see, e.g.,][]{brainerd_etal96, fischer_etal00, hoekstra_etal01, hoekstra_etal04, mckay_etal01, sheldon_etal01, sheldon_etal04,smith_etal01,wilson_etal01}.  Several theoretical studies of the galaxy--mass correlations in cosmological simulations have been carried out in the last three years \\citep{guzik_seljak01,guzik_seljak02,yang_etal03,weinberg_etal04}. Still, there appear to be several discrepancies between the various model results and observations, which are yet to be properly understood.  \\cite{guzik_seljak01} used intermediate-resolution dark matter only simulations combined with a semi-analytic model to identify galaxies and make predictions for the galaxy--mass correlation function.  They compared these results to the first weak lensing detection from SDSS \\citep{fischer_etal00}, and found the amplitude of the galaxy--mass correlation in their model to be systematically higher.  They attributed this discrepancy to the differences between the luminosity function used in their semi-analytic prescription and the observed luminosity function. \\cite{yang_etal03}, using the same theoretical model (intermediate-resolution dark matter simulations combined with the \\citealt{kauffmann_etal99a} semi-analytic galaxy formation prescription), compared to more extensive observational results from \\cite{mckay_etal01}.  They also found that the galaxy--mass correlation function in the simulations was systematically higher than the observational measurements, by about a factor of two. Consequently, they found mass-to-light ratios about a factor of two higher than those observed.  The study of \\cite{weinberg_etal04} investigated several of the same statistics, based on galaxies identified in SPH simulations.  They found that for systems with small baryonic masses, their dark matter-to-baryon ratios agreed with the mass-to-light ratios, derived from weak lensing data, presented by \\cite{mckay_etal01}.  For large masses, however, the baryonic masses of their simulated galaxies were high by a factor of 1.5-2 compared to what is required for agreement with the SDSS data.  In each of these studies, discrepancies with the data may have been due to discrepancies of the model luminosity functions with that observed. \\cite{guzik_seljak02} used a more phenomenological approach, and developed a formalism in the context of the halo model to fit the \\cite{mckay_etal01} data. They extracted information on, e.g., the dependence of the virial mass-to-light ratio on luminosity, the typical mass for galaxies, and the fraction of galaxies in groups and clusters.  The most recent observational study of galaxy--galaxy lensing by the SDSS collaboration \\citep[][hereafter S04]{sheldon_etal04}, has significantly improved the accuracy of galaxy correlation measurements, and may shed light on many of the previous uncertainties.  The spectroscopic sample of 127,001 lensing galaxies was combined with more than $9 \\times 10^{6}$ source galaxies with photometric redshifts, which allowed a detailed study of the relation between mass and light for several luminosity bands, morphological types, and on scales from 20 $\\hkpc$ to 10 $\\hMpc$ (physical). In addition to verifying the findings of previous studies, S04 identified new features in the data, including a scale-dependent luminosity (cross) bias.  Besides being a general test of the Cold Dark Matter paradigm, comparisons with cosmological simulations can be used to gain a deeper insight into the interpretation of these new observational results. Furthermore, they can shed light on to what extent weak lensing results can be used to estimate halo masses, or to learn about the average dark matter halo density profiles. In this paper we use high-resolution dissipationless simulations of the concordance $\\Lambda$CDM model to study the galaxy--mass correlation function with specific emphasis on comparing to the observational measurements of S04. Results of several recent studies suggest that gravitational dynamics is the dominant mechanism shaping galaxy clustering, at least in the simple case of galaxies selected above a luminosity or mass threshold \\citep{kravtsov_klypin99, kravtsov_etal04,zentner_etal04}. \\citet{kravtsov_etal04}, using dark-matter only simulations which resolve galactic mass subhalos, matched galaxies of a given luminosity to a population of halos and subhalos of a given circular velocity and the same number density, and found excellent agreement with the galaxy-galaxy correlation functions measured in the SDSS. This study extends that approach to investigate the galaxy--mass correlations, and makes detailed comparisons to the new \\citet{sheldon_etal04} results.  The paper is organized as follows. In \\S~\\ref{sec:sim} and \\S~\\ref{sec:haloid} we describe the simulations and the halo identification algorithm. The halo samples used in our analysis are described in \\S~\\ref{sec:samples}. The main results, including a detailed comparison with the most recent SDSS measurements, are presented in \\S~\\ref{sec:results}.  In \\S~\\ref{sec:discussion} and \\S~\\ref{sec:conclusions} we discuss and summarize our results and conclusions.  \\begin{table}[tb] \\label{tab:sim} \\caption{Simulation parameters} \\begin{center} \\small \\begin{tabular}{cccccc} \\tableline\\tableline\\\\ \\multicolumn{1}{c}{Name}& \\multicolumn{1}{c}{$\\sigma_8$}& \\multicolumn{1}{c}{$L_{\\rm box}$}  & \\multicolumn{1}{c}{$N_{\\rm p}$}  & \\multicolumn{1}{c}{$m_{\\rm p}$}  & \\multicolumn{1}{c}{$h_{\\rm peak}$} \\\\ \\multicolumn{1}{c}{}& \\multicolumn{1}{c}{}& \\multicolumn{1}{c}{$h^{-1}\\rm Mpc$} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{$h^{-1}\\rm\\ M_{\\odot}$} & \\multicolumn{1}{c}{$h^{-1}\\rm\\ kpc$} \\\\ \\\\ \\tableline \\\\ $\\Lambda$CDM$_{\\rm 60a}$ & 1.0  & 60  & $256^3$ & $1.07\\times 10^9$ & $1.9$\\\\ $\\Lambda$CDM$_{\\rm 60b}$ & 0.9  & 60  & $256^3$ & $1.07\\times 10^9$ & $1.9$\\\\ $\\Lambda$CDM$_{\\rm 60c}$ & 0.9  & 60  & $256^3$ & $1.07\\times 10^9$ & $1.9$\\\\ $\\Lambda$CDM$_{\\rm 80a}$ & 0.75 & 80  & $512^3$ & $3.16\\times 10^8$ & $1.2$\\\\ $\\Lambda$CDM$_{\\rm 80b}$ & 0.9  & 80  & $512^3$ & $3.16\\times 10^8$ & $1.2$\\\\ $\\Lambda$CDM$_{\\rm 120}$ & 0.9  & 120  & $512^3$ & $1.07\\times 10^9$ & $1.8$\\\\ \\\\ \\tableline \\end{tabular} \\end{center} \\label{tab:simparam} \\end{table} ",
        "conclusions": "\\label{sec:conclusions}  We have presented a detailed comparison of galaxy--mass correlation functions measured in high-resolution cosmological simulations of the concordance $\\Lambda$CDM cosmology with the most recent weak lensing SDSS observations. We found that the simple recipe of assigning a luminosity to a halo of a certain maximum circular velocity by matching the subhalo velocity function in simulations to the observed luminosity function leads to good agreement with the observed galaxy--mass correlation and its dependence on luminosity, if an observationally-motivated amount of scatter is introduced in the $V_{\\rm max}-L_{r}$ relation. Our main results and conclusions can be summarized as follows:  \\begin{itemize} \\item[$\\bullet$] The simulations reproduce the galaxy--mass correlation function measured by \\citet{sheldon_etal04} in SDSS and the observed dependence of its shape and amplitude on luminosity. Interestingly, the agreement for bright samples seems to require some scatter between luminosity and circular velocity.  The amount of scatter required is consistent with the observed scatter.  \\item[$\\bullet$] The galaxy--mass correlation function for central galaxies is a strong function of the galaxy luminosity (halo velocity), while $\\xi_{\\rm gm}$ for satellite galaxies is only weakly sensitive to luminosity.  The luminosity dependence of the correlation function as a whole is thus determined primarily by the increasing contribution of bright central galaxies relative to the satellite galaxies at bright luminosities.  Conversely, the correlation function gets shallower with decreasing luminosity because it is increasingly dominated by the contribution from the satellite (non-central) galaxies in halos at intermediate and large scales.  \\item[$\\bullet$] We use the color-density correlation observed in the SDSS to assign colors to the galactic halos in simulations and compare the simulation results to observations in other SDSS bands, from $u$ to $z$. In each band the agreement between galaxy--mass correlation function in simulations and observations is remarkably good. The simulations also reproduce the observed trend of the galaxy--mass correlation function with the $g-r$ color.  \\item[$\\bullet$] In agreement with previous studies, we find that the aperture mass-to-light ratio, $M_{\\Delta\\Sigma}/L_r$, is independent of galaxy color. For $L_r\\gtrsim 10^{10} h^{-2} L_{\\odot}$, $M_{\\Delta\\Sigma}/L_r$ is approximately independent of luminosity. For both central galaxies and galaxies in low density environments the best fit virial mass $M_{\\rm vir}^{\\Delta \\Sigma}$ correlates tightly with the observed aperture mass. The best fit virial mass lies in between the median and the mean actual virial mass when scatter between $V_{max}$ and luminosity is used.  \\item[$\\bullet$] We compare the cross bias, $b_x\\equiv b/r$, measured in simulations and observations and find a good agreement at all probed scales.  The correlation coefficient, $r$, obtained for a volume-limited sample similar to that analyzed by S04, is close to unity on scales $\\gtrsim 1 \\hMpc$. This indicates that the cross bias measured in weak lensing observations does measure the actual bias $b$ of galaxy clustering on these scales.  \\item[$\\bullet$] We show that for luminous galaxies ($M_r<-21$) the galaxy--mass correlation function at $r\\lesssim 100-200h^{-1}\\ \\rm kpc$ can be interpreted as the average density profile of these galaxies. We also show that the masses obtained by the $\\Delta\\Sigma(R)$ measured in weak lensing observations cannot be interpreted in a straightforward way as the mass for galaxies of a given luminosity. In the presence of scatter between mass and light, the obtained mass-luminosity relations should be quoted and interpreted with caution.  \\end{itemize}"
    },
    "0404/astro-ph0404442_arXiv.txt": {
        "abstract": "During the Core Programme, {\\it INTEGRAL}  has surveyed a large portion of the sky (around 9000 square degrees); although {\\it INTEGRAL} is not optimized for extra-galactic studies its observations have nevertheless given us  the opportunity to explore the sky behind our Galaxy, something which is impossible in some wavebands due to the presence of strong absorption. Preliminary results from this exploration are presented and compared with the pre-launch expectations. In particular, we detail all extragalactic detections obtained so far: 10 active galaxies and one cluster of galaxies. Of these sources, many have previously been studied at energies above 10-20 keV while a few are new hard X-ray discoveries. Since the number of detections is smaller than estimated on the basis of the IBIS/ISGRI sensitivity, it is likely that some of the new ISGRI sources found in this survey are extragalactic objects; a few of these are likely to be AGN and are described in detail. ",
        "introduction": "The so-called \"Zone of Avoidance\" refers to the area contained within $\\pm$ 10-15$^\\circ$ of the disk plane of the Milky Way. Gas and dust obscure starlight within this region and screen nearly all background extragalactic objects from traditional optical-wavelength surveys; in the optical, as much as 20$\\%$ of the extragalactic sky is obscured by the Galaxy.  As a consequence, the Galactic plane historically has been neglected by extragalactic astronomers. Hard X-rays ($\\ge$ 10 keV) are able to penetrate this zone thereby providing a \"window\" that is virtually free of obscuration not only relative to optical wavelengths but also partly in comparison to soft X-ray observations. Unfortunately the hard X-ray band is still poorly explored and the only truly all-sky survey conducted so far dates back to the 1980's (Levine et al. 1984). This pioneering work, made with the {\\it HEAO1}-A4 instrument yielded a catalogue of about 70 sources down to a flux level of typically 1/75 of the Crab (or 2-3 $\\times$ 10$^{-10}$ erg cm$^{-2}$ s$^{-1}$) in the 13-80 keV band. Only 7 extragalactic objects are reported in the A4 survey : none of these objects is within 10$^\\circ$ of the galactic plane and only two (Centaurus A and the Perseus cluster ) are located below 20$^\\circ$ in galactic latitude. Pointed observations by {\\it BeppoSAX}/PDS have unveiled more sources but observations were sometimes limited by the lack of imaging capability of the instrument which is particularly crucial in the galactic plane region. A step forward in the study of the zone of avoidance is possible with the imager on board {\\it INTEGRAL}, which allows detection with a sensitivity up to a few mCrab in the most exposed regions (i.e. the galactic center) and provides an angular resolution of 12' and a point source location accuracy of 2-3'.  Here we present a compendium of the extragalactic results obtained so far within the first year of the {\\it INTEGRAL} Core Programme. \\\\  \\begin{small} \\begin{figure*}[t] \\centering \\includegraphics[width=0.8\\linewidth]{sky.eps} \\caption{Sky coverage so far analysed in the present survey work (-180$^{\\circ}$ $<$l$<$+180$^{\\circ}$, -30$^{\\circ}$$<$b$<$+30$^{\\circ}$). The circles have10$^\\circ$  radius to roughly match the half coded field of view of ISGRI.  \\label{fig:single}} \\end{figure*} \\end{small}   \\begin{small} \\begin{figure*} \\centering \\includegraphics[width=0.8\\linewidth]{figure2.eps} \\caption{A typical 20-100 keV ISGRI image showing 2 Seyfert 2s (ESO103-G035 and NGC6300) just above the Galactic Plane.\\label{fig:single}} \\end{figure*} \\end{small} ",
        "conclusions": "We have presented a compendium of the extragalactic results obtained so far within the first year of the {\\it INTEGRAL} Core Programme. Observations of the galactic plane and center have revealed so far 11 extragalactic objects: 10 are  active galaxies  and one is a cluster of galaxies. Of the AGN sample, 2 objects are of type 1  and 7 are of type 2. This provides a type 2 over type 1 ratio which is in line with optical spectroscopic data and furthermore agrees with the expectations of the unified theory. Furthermore it implies a torus half opening angle of $\\sim$40$^{\\circ}$. Many of the objects reported in this work are known to emit at high energies while a few are new hard X-ray discoveries. In particular we find PKS 1830-211 a low frequency (i.e. MeV) peaked or red blazar and Oph Cluster the first and only cluster so far  reported by {\\it INTEGRAL}. We also argue that a few of the new ISGRI sources discovered during the Core Programme could also be of extragalactic origin and describe two likely cases, IGR J21247+5058 and IGR J18027-1455.    \\begin{small} \\begin{figure} \\centering \\includegraphics[width=0.8\\linewidth]{igrj1802_H.eps} \\caption{H-band of the field (1' $\\times$ 1') containing IGR J18027-1455 showing the extended 2MASS object at the center nearby the radio source NVSS J180247-145451.\\label{fig:single}} \\end{figure} \\end{small}"
    },
    "0404/astro-ph0404504_arXiv.txt": {
        "abstract": "There is overwhelming evidence for the presence of supermassive black holes (SMBHs) in the centers of most nearby galaxies. The mass estimates for these remnant black holes from the stellar kinematics of local galaxies and the quasar phenomenon at high redshifts point to the presence of assembled SMBHs. The accretion history of SMBHs can be reconstructed using observations at high and low redshifts as model constraints. Observations of galaxies and quasars in the submillimeter, infrared, optical, and X-ray wavebands are used as constraints, along with data from the demography of local black holes. Theoretical modeling of the growth of black hole mass with cosmic time has been pursued thus far in two distinct directions: a phenomenological approach that utilizes observations in various wavebands, and a semi-analytic approach that starts with a theoretical framework and a set of assumptions with a view to matching observations. Both techniques have been pursued in the context of the standard paradigm for structure formation in a Cold Dark Matter dominated universe. In this chapter, we examine the key issues and uncertainties in the theoretical understanding of the growth of SMBHs. ",
        "introduction": "The local demography of black holes (Ferrarese, this volume) has established that most galaxies harbor a supermassive black hole (SMBH; Kormendy \\& Richstone 1995; Magorrian et al.\\ 1998; van der Marel 1999), most likely assembled via a combination of accretion and mergers. These nuclear SMBHs are ``dead quasars'', relics of quasar activity that might have occurred in many galaxies over their history (Lynden-Bell 1969; So\\l tan 1982; Rees 1990; Richstone et al.\\ 1998). Early attempts to interlink the properties of these remnant black holes with those of their host galaxy luminosities (Magorrian et al.\\ 1998) yielded a relation between the bulge luminosity and the black hole mass of about 0.5 dex in the ratio of $\\mh/\\mb$. A tighter correlation has since been measured between the velocity dispersion of the bulge and the black hole mass (Ferrarese \\& Merritt 2000; Gebhardt et al.\\ 2000), suggesting that the formation and evolution of SMBHs is inextricably linked to that of the stellar component of galactic bulges.  Quasar activity is powered by gas accretion onto SMBHs (Lynden-Bell 1969), and hence the build-up of these SMBHs is likely to have commenced at fairly high redshifts. Indeed, optically bright quasars have now been detected at redshifts greater than 6 (e.g., Fan et al.\\ 2001a, 2003). Quasars at these high redshifts provide an efficient tool to investigate the relation between black hole and early spheroid assembly. There are indications that high redshift quasar hosts are often strong sources of dust emission (Omont et al.\\ 2001; Cox et al.\\ 2002; Carilli et al.\\ 2002; Walter et al.\\ 2003), suggesting that quasars were common in massive galaxies at a time when the galaxies were undergoing copious star formation.  The growth of black hole mass in the universe can therefore be traced using quasar activity. The phenomenological approach to understanding the assembly of SMBHs involves using observational data from both high and low redshifts as a starting point to construct a viable and consistent picture that is consonant with the larger framework of the growth and evolution of structure in the universe. Another approach that has been pursued is semi-analytic modeling, in which one starts from a set of ab initio assumptions and attempts to explain the observations.  Both approaches have proved to be fruitful and, in fact, share many common features. They are both grounded in the framework of the standard paradigm that involves the growth of structure via gravitational amplification of small perturbations in a Cold Dark Matter (hereafter, CDM) universe---a model that has independent validation, most recently from {\\it Wilkinson Microwave Anisotropy Probe (WMAP)} measurements of the anisotropies in the cosmic microwave background (Spergel et al.\\ 2003; Page et al.\\ 2003). Structure formation in both modeling schemes is tracked in cosmic time by keeping a census of the number of collapsed dark matter halos of a given mass that form; these provide the sites for harboring black holes. The computation of the mass function of dark matter halos is done using either the Press-Schechter (Press \\& Schechter 1974) or the extended Press-Schechter theory (Lacey \\& Cole 1993), or, in some cases, directly from cosmological N-body simulations (Di Matteo et al.\\ 2003). In this chapter, we present the detailed modeling procedures and key parameters and uncertainties for the phenomenological approach, as discussed by Haehnelt, Natarajan, \\& Rees (1998). We also summarize the results from the semi-analytic modeling. ",
        "conclusions": ""
    },
    "0404/astro-ph0404262_arXiv.txt": {
        "abstract": "{  A new method to evaluate the dust-to-gas ratios in the Kepler SNR is presented. Dust emission  in the  infrared and bremsstrahlung are calculated consistently, considering    that    dust grains are collisionally heated by  the gas throughout the  front and downstream of both the expanding and the reverse shocks. The  calculated continuum SED is  constrained by the observational data. The dust-to-gas ratios are determined by the ratio of the dust emission bump and bremsstrahlung in the  infrared. The shell-like morphological similarity  of  X-ray    and radio emission, and of    the \\Ha ~and  infrared images confirms that both radio  and X-ray emissions are created at the front of the  expanding  shock and  that dust and gas are coupled   crossing  the expanding and reverse shock fronts. The results show that large  grains with radius of $\\sim$ 1 \\mum ~with dust-to-gas ratios  $<$ 4 10$^{-3}$ survive sputtering and are  heated to a maximum temperature of 125 K downstream of the shock expanding outwards with  velocity of about 1000 \\kms. The high velocity shocks become radiative for dust-to-gas ratios $>$ 10$^{-3}$. Such shocks do not appear in the NE region, indicating that dust grains are not homogeneously distributed throughout the remnant. Smaller grains with radius of about  0.2\\mum ~and   dust-to-gas ratios of  $\\sim$ 4 10$^{-4}$ are heated to a maximum temperature of $\\sim$ 50 K downstream of the reverse shock corresponding to velocities of about 50 \\kms. A maximum dust mass   $<$  0.16 \\msol is calculated. ",
        "introduction": "Since   the first identification by Baade (1943) the remnant of the Kepler  supernova (SN)  which exploded in 1604, has  been  observed in the different wavelength ranges.  van den Bergh \\& Kemper (1977) studied brightness variations and proper motions of the filaments in the supernova remnant (SNR). Spectrophotometry data of the bright-western knots in the spectral range 3700-10,500 \\AA ~were  presented  by Dennefeld (1982). Observations in the optical range for single bright knots were obtained by Leibowitz \\& Danziger (1983) and by Blair, Long \\& Vancura (1991).  Data in the X-ray  (e.g. Hughes 1999) and in the radio (e.g. DeLaney et al. 2002) ranges were also recently presented. Morgan et al (2003) provided  SCUBA data of the continuum in the far infrared  (450 and 850 \\mm), so the information about the infrared (IR) emission is rather complete.  The main results obtained from the interpretation of the line spectra (Blair et al) are, in particular, that the emitting gas has high densities ($\\geq$ 1000 \\cm3), that the N/H relative abundance is $\\sim$ 3.5 times higher than solar, and \\Ha ~line profiles display both broad and narrow components. A distance of  5 kpc  was consistently determined.  The   origin of Kepler SNR emitting features  is explained by the interaction of the high velocity ejecta with the ambient medium (Decourchelle \\& Pretre 1999) and is very similar to that found for other SNRs , e.g. Cassiopeia A (Chevalier \\& Oishi 2003), namely, \"that the observed shock wave positions and expansions can be interpreted in a model of supernova interaction with a freely expanding stellar wind.\" The  collision  with ambient matter leads to the formation of two shock fronts : a shock propagating outwards throughout the circumstellar/interstellar medium  and one propagating  in reverse throughout the ejecta in the direction opposite to the propagation of the blast wave. The gas in the knotty regions  of Kepler SNR is ionized and heated to relatively high temperatures ($\\leq$ 3 10$^7$ K)  downstream of the shock which expands  with velocities  of 1000 -1500 \\kms. The reverse  shock, on the other hand, propagating with velocities of $\\sim$ 50 \\kms in the opposite direction, leads to lower gas temperatures. The velocities are reduced by the  density gradient of the ejected matter. The circumstellar origin of the optical emitting gas  suggested that Kepler's SN was actually of Type Ib (Bandiera 1987).  Very recently Morgan et al (2003) addressed  the problem of dust in Kepler SNR in terms of  its total mass and its origin from a massive star. They claim that dust formation in SNe   is an important process,  but dust in SNRs has been detected in small quantities.  In this work we will use the method adopted by e.g. Contini \\& Contini (2003) to calculate the dust-to-gas ratios in the starburst regions of luminous infrared galaxies  and by Contini, Viegas \\& Prieto (2004) for Seyfert galaxies. Namely, dust-to-gas ratios  are  calculated by  modeling the spectral energy distribution (SED) of the continuum in  the optical-IR frequency range.  At the shock front edge of the emitting nebulae, the gas is collisionally heated to relatively high temperatures which depend on the shock velocity. Dust and gas  are coupled crossing the shock front, mutually heating and cooling by collisions (Dwek \\& Arendt 1992, Viegas \\& Contini 1994, Contini et al. 2004). The temperature of the grains  depends  on the  temperature of the gas and, therefore, the maximum temperature of dust depends on the shock velocity (Contini et al. 2004). The observational data constrain the model which better  explains both the bremsstrahlung and reradiation by dust. Particularly, the ratio between  the dust radiation flux  in the IR and bremsstrahlung depends  on the dust-to-gas ratio.  However,  the modeling of the continuum SED must be  cross-checked by the modeling of the line spectra, which is strongly constraining.  In this paper I would like to  find out the physical conditions in the different  regions of Kepler SNR, as well as  the dust-to-gas ratios, by consistent model calculations,   which take into account the results of previous investigations.  The SUMA code (Viegas \\& Contini 1994, Contini \\& Viegas 2001 and references therein), which is adapted  to the calculation of gas and dust spectra emitted downstream of a shock front, is  adopted.  The  general models are presented in Sect. 2. The models specifically calculated for the Kepler SNR appear  in Sect. 3.1. The comparison of  the spectra with the data  is discussed in Sect. 3.2  and cross checked by the fit of the continuum SED in Sect. 3.3. Discussion and concluding remarks  follow in Sect. 4. ",
        "conclusions": "In this work the calculation of dust emission  is carried out consistently with emission from gas for  both  line and continuum spectra, leading to the evaluation of the dust-to-gas ratios in   the Kepler SNR. A successful application of this method (Contini et al 2004)  was presented for luminous IR galaxies (Contini \\& Contini 2003) and for AGN (e.g. Contini \\& Viegas  2000).  The  two  component emission  observed in  the FWHM of the   \\Ha ~line profiles corresponding to velocities of $\\sim$ 50 \\kms and $\\sim$ 1000 \\kms is  confirmed by modeling. They correspond to the reverse and the expanding shocks, respectively,   and are explained by the supernova interaction with a stellar wind.  A grid of single-cloud models is calculated. The models  cover the set  of physical conditions in the emitting gas that are revealed  by the line ratios and by the FWHM of the line profiles. Multi-cloud models which account for the reverse and expanding shock in each region, are adopted in order to fit the line spectra in the different observed regions and the corresponding continuum SED. The total continuum  which results from summing up the contributions  of the most luminous regions, is compared with the observed  continuum SED. Dust emission  in the IR is consistently calculated. The ratio between    dust emission  and   bremsstrahlung in the IR range depends on the dust-to-gas ratios adopted in the models.  The  high dust-to gas ratios  which are found  by modeling the IR  bump in  models corresponding to high velocity shocks, speed up the cooling rate downstream, leading to radiative  shocks even   in  a relatively low density gas. It is found that radiative high velocity shocks are absent in the NE region of the remnant, suggesting that  d/g is low and/or the grains are small enough to be destroyed by sputtering. This result may indicate that dust  is not homogeneously distributed  throughout the remnant.  Moreover, Kepler SNR  morphological structure in the different wavelengths can be   explained by model  results. In fact, the X-ray emission has the same shell-like morphology and is qualitatively similar to the radio emission (Blair et al, DeLaney et al.), confirming that both radio  and X-ray are created at the shock front of the expanding shock. Indeed, the   high temperatures ($>$ 10$^{7}$ K)  of the gas in the immediate post-shock region downstream, correspond to the X-ray emission.  The \\Ha ~and IR images are  also similar (Blair et al.) and are similar to the X-ray image  as well. Particularly,  the striking similarity between the \\Ha ~and IR images suggests that the 12 \\mm ~thermal dust emission  and  the optical emission  have the same origin (DeLaney et al. 2002)  confirming that    dust  grains  are collisionally heated by the gas across the shock fronts and downstream. The  regions of dissimilarity  could be explained by the inhomogeneous distribution of  dust grains.  Notice that the modeling of the Kepler SNR through the line and continuum spectrum analysis in the different regions is presented in previous sections  with the aim of  calculating the dust-to-gas ratios. Actually, this work   has been inspired   by Morgan et al (2003) who claim that dust formation in supernovae is required to be an important process relative to the age of the Universe.  Morgan et al calculate that the maximum  dust  mass swept up by the SNR is $\\sim$ 10$^{-3}$ \\msol, assuming a  maximum density  n = 0.1 \\cm3 and a gas to dust ratio of 160. However, adopting a  higher n (100 \\cm3) the dust mass swept up is $\\sim$ 1 \\msol. Douvion et al (2001) claim that by using  dust grains of 'astronomical silicates' both mid-IR  ISOCAM and IRAS data can be fitted by a single grain temperature of 107 K, while previous models required two dust components : a hot dust of 140 K and a cold one of 54 K (Saken et al. 1992). The total mass calculated by them is 10$^{-4}$ \\msol.  The  model presented in this work  shows that  two  main  types of dust  lead to a good fit of the IR bump consistently with the bremsstrahlung. Large  grains with \\agr=1 \\mum ~ with d/g $<$ 4 10$^{-3}$ survive sputtering downstream of the expanding shock and  are collisionally heated to a maximum temperature of 125 K. Smaller grains (\\agr=0.2\\mum)  downstream of the reverse shock (Fig. 1) with   dust-to-gas ratios of  $\\sim$ 4 10$^{-4}$ are heated to a maximum temperature of $\\sim$ 50 K.  Adopting a SNR  diameter of 3.8 pc, a pre-shock density $\\sim$ 50 \\cm3  by the average of models m7, m8, m9 (Table 1) (Sect. 3.3), we obtain a swept up gas mass of  40 $\\it ff$ \\msol and a maximum dust mass of   0.16 $\\it ff$ \\msol, where $\\it ff \\leq 1$ is the filling factor.  Such  relatively high  pre-shock densities in single clouds are  justified  by the relatively young age of Kepler SNR. In older SNR  (e.g. the  Cygnus Loop), \\n0 is less than 10 \\cm3 (Contini et al 1980). Moreover, the clumpy aspect of the remnant suggests that average densities on large scales are not realistic and that $\\it ff$ should be $\\leq$ 0.1.  Comparing with  other SNR, Douvion et al (2001a) found that a dust mass of 10$^{-4}$ \\msol  results from the warm component (140 K as for Kepler SNR) and of 4 10$^{-3}$ \\msol for the cold component (55 K) of Tycho. They claim that the Crab nebula IR emission is dominated by synchrotron radiation and no dust is detected. In Cassiopeia A the dust is  made  of more  components,  e.g. quartz and aluminum oxid, besides silicates which have a mass of 7.8 10$^{-9}$ \\msol (Douvion et al 2001b)."
    },
    "0404/astro-ph0404054_arXiv.txt": {
        "abstract": "We present a new infrared extinction study of Globule 2, the most opaque molecular cloud core in the Coalsack complex.  Using deep near-infrared imaging observations obtained with the ESO NTT we are able to examine the structure of the globule in significantly greater detail than previously possible.  We find the most prominent structural feature of this globule to be a strong central ring of dust column density which was not evident in lower resolution studies of this cloud. This ring represents a region of high density and pressure that is likely a transient structure.  For a spherical cloud geometry the ring would correspond to a dense inner shell of high pressure that could not be in dynamical equilibrium with its surroundings since there appear to be no sources of pressure in the central regions of the cloud that could support the shell against gravity and prevent its inward implosion. The timescale for the inward collapse of the ring would be less than 2 x 10$^5$ years, suggesting that this globule is in an extremely early stage of evolution, and perhaps caught in the process of forming a centrally condensed dense core or Bok globule.  Outside its central regions the globule displays a well-behaved density profile whose shape is very similar to that of a stable Bonnor-Ebert sphere.  Using SEST we also obtained a \\ceio spectrum toward the center of the cloud.  The CO observation indicates that the globule is a gravitationally bound object.  Analysis of the CO line profile reveals significant non-thermal gas motions likely due to turbulence.  As a whole the globule may be evolving to a global state of quasi-static dynamical equilibrium in which thermal and turbulent pressure balance gravity. ",
        "introduction": "Low mass dense cores found within molecular cloud complexes and isolated dark globules (also known as Bok globules) are the simplest configurations of dense molecular gas and dust known to form stars (e.g, Benson \\& Myers 1989; Yun \\& Clemens 1990).  They have been long recognized as important laboratories for investigating the physical processes which lead to the formation of stars and planets (e.g., Bok 1948).  In recent years deep infrared imaging surveys of such globules have provided an important new tool for detailed examination of their structure.  For example, infrared observations of the starless globule B 68 produced an exquisitely detailed measurement of the cloud's radial density distribution (Alves, Lada \\& Lada 2001, hereafter ALL01).  This density distribution was shown to match very closely that predicted for a marginally stable, pressure bounded, isothermal sphere in hydrostatic equilibrium (i.e., a Bonnor-Ebert sphere).  Similar deep near-infrared observations of the darker globule B 335 enabled a detailed investigation of the radial density profile of a globule containing a central protostar.  The profile in this star forming globule was found to be significantly steeper than that of B 68, consistent with expectations for a collapsed or collapsing core (Harvey et al. 2001).  These observations suggest that the evolution of dense cores to form stars can be effectively traced and measured by detailed observations of the cloud's density structure, best obtained by infrared extinction observations.  Two of the least understood aspects of the star formation process are the initial conditions that describe dense cores that ultimately form stars and the origin of such dense cores from more diffuse atomic and molecular material.  Deep infrared and millimeter-wave observations of starless cores or globules offer the best opportunity to investigate these issues.  Starless cores account for about 70\\% of optically selected dark cores and globules (Lee \\& Myers 1999).  A particularly interesting group of cores in this regard are those associated with the conspicuous Coalsack dark cloud complex in the southern Milky Way (Tapia 1973; Bok, Sim \\& Hawarden 1977; Bok 1977).  The entire Coalsack complex subtends an angle of roughly 6$^o$ on the sky corresponding to a linear dimension of nearly 15 pc at the distance of 150 pc estimated for this cloud (e.g., Cambresy 1999, and references therein).  A survey of $^{12}$CO emission found the cloud to be characterized by complex, filamentary structure and to be relatively massive containing about 3500 \\msunsp of material (Nyman, Bronfman \\& Thaddeus 1989).  However, a survey of emission from the rarer $^{13}$CO isotopic line found the ratio of $^{13}$CO to $^{12}$CO emitting gas in the cloud to be considerably lower (17\\%) than that (50-80\\%) which characterizes nearby star forming molecular clouds (Kato et al.  1999). This observation indicated that the fraction of dense gas in the Coalsack complex is considerably smaller than that which characterizes typical star forming cloud complexes. The paucity of dense gas coupled with the lack of the usual signposts of star formation activity, such as emission-line stars, HH objects, embedded infrared sources, etc. (e.g., Kato et al.  1999) suggests that the Coalsack may be a molecular cloud complex in the earliest stages of evolution.  The globules within the Coalsack are apparently all starless and also may be in the early phases of development.  Knowledge of the detailed structure of globules in the Coalsack may provide further insight concerning their physical nature and evolutionary status.  In particular, it would be interesting to know to what extent the structure of a Coalsack globule resembles that of a starless cloud like B 68 or a star forming globule such as B 335. Tapia's Globule~2 is the most prominent and likely densest globule in the Coalsack (Bok et al.  1977; Kato et al.  1999).  Located slightly below ($\\sim$ 1$^o$) the galactic plane, it is projected against a rich star field and is a prime candidate for infrared extinction studies.  Jones et al.  (1980) obtained JHK photometry of 75 stars located behind the Coalsack cloud and derived individual extinctions to all these stars.  Their observations suggested a radial column density distribution that was not highly centrally concentrated.  Racca, Gomez and Kenyon (2002) obtained JHK infrared images of the globule that were more sensitive and resulted in the detection of a few thousand stars behind the cloud.  They used H-band star counts to construct a low resolution extinction map of the globule and derived an azimuthally averaged radial density profile which confirmed the shallow nature of the radial column density distribution.  Moreover, they found that the radial density distribution could be fit by a Bonnor-Ebert configuration with a density profile quite similar to that of B 68 and considerably shallower than that of B 335.  In this paper we report new and deeper near-infrared imaging observations of Globule~2 (hereafter G2).  We use these observations to determine the individual line-of-sight extinctions to thousands of stars behind the cloud.  These measurements enable an examination of the structure of the globule to be made in a degree of detail considerably greater than that obtained by previous studies and comparable to that acheived earlier for B 68.  The high angular resolution achieved by our observations reveal the globule to be more structured than suspected previously.  Moreover, these observations indicate that the spatial distribution of column density in this globule differs significantly from that of other well studied examples, such as B 68 and B 335. In particular, the radial dust column density profile of G2 is not only found to be shallow, but also to be characterized by a significant central depression that, in turn, is surrounded by a prominent ring of high column density.  These characteristics suggest that Globule~2 may represent a very early stage in the evolution of cloudy material to form a centrally condensed dense core. ",
        "conclusions": "A number of considerations from the analysis of our observations of Tapia's Globule~2 in the Coalsack suggest that the globule is a bound core in a state very close to overall dynamical equilibrium.  Perhaps the most compelling of these is the finding that the measured gas motions in the globule, derived from the linewidth of \\ceio emission, indicate that the cloud is gravitationally bound and may even be near virial equilibrium.  Moreover, the density gradient of the cloud, beyond the central ring, is consistent with that of a sub-critical BE cloud in which thermal and non-thermal pressure balance self-gravity.  This, in turn, is consistent with the fact that the cloud's mass is found to be less than the Jean's mass.  In addition the globule's structure is well resolved by our extinction observations and appears to be relatively smooth, though less smooth than B 68, the only other cloud studied in such detail by infrared extinction measurements.  However, the G2 core differs from typical bound cloud cores in some significant ways. The peak extinction (A$_V$ $\\approx$ 12 mag.)  is relatively low for a dense core.  Most importantly, the density structure of the central regions of the G2 cloud departs significantly from that of a typical cloud core and from the predictions for a centrally condensed BE configuration in hydrostatic equilibrium.  In particular, the observed column density distribution is characterized by a deep central depression bordered by a well-defined ring of high column density and high pressure.  The central ring is likely out of equilibrium with the gas in the very center of the cloud and consequently appears to be a transient feature with an evolutionary time scale of order 10$^5$ years which is less than the expected million year lifetime of a starless core.  This would seem to indicate that the central regions of the cloud are in a phase of structural transition and evolving to a new physical configuration.  Given that the central column density depression represents a mass deficit between the observed and best-fit BE density profile of only about 1-2\\% the total cloud mass, it seems plausible to assume that the cloud will evolve to a more relaxed, BE-like, configuration. The gradual dissipation of turbulence in the central regions of the cloud may be driving this transition. Thus, G2 appears to be a cloud core in an early stage of formation.   The Coalsack complex has long been known to be unusual among molecular clouds in its lack of any significant star forming activity (Nyman et al.  1989).  Earlier CO observations have shown that it is also unusual in its relative paucity of dense gas (Kato et al.  1999).  These facts are clearly not unrelated since surveys of molecular clouds have demonstrated that stars form exclusively in dense gas (e.g., Lada 1992). The Coalsack complex is plausibly a very young molecular complex.  The fact that we find the structure of the most opaque globule in the Coalsack cloud to clearly differ from the steep centrally condensed configurations of star forming cores and globules in other clouds, appears consistent with the idea that this core is also quite young.  Exactly how early this core is in its development is difficult to determine.  However, many starless cores, for example, B68 (ALL01) and L694-2 (Harvey et al 2004), are known to be characterized by a high degree of central concentration, often similar to star forming cores.  If such structure is more typical of starless cores then we would suspect that G2 represents a rare example of an extremely early phase of dense core evolution. Investigation of the detailed structures of a larger sample of starless cores is clearly needed to better quantify this conjecture.  However, we note that although starless cores in molecular clouds are relatively common (Lee \\& Myers 1999), cloud complexes without star formation are extremely rare.  Together these factors suggest the extreme youth of the Coalsack complex and indicate that this region constitutes an important laboratory for future studies of the poorly understood processes of cloud and core formation."
    },
    "0404/astro-ph0404324_arXiv.txt": {
        "abstract": "{We briefly discuss a handful of topics in pulsar astrophysics, first some general well-known features, then an overview of the glitch phenomenon and the sort of information gathered about the internal structure and dynamics, and finally the quandary posed by the precession of PSR B1828-11 a very important clue pointing towards a novel paradigm for structure of the core regions. We point out that ``exotic\" solutions for the precession puzzle would force a consideration of exotic {\\it glitch} mechanisms as well. } ",
        "introduction": "Pulsars are now a ``classical\" subject of modern astrophysics after almost four decades of intense work. Shortly after their discovery in 1967 \\cite{HB}, beautiful theoretical work \\cite{Pac,Gold} convincingly argued that neutron stars (and not, for example, white dwarfs) were responsible for the emission. As a brand new field at that time, several ideas were put forward and contributed to fundament the broad-brush picture of pulsars available today. Thus, concepts such as the charged magnetosphere, light cylinder and so on form now (in spite of the lack of exact solutions for this complicated plasma problem), a body of concepts subject to continuous testing (see Ref.4 for a comprehensive discussion.  This is not the appropriate place to recall the spectacular advances in high-resolution instrumentation (see this volume), but the availability of enhanced ground (Keck, Arecibo, etc.) and space (HST, Chandra, XMM, etc.) facilities, coupled with the intensive long-term monitoring (radio) of a handful of objects and targeted searches have now revealed a wealth of phenomena not always fitting into the ``standard\" view. This has in turn enriched our vision of pulsars, and also created puzzles for the models which are being worked out, as is the case of the glitch phenomenon in particular. We may say that pulsar physics is definitely entering its maturity where more detailed models can be constructed and tested. ",
        "conclusions": "We have given a broad overview of one of the most spectacular dynamical features of pulsars (glitches), repeatedly associated with perhaps the most gigantic quantum fluids found in nature, namely the components of the crusts of neutron stars. Accurate timing for over three decades have provided very good data still waiting for a comprehensive explanation. Strong limits to the idea of heat release in a glitch have been set recently by Helfand et al. \\cite{Helf} and Pavlov et al. \\cite{tocayo} using Chandra data to show that the temperature of the pulsar did not change by more than $0.2 \\%$ one month after the Vela glitch of January 2000. Even though it is not impossible that some mechanism can get rid of the heat very quickly, these observations constrain the thermal models in which a large energy input is needed. An even more serious challenge has been posed by some authors (notably Jones , \\cite{PBJ} and references therein) suggesting that vortices in the crust do {\\it not} pin at all (see also Donati and Pizzochero \\cite{Pizz} for a general analysis). May be the core plays a role \\cite{JM}, but as discussed above, this component surely hides some (big ?) surprises and would require extensive studies. We are far from a thorough understanding of the body of evidence of glitches, whereas additional complications from related observations have enriched the general picture recently. A whole new synthesis is needed soon, and perhaps a change in the paradigm as well to pin down the essentials of pulsar dynamics."
    },
    "0404/astro-ph0404112_arXiv.txt": {
        "abstract": "We consider the degree to which ``21 cm tomography'' of the high-redshift Universe can distinguish different ionization histories. Using a new analytic model for the size distribution of \\ion{H}{2} regions that associates these ionized bubbles with large-scale galaxy overdensities, we compute the angular power spectrum and other statistical properties of the 21 cm brightness temperature during reionization.  We show that the \\ion{H}{2} regions imprint features on the power spectrum that allow us to separate histories with discrete bubbles from those with partial uniform ionization (by, for example, X-rays).  We also show that ``double reionization'' scenarios will modify the morphology of the bubbles in ways that depend on the mechanism through which the first generation of sources shuts off. If, for example, the transition occurs globally at a fixed redshift, the first generation imprints a persistent feature on the 21 cm power spectrum.  Finally, we compare our model to one in which voids are ionized first.  While the power spectra of these two models are qualitatively similar, we show that the underlying distributions of neutral hydrogen differ dramatically and suggest that other statistical tests can distinguish them.  The next generation of low-frequency radio telescopes will have the sensitivity to distinguish all of these models and strongly constrain the history and morphology of reionization. ",
        "introduction": "\\label{intro}  Recently, a great deal of effort -- both observational and theoretical -- has been focused on understanding the reionization of hydrogen in the intergalactic medium (IGM) at $z \\ga 6$.  Several independent observational techniques offer complementary views of this landmark event.  The most straightforward method is to seek extended troughs of complete Ly$\\alpha$ absorption in the spectra of high-redshift quasars.  Evidence for this \\citet{gp} effect has been found near $z \\sim 6.5$ \\citep{becker,white03}. Unfortunately, the optical depth of a fully neutral medium is so high that current measurements only require a mean neutral fraction $\\bxh \\ga 10^{-3}$, and even this limit has been disputed \\citep{songaila04}.  Studies of the rapidly growing ionizing background \\citep{fan} and the Str{\\\" o}mgren spheres surrounding the quasars \\citep{wyithe04-prox} suggest larger neutral fractions but depend on detailed modeling.  A second constraint arises because free electrons scatter cosmic microwave background (CMB) photons, washing out intrinsic anisotropies and generating a polarization signal \\citep{zal97}. This provides an integral constraint on the reionization history; recent measurements of the CMB imply that $\\bxh$ must have been small at $z \\ga 14$ \\citep{kogut03,spergel03}.  Finally, the relatively high temperature of the Ly$\\alpha$ forest at $z \\sim 2$--$4$ suggests an order unity change in $\\bxh$ at $z \\la 10$ \\citep{theuns02-reion,hui03}, although this argument depends on the characteristics of \\ion{He}{2} reionization (e.g., \\citealt{sokasian02}).  Taken together, these three sets of observations suggest a complex ionization history extending over a large redshift interval ($\\Delta z \\sim 10$).  This is inconsistent with a generic picture of fast reionization (e.g., \\citealt{barkana01}, and references therein).  The results seem to indicate strong evolution in the sources responsible for reionization, and a precise measurement of the ionization history would strongly constrain early structure formation \\citep{wyithe03,cen03,haiman03,sokasian03b}.  Unfortunately, extracting such detailed information requires new observational techniques.  One of the most exciting possibilities is ``21 cm tomography'' of the high-redshift IGM \\citep{scott,kumar,mmr}, in which one maps the distribution of neutral hydrogen on large scales through its hyperfine transition.  By probing a specific spectral line, 21 cm surveys measure the time history of reionization, and unlike the Ly$\\alpha$ forest they do not suffer from saturation problems.  In fact, the signal is sufficiently weak that foregrounds from the Galaxy and a wide variety of extragalactic objects pose substantial challenges to these experiments \\citep{oh03,dimatteo04}. Fortunately, the known foregrounds all have smooth continuum spectra, which should allow frequency cleaning to high accuracy because the 21 cm signal has spectral structure on small scales (\\citealt{zald04}, hereafter ZFH04; \\citealt{morales03,cooray04}).  However, predicting the 21 cm signals from reionization has proven difficult.  One method is to use numerical simulations of reionization \\citep{ciardi03,furl-21cmsim,gnedin03}, but computational costs have so far limited the simulations to subtend (at best) a few resolution elements of the 21 cm maps.  Analytic models of reionization can extend to larger scales but require some way to describe the complexities of structure formation and radiative transfer.  The simplest approach, in which we assign each galaxy its own \\ion{H}{2} region, does a poor job of matching the large ionized bubbles found in simulations \\citep{ciardi03-sim,sokasian03}.  We have recently developed a model that associates \\ion{H}{2} regions with large-scale fluctuations in the density field and reproduces the major features seen in simulations (\\citealt{furl04a}, hereafter Paper I).  This allows us, for the first time, to make quantitative predictions about the 21 cm signal at reionization.  With high signal-to-noise 21 cm maps, we can directly measure the distribution of \\ion{H}{2} regions in order to test our model.  However, the signals are sufficiently weak that such maps will be difficult to make; fortunately, statistical measurements of the data also contain a great deal of information about reionization (ZFH04). In this paper, we use the model of Paper I to predict the 21 cm angular power spectrum for several reionization scenarios, with an emphasis on how the measurements help distinguish the crucial features of these different pictures.  In \\S \\ref{model} we briefly review our model for reionization.  We then show how 21 cm measurements can distinguish different reionization histories in \\S \\ref{complex} and how they can distinguish different models of reionization in \\S \\ref{void}.  We conclude in \\S \\ref{disc}.  In our numerical calculations, we assume a $\\Lambda$CDM cosmology with $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$, $\\Omega_b=0.046$, $H=100 h \\hunits$ (with $h=0.7$), $n=1$, and $\\sigma_8=0.9$, consistent with the most recent measurements \\citep{spergel03}. ",
        "conclusions": "21 cm measurements contain an unprecedented wealth of information about the timing and morphology of reionization.  Our simple model shows that the qualitative statistical features of the signal can be related to the source and IGM properties.  The next generation of low-frequency radio telescopes, such as the Primeval Structure Telescope,\\footnote{ See http://astrophysics.phys.cmu.edu/$\\sim$jbp for details on PAST.} the Low Frequency Array,\\footnote{ See http://www.lofar.org for details on LOFAR.} and the Square Kilometer Array,\\footnote{ See http://www.skatelescope.org for details on the SKA.} should be able to measure these statistical properties (ZFH04) and place strong constraints on the reionization era.  This work was supported in part by NSF grants ACI AST 99-00877, AST 00-71019, AST 0098606, and PHY 0116590 and NASA ATP grants NAG5-12140 and NAG5-13292 and by the David and Lucille Packard Foundation Fellowship for Science and Engineering."
    },
    "0404/astro-ph0404438_arXiv.txt": {
        "abstract": "The spectral energy distributions (SEDs) of dust-enshrouded galaxies with powerful restframe far-infrared(IR) emission have been constrained by a range of ground-based and space-borne surveys. The {\\it IRAS} catalog provides a reasonably complete picture of the dust emission from nearby galaxies (redshift $z \\simeq 0.1$) that are typically less luminous than about $10^{12}\\,L_\\odot$. However, at higher redshifts, the observational coverage from all existing far-IR and submillimeter surveys is much less complete. Here we investigate the SEDs of a new sample of high-redshift submillimeter-selected galaxies (SMGs), for which redshifts are known, allowing us to estimate reliable luminosities and characteristic dust temperatures. We demonstrate that a wide range of SEDs is present in the population, and that a substantial number of luminous dusty galaxies with hotter dust temperatures could exist at similar redshifts ($z \\simeq 2-3$), but remain undetected in existing submillimeter surveys. These hotter galaxies could be responsible for about a third of the extragalactic IR background radiation at a wavelength of about 100\\,$\\mu$m. The brightest of these galaxies would have far-IR luminosities of the order of $10^{13}\\,L_\\odot$ and dust temperatures of the order of 60\\,K. Galaxies up to an order of magnitude less luminous with similar SEDs will be easy to detect and identify in the deepest {\\it Spitzer Space Telescope} observations of extragalactic fields at 24\\,$\\mu$m. ",
        "introduction": "Deep submillimeter-wave surveys for distant galaxies offer direct access to high redshifts, because their sensitivity is almost independent of galaxy luminosity for a fixed spectral energy distribution (SED) over the wide redshift range $1<z<5$ (Smail et al. 2002). The positional accuracy of these surveys is only $\\simeq 8$\\,arcsec, and optical counterparts are typically faint. Hence it has been difficult to identify and measure the redshifts of a substantial fraction of these submillimeter-selected galaxies (SMGs).  These difficulties have now been overcome with rest-UV spectroscopy of SMGs located in very deep ($\\simeq 10\\,\\mu$Jy) radio images by Chapman et al. (2003b; 2004a). Redshifts have been found for about half of the SMG population using the Keck I telescope's blue-sensitive multiobject spectrograph LRIS-B. The resulting sample of 73 spectroscopic redshifts is very valuable, allowing a fairly accurate restframe SED to be determined for each galaxy, under reasonable assumptions about their radio--far-IR properties. The SED model is most strongly affected by the value of the effective dust temperature of the galaxies $T_{\\rm d}$, which can be treated as a proxy for the peak restframe wavelength of the SED. Just two parameters, $T_{\\rm d}$ and the far-IR luminosity $L$\\footnote{ Defined as the luminosity between wavelengths of 3\\,mm and 10\\,$\\mu$m, and typically close to the bolometric luminosity.} allow the parameter space of SEDs that describe observational data at wavelengths longer than about 20\\,$\\mu$m to be spanned reasonably well (Dale et al.\\ 2001; Blain, Barnard \\& Chapman 2003; hereafter BBC03). Two other parameters are required to describe the detailed shape of the SED. The Rayleigh--Jeans spectral slope of dust emission is fixed by the emissivity index of dust $\\beta$ ($f_\\nu \\propto \\nu^{2+\\beta}$). The value of $\\beta$ is degenerate with temperature, in the sense that models in which $T_{\\rm d}\\beta$ is approximately constant provide similarly accurate descriptions of the SED (see Fig.\\,2 of BBC03). This effect could be important if $T_{\\rm d}$ were treated as a real physical temperature, but when $T_{\\rm d}$ is used as the key descriptive parameter in the SED the effect of this degeneracy between the fitted values of $T_{\\rm d}$ and $\\beta$ is not important. The slope of the mid-IR spectrum $f_\\nu \\propto \\nu^\\alpha$ is set by a second parameter, which has little effect on the value of $L$ at fixed $T_{\\rm d}$.  Spectroscopic redshifts are especially important for understanding the SEDs of SMGs as there is a very significant, pernicious degeneracy between the inferred dust temperature $T_{\\rm d}$ and redshift $z$, which only allows accurate measurement of the quantity $T_{\\rm d}/(1+z)$ (Blain 1999; BBC03). The large Chapman et al. sample of secure spectroscopic redshifts for SMGs provides a way to break this degeneracy, and thus unambiguously reveals the range of SED properties of the high-redshift population of SMGs. This is the population of galaxies for which radio--submillimeter--far-IR photometric redshifts are often sought (Carilli \\& Yun 1998; Hughes et al.\\ 2002; Yun \\& Carilli 2002; Aretxaga et al. 2003; Wiklind 2003), and we use this new, accurate SED information, to address the precision of these radio--far-IR photometric redshift techniques. We then compare the SEDs of low- and high-redshift samples of luminous dusty galaxies to highlight a region of the $L$--$T_{\\rm d}$ parameter space where unrecognized populations could still lurk, and discuss ways to find them.  Throughout this paper we assume a cosmology with $\\Omega_0=0.3$, $\\Omega_\\Lambda=0.7$ and $H_0=100h$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$ with $h=0.65$ to remain consistent with earlier work. In the {\\it WMAP} cosmology with $\\Omega_0=0.27$, $\\Omega_\\Lambda=0.73$ and $h=0.71$ the quoted luminosities would be reduced by a factor of about 15\\% at $z \\simeq 2.5$. ",
        "conclusions": "The recent acquisition of a large sample of SMG redshifts provides the opportunity to constrain the SEDs of the most luminous ($> 10^{12}\\,L_\\odot$) high-redshift galaxies. It is now reasonably certain that the typical high-redshift SMGs selected at 850-$\\mu$m are systematically cooler than local dusty galaxies of comparable luminosity, and are thus without direct low-redshift analogs. The scatter in their dust temperatures is about 50\\%, sufficient to introduce a $\\sim 30\\%-40$\\% uncertainty in {\\emph any} photometric redshifts derived from far-IR, submillimeter, or radio photometry, no matter how accurate. The true scatter of the dust temperatures in high-redshift dusty galaxies certainly could be still greater, but hotter examples cannot be selected by existing submillimeter surveys. Faint radio galaxies that are undetected in the submillimeter, but whose optical properties are like those of the SMGs hint at the existence of a hotter population, with luminosities $L \\simeq 10^{13}\\,L_\\odot$ and temperatures $T_{\\rm d} \\simeq 60$\\,K. These galaxies would produce a luminosity density comparable to that of the existing populations of SMGs. The deepest forthcoming mid-IR {\\it Spitzer} surveys should soon confirm whether this population exists, while providing a complete inventory of the range of SEDs present in high-redshift dusty galaxies in fields that have excellent supporting long-wavelength data."
    },
    "0404/astro-ph0404180_arXiv.txt": {
        "abstract": "{ The availability of a number of new interferometric measurements of Main Sequence and subgiant stars makes it possible to calibrate the surface brightness relations of these stars using exclusively direct angular diameter measurements. These empirical laws make it possible to predict the limb darkened angular diameters $\\theta_{\\rm LD}$ of dwarfs and subgiants using their dereddened Johnson magnitudes, or their effective temperature. The smallest intrinsic dispersions of $\\sigma \\le 1\\,\\%$ in $\\theta_{\\rm LD}$ are obtained for the relations based on the $K$ and $L$ magnitudes, for instance $\\log \\theta_{\\rm LD} = 0.0502\\,(B-L) + 0.5133 - 0.2\\,L$ or $\\log \\theta_{\\rm LD} = 0.0755\\,(V-K) + 0.5170 - 0.2\\,K$. Our calibrations are valid between the spectral types A0 and M2 for dwarf stars (with a possible extension to later types when using the effective temperature), and between A0 and K0 for subgiants. Such relations are particularly useful for estimating the angular sizes of calibrators for long-baseline interferometry from readily available broadband photometry. ",
        "introduction": "\\label{sec:int} The surface brightness (hereafter SB) relations link the emerging flux per solid angle unit of a light-emitting body to its color, or effective temperature. These relations are of considerable astrophysical interest, as a well-defined relation between a particular color index and the surface brightness can provide accurate predictions of the stellar angular diameters. Such predictions are essential for the calibration of long-baseline interferometric observations. We propose in the present paper new and accurate calibrations of the SB-color relations of dwarfs and subgiants based on direct interferometric measurements of nearby members of these two luminosity classes. Our primary purpose is to establish reliable relations that can be used to predict the angular sizes of calibrator stars for long-baseline interferometry.  After defining the surface brightness relations (Sect.~\\ref{def_relations}), we discuss in Sect.~\\ref{section_sample} the sample of measurements that we selected for our calibrations (interferometric and photometric data). Sect.~\\ref{section_relations} is dedicated to the calibration of the empirical SB relations, relative to the color indices and to the effective temperature, for stars of spectral types A0 to M2. We also derive inverse relations to estimate the effective temperature from broadband photometry and angular diameter measurements. As the established relations are intended to be used primarily to predict angular diameters, we discuss in Sect.~\\ref{estim_errors} their associated errors in this context. In Sect.~\\ref{comparison_section}, we search for a possible instrumental bias linked to one of the five interferometric instruments represented in our sample. Numerous versions of the SB relations have been established in the literature, mostly for giants and supergiants, and we discuss them in Sect.~\\ref{compare_rel}. Main Sequence stars are potentially very good calibrators for long-baseline interferometry, and we discuss this particular application of our SB relations in Sect.~\\ref{calib_interf}. ",
        "conclusions": "The laws that we established between the angular size and broadband colors (or effective temperature) are strictly empirical. Our best relations present a very small intrinsic dispersion, down to less than 1\\,\\%. They can be used to predict the angular sizes of A0--M2 dwarfs and A0--K0 subgiants from simple, readily available broadband photometry. On the one hand, Gray et al.~(\\cite{gray03}) have recently published an extensive survey of the spectral properties of nearby stars within 40\\,pc, including estimates of their effective temperatures. On the other hand, several large catalogues (2MASS, DENIS,...) provide high precision magnitudes of these stars in the infrared. From the cross-comparison of these sources, the SB relations determined in the present paper make it possible to assemble a catalogue of calibrators for interferometry that will be practically unaffected by interstellar extinction, multiplicity or circumstellar material biases. These resulting angular diameter predictions will provide a reliable basis for the calibration of long-baseline interferometric observations."
    },
    "0404/astro-ph0404149_arXiv.txt": {
        "abstract": "We report on the measurement at 820- and 1400-MHz of orbital modulation of the diffractive scintillation timescale from pulsar A in the double-pulsar system \\psr\\ using the Green Bank Telescope. Fits to this modulation determine the systemic velocity in the plane of the sky to be \\viss\\ $\\simeq$140.9$\\pm$6.2$\\kms$.  The parallel and perpendicular components of this velocity with respect to the line of nodes of the pulsar's orbit are \\vplane\\ $\\simeq$96.0$\\pm$3.7$\\kms$ and \\vperp\\ $\\simeq$103.1$\\pm$7.7$\\kms$ respectively.  The large \\vperp\\ implies that pulsar B was born with a kick speed of $\\ga$100$\\kms$.  Future VLBA determination of the angular proper motion in conjunction with improved \\viss\\ measurements should provide a precise distance to the system.  Using high-precision timing data and the \\viss\\ model, we estimate a best-fit orbital inclination of $i=88\\fdg7\\pm0\\fdg9$. ",
        "introduction": "For over twenty years observers have known that measurements of the decorrelation bandwidth \\nud\\ and scintillation timescale \\td\\ of pulsars undergoing strong Diffractive Interstellar Scintillation (DISS) can be used to estimate their velocity in the plane of the sky \\citep{ls82}.  \\citet{cr98} examined DISS-derived pulsar velocities in detail and found that the measurements depend heavily on the observing frequency, the direction and distance to the pulsar, and the detailed distribution of the interstellar material causing the scintillation. This last point makes measurements of pulsar velocities particularly difficult since different models for the electron distribution and its irregularities can cause differences in the ``measured'' scintillation velocities (\\viss) by factors of a few.  However, for binary pulsars in compact orbits (orbital periods $P_{\\rm orb} \\la 1$\\,day), the pulsar-timing-derived orbital velocities can be used to calibrate \\viss\\ measurements and remove many model-dependent and/or systematic effects.  Unfortunately, suitable binary pulsars are rare and successful measurements of this kind have only been made for two pulsars: PSR~B0655$+$64 by \\citet{lyn84} and PSR~J1141$-$6545 by \\citet*[hereafter OBvS]{obvs02}.  \\psrs A \\& B \\citep[hereafter A and B;][]{bdp+03,lbk+04} comprise a fantastic double-pulsar binary (hereafter 0737) consisting of the 22.7-ms pulsar A and the 2.77-s pulsar B.  It is nearby ($\\sim$0.6\\,kpc), mildly eccentric ($e\\sim0.088$), compact ($P_{\\rm orb}\\sim2.45$\\,h), highly inclined ($i\\sim87\\degr$), strongly relativistic, and displays eclipses of A and very large but systematic flux variations of B.  Its proximity, relatively high measured flux density, and rapidly moving pulsars (orbital velocities $\\sim$300$\\kms$), make it an ideal target for \\viss\\ studies.  In this paper we report measurements of the orbital modulation of \\viss\\ from A at 820- and 1400-MHz using data from the 100-m Green Bank Telescope (GBT). ",
        "conclusions": "We have measured the systemic velocity and inclination of 0737 using the ``self-calibrating'' method of binary scintillation velocity measurements.  The inferred high velocity \\viss$\\simeq$140$\\kms$ strongly suggests that a substantial ($\\ga$100$\\kms$) kick was imparted to pulsar B at its birth.  The large \\viss\\ will impact long-term timing of the system and will allow a precise geometric distance measurement when combined with a VLBA proper motion.  Future scintillation observations using the high frequency resolution and wide bandwidth of the GBT$+$SPIGOT will allow substantially improved measurements of \\viss\\ at 1400 and even 2200\\,MHz where we have already detected the orbital modulation of A in BCPM data.  These measurements should also allow the removal of the contaminating effects of the Earth's motion and the differential Galactic rotation. Finally, we have already detected scintillation from B during the bright portions of its orbit, but future observations may allow ``snapshot'' calibrations of \\viss\\ based on measurements of \\td\\ from A and B at an instant in time given our knowledge of the relative orbital velocities of the pulsars.  Such measurements would demand far less telescope time.  {\\em Acknowledgements} We extend special thanks to Karen O'Neil, Rich Lacasse, Glen Langston, and Chris Clark for help with the observations.  The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.  V.M.K. acknowledges support from NSERC Discovery Grant 228738-03, NSERC Steacie Supplement 268264-03, a Canada Foundation for Innovation New Opportunities Grant, FQRNT Team and Centre Grants, and NASA Long-Term Space Astrophysics Grant NAG5-8063."
    },
    "0404/gr-qc0404081_arXiv.txt": {
        "abstract": "Kerr--Schild solutions to the vacuum Einstein equations are considered from the viewpoint of integral equations.  We show that, for a class of Kerr--Schild fields, the stress-energy tensor can be regarded as a total divergence in Minkowski spacetime.  If one assumes that Minkowski coordinates cover the entire manifold (no maximal extension), then Gauss' theorem can be used to reveal the nature of any sources present.  For the Schwarzschild and Vaidya solutions the fields are shown to result from a $\\del$-function point source.  For the Reissner--Nordstrom solution we find that inclusion of the gravitational fields removes the divergent self-energy familiar from classical electromagnetism.  For more general solutions a complex structure is seen to arise in a natural, geometric manner with the role of the unit imaginary fulfilled by the spacetime pseudoscalar. The Kerr solution is analysed leading to a novel picture of its global properties.  Gauss' theorem reveals the presence of a disk of tension surrounded by the matter ring singularity.  Remarkably, the tension profile over this disk has a simple classical interpretation.  It is also shown that the matter in the ring follows a light-like path, as one expects for the endpoint of rotating, collapsing matter.  Some implications of these results for physically-realistic black holes are discussed. ",
        "introduction": "Many of the important solutions to the Einstein field equations can be expressed in Kerr--Schild form (see, for example, the discussion in~\\cite{kra-exact}).  These include all black hole solutions, and a range of solutions representing radiation.  Here we analyse solutions of Kerr--Schild type from the viewpoint of the gauge theory approach to gravity~\\cite{DGL98-grav,gap,DLkerr03}.  In this approach the gravitational fields are gauge fields defined over a flat Minkowski spacetime.  These fields ensure that all relations between physical quantities are independent of the position and orientation of the matter fields --- a scheme that ensures that the background spacetime plays no dynamic role in the physics and has no measurable properties. Kerr--Schild metrics are constructed from a null vector field in the background Minkowski spacetime, so are particularly well-suited to analysis via this gauge-theoretic approach.  In this paper we show that, for all fields of Kerr--Schild type, the Einstein tensor is a total divergence in the background Minkowski spacetime.  Various consequences of this result are explored.  Gauss' theorem is used to convert volume integrals of the Einstein tensor to surface integrals, enabling us to probe the nature of the matter singularities generating the gravitational fields.   The gauge-theory viewpoint always produces a metric that satisfies the Einstein equations (or their generalisation to include torsion).  But working with fields defined over a Minkowski background does place additional restrictions on the form of the solutions.  For example, general relativity admits two possibilities when dealing with the Kerr solution~\\cite{isr70} \\begin{enumerate} \\item The complete Kerr manifold can be covered by a single set of Minkowski coordinates.  This implies a discontinuity in the fields over the entire disk region bounded by the matter singularity. \\item The fields are smooth everywhere away from the ring, but an observer passing through the ring emerges in a new, asymptotically flat, region.  This is achieved by extending the radial coordinate $r$ to negative values, producing the maximally-extended Kerr spacetime. \\end{enumerate} By adopting a flat-spacetime, gauge-theory formulation we restrict ourselves to considering case~1 only.  This can be justified on the grounds that the full, maximally-extended Kerr solution is not thought to be a feasible endpoint for any collapse process.  Similar comments apply to the maximally-extended Schwarzschild and Reissner--Nordstrom solutions, neither of which are considered here.   The first applications we consider are to spherically-symmetric fields, concentrating on the Schwarzschild, Reissner--Nordstrom and Vaidya solutions.  In all cases the integrals provide sensible results for the total energy contained in the fields, with the mass contribution to the energy residing in a point-source $\\del$-function. For the Reissner--Nordstrom solution the inclusion of gravitational fields removes the infinite electromagnetic self-energy for a point charge familiar from classical electromagnetism~\\cite{fey-lectII}. This result is achieved without requiring any form of regularisation procedure, and ensures that the total electromagnetic self-energy is zero.   We next turn to more general fields following the work of Schiffer \\etal~\\cite{sch73}.  These authors showed that stationary Kerr--Schild vacuum solutions are generated by a single, complex generating function.  This complex structure underlies the `trick' by which the Kerr solution is obtained from the Schwarzschild solution via a complex `coordinate transformation'~\\cite{new65}.  (This is a trick because there is no \\textit{a priori} justification for expecting the complex transformation to result in a new vacuum solution.)  The complex structure associated with vacuum Kerr--Schild fields is shown here to have a simple geometric origin, with the role of the unit imaginary fulfilled by the spacetime pseudoscalar --- the same entity that is responsible for duality transformations of the Riemann tensor.   The remainder of this paper deals with a detailed analysis of the Kerr solution.  For this we require a careful choice of branch cut in the complex square route in the generating function.  Once this is made, Gauss' theorem reveals the detailed structure of the singular region, confirming that the matter is concentrated in a ring that circulates on a lightlike trajectory.  This is as one would expect, since the Riemann tensor only diverges on a ring, and special relativity alone is sufficient to predict that rotating collapsing matter will fall inwards until its velocity becomes lightlike.  A more surprising result is obtained from considering integrals inside the ring, which reveal the presence of a disc of planar tension~\\cite{DGL96-erice}. This tension is isotropic over the disk and has a simple radial dependence, rising to $\\infty$ at the ring.  Remarkably, the functional form of the tension has a simple non-gravitational interpretation.  In non-relativistic dynamics a membrane holding together a rotating ring of disconnected particles would be under a constant tension.  When special-relativistic effects are included the picture is altered by the fact that tension can act as a source of inertia.  This introduces a radial dependence into the tension, the functional form of which is precisely that which lies at the heart of the Kerr solution.  These conclusions are gauge invariant and are not artifacts of the use of the background spacetime.  This is demonstrated by eigen-decompositions of the stress-energy and Riemann tensors, from which we extract the gauge-invariant information.   There has been considerable debate over many years surrounding the nature of sources for the Kerr metric.  Many physicists have attempted to construct extended sources for which the Kerr metric could represent the external geometry (see Krasinski for an early review~\\cite{kra78}).  More recently, a series of authors have constructed disk sources for the Kerr metric~\\cite{bic93,neu93,pic96}. These solutions represent extended sources and do not have horizons present.  The present work is of a different nature, dealing solely with the structure of the singular region --- the endpoint of a collapse process.  The first authors to consider this were Newman \\& Janis~\\cite{new65} and Isreal~\\cite{isr70}.  We disagree with Isreal's result for the energy distribution over the disk, agreeing instead with Hamity's later result~\\cite{ham76}.  Our techniques enable us to go some way beyond Hamity's description, both in revealing the physical properties of the disk and in understanding the nature of the singularity around the ring.  The simple physics of the disk was first pointed out in~\\cite{DGL96-erice}.   Many of the calculations here are simplified by using the language of `spacetime algebra'~\\cite{gap,hes-sta,hes-gc}.  This is crucial to understanding the geometric nature of the complex structure at the heart of Kerr--Schild solutions.  The algebraic structure of the spacetime algebra is that of the Dirac $\\gam$-matrices.  Using this algebraic structure one can develop a mathematical language that is adept at describing many aspects of relativistic physics.  This language includes a calculus that is somewhat more powerful than any available in alternative languages.  The gauge theory of gravity developed in~\\cite{DGL98-grav} takes on its most natural and compelling form when expressed in the spacetime algebra.  We start with an introduction to the spacetime algebra, giving the necessary conventions and notations.  Further details can be found in~\\cite{DGL98-grav,gap,hes-gc,DGL95-elphys} and references contained therein.  Reference~\\cite{DGL98-grav} includes an appendix describing how to convert between spacetime algebra and more conventional tensor calculus.  Natural units ($G=c=\\eps_0=1$) are employed throughout this paper.   \\subsection{Spacetime algebra}  The basic algebraic structure behind the spacetime algebra will be familiar to most physicists in the guise of the algebra of the Dirac $\\gamma$-matrices.  The geometric interpretation the spacetime algebra attaches to this algebra may be less familiar, though it is remarkably well-suited to most problems in relativistic physics~\\cite{gap,hes-sta,DGL95-elphys}.  The spacetime algebra is generated by four vectors $\\{\\gamdm\\}$, $\\mu=0\\ldots 3$, equipped with an associative (Clifford) product denoted by juxtaposition.  The symmetric and antisymmetric parts of this product define the inner and outer products, and are denoted with a dot and a wedge respectively, so \\begin{equation} \\gamdm \\dt \\gamdn \\eqv \\half (\\gamdm \\gamdn + \\gamdn\\gamdm) = \\eta_{\\mu \\nu} = \\mbox{diag($+$\\ $-$\\ $-$\\ $-$)} \\end{equation} and \\begin{equation} \\gamdm \\wdg \\gamdn \\eqv \\half (\\gamdm \\gamdn - \\gamdn\\gamdm). \\end{equation} The outer product of two vectors defines a bivector --- a directed plane segment representing the plane defined by the two vectors.  A full basis for the spacetime algebra is provided by the set \\begin{equation} \\begin{array}{ccccc} 1 & \\{\\gamdm\\} & \\{\\bsig_k, I\\bsig_k\\} & \\{I\\gamdm\\} & I \\\\ \\mbox{1 scalar} & \\mbox{4 vectors} & \\mbox{6 bivectors} & \\mbox{4 trivectors} & \\mbox{1 pseudoscalar}, \\\\ \\mbox{grade 0} & \\mbox{grade 1} & \\mbox{grade 2} & \\mbox{grade 3} & \\mbox{grade 4} \\end{array} \\label{basis} \\end{equation} where \\begin{equation} \\bsig_k \\eqv \\gam_k \\go, \\quad  k=1\\ldots 3 \\end{equation} and \\begin{equation} I\\eqv\\go\\gi\\gj\\gk=\\bsi\\bsj\\bsk. \\end{equation} The pseudoscalar $I$ squares to $-1$, anticommutes with all odd-grade elements and commutes with even grade elements.  Both the $\\{\\bsig_k\\}$ and $\\{\\gamdm\\}$ are algebraic entities with clear geometric significance.  They should not be thought of as matrices acting on an internal spin space.  (The same symbols as employed in quantum theory are used here simply because the algebraic relations are the same.)   An arbitrary real superposition of the basis elements~\\eqref{basis} is called a `multivector' and these inherit the associative Clifford product of the $\\{\\gamdm\\}$ generators.  The inner and outer products with a vector $a$ are of particular importance.  For these we write \\begin{equation} a \\dt A_r = \\half(a A_r - (-1)^r A_r a), \\qquad a \\wdg A_r = \\half(a A_r + (-1)^r A_r a). \\end{equation} The outer and geometric products are associative, but the inner product is not.  We also employ the commutator product, \\begin{equation} A \\crs B \\eqv \\half(AB - BA). \\end{equation} Vectors are usually denoted in lower case Latin, $a=a^\\mu \\gamdm$, or Greek for basis frame vectors.  In the absence of brackets the inner, outer and commutator products take precedence over geometric products.   An inertial system is picked out by a future-pointing timelike (unit) vector. If this is chosen to be the $\\go$ direction then the $\\go$-vector determines a map between spacetime vectors $a=a^\\mu\\gamdm$ and the even subalgebra of the full spacetime algebra via \\begin{equation} a \\go = a_0 + \\ba, \\label{1sptsplt} \\end{equation} where \\begin{equation} a_0 = a \\dt \\go, \\hs{0.5} \\mbox{and} \\hs{0.5} \\ba = a \\wdg \\go. \\end{equation} The `relative vector' $\\ba$ can be decomposed in the $\\{\\bsig_k\\}$ frame and represents a spatial vector as seen by an observer in the $\\go$-frame.  Relative (or spatial) vectors in the $\\go$-system are written in bold type to record the fact that they are in fact spacetime bivectors.  This distinguishes them from spacetime vectors, which are left in normal type.  The $\\{\\bsig_k\\}$ generate the (Pauli) algebra of three-dimensional space, and we occasionally require that the dot and wedge symbols define the three-dimensional inner and outer products.  The convention we adopt is that, if both arguments of a dot or wedge product are written in bold, then the product takes its three-dimensional meaning.  For example, $\\ba \\wdg \\bb$ is a relative bivector, and so also a spacetime bivector, and not a spacetime four-vector.   The vector derivative, $\\grad$, is defined by \\begin{equation} \\grad \\eqv \\gamum \\deriv{}{x^\\mu} \\end{equation} where the $\\{x^\\mu\\}$ are a set of Cartesian coordinates and the $\\{\\gamum\\}$ are the reciprocal frame to the associated coordinate frame $\\{\\gamdm\\}$, \\textit{i.e.\\/} $\\gamma^\\mu \\dt \\gamma_\\nu = \\delta^\\mu_\\nu$.  The spacetime split of the vector derivative $\\grad$ goes through slightly differently, since we require that the $\\bgrad$ symbol agrees with its conventional three-dimensional meaning.  This is achieved by writing \\begin{equation} \\go \\grad = \\dift + \\bgrad, \\end{equation} so that $\\bgrad=\\bsig_i\\partial_{x^i}$.  The $\\grad$ operator has the algebraic properties of a vector, and often acts on objects to which it is not adjacent.  The `overdot' notation is a convenient way to encode this: \\begin{equation} \\dgrad A \\dot{B} \\eqv \\gamum A \\deriv{B}{x^\\mu}. \\end{equation} The $\\grad$ operator acts on the object to its immediate right unless brackets or overdots are present.  If brackets are present then $\\grad$ operates on everything in the bracket, so that, for example, $\\grad(AB)= \\grad A B + \\dgrad A \\dot{B}$.  The same rules apply to $\\bgrad$.   One of the two gravitational gauge fields is the (position-dependent) linear function $\\hu(a)$, which maps vectors to vectors (where $a$ is the vector argument).  Linear functions of this type have their action extended to general multivectors via the rule \\begin{equation} \\hu(a \\wdg b \\cdots \\wdg c) \\eqv \\hu(a) \\wdg \\hu(b) \\wdg \\cdots \\wdg \\hu(c), \\end{equation} which defines a grade-preserving linear operation.  The pseudoscalar is unique up to a scale factor, and the determinant is defined by \\begin{equation} \\hu(I) = \\det(\\hu) I. \\end{equation} The adjoint is denoted with an overbar, $\\ho(a)$.  The function $\\hu(a)$ and its adjoint are related by~\\cite{hes-gc} \\begin{align} A_r \\dt \\ho(B_s) &= \\ho( \\hu(A_r) \\dt B_s) \\qquad r \\leq s  \\nn \\\\ \\hu(A_r) \\dt B_s &= \\hu( A_r \\dt \\ho(B_s)) \\qquad r \\geq s . \\end{align}   A number of manipulations in linear algebra are simplified by using the vector derivative in place of frame contractions.  For example, the trace of $\\hu(a)$ can be written as \\begin{equation} \\mbox{Tr} (\\hu) = \\gamum \\dt \\hu (\\gamdm) = \\da \\dt \\hu (a), \\end{equation} where $\\da$ is the vector derivative with respect to $a$.  The following results are also useful: \\begin{align} \\da \\, a \\dt A_r &= r A_r \\\\ \\da \\, a \\wdg A_r &= (n-r) A_r \\\\ \\da A_r a &= (-1)^r (n-2r) A_r , \\end{align} where $A_r$ is a multivector of grade $r$ and $n$ is the dimension of the space.    \\subsection{The field equations}  The gravitational gauge fields are a linear function $\\ho(a)$ mapping vectors to vectors and a linear function $\\Om(a)$ mapping vectors to bivectors.  Both of these gauge fields have an arbitrary position dependence.  The gauge-theoretic origin of these fields is described in~\\cite{DGL98-grav,gap}.  The gauge fields are related by the equation \\begin{equation} 2\\Om(a) = -\\ho(\\grad \\wdg \\lig(a)) + \\hu^{-1}(\\db) \\wdg (a \\dt \\grad \\ho(b)), \\end{equation} where \\begin{equation} \\lig(a) \\eqv \\ho^{-1} \\hu^{-1} (a). \\end{equation} The argument of the linear function, usually denoted by a vector $a$ or $b$, is always assumed to be independent of position. To recover the more conventional representation of general relativity we introduce an arbitrary set of coordinates $x^\\mu$, with $e_\\mu$ the associated coordinate frame vectors, \\begin{equation} e_\\mu \\eqv \\deriv{x}{x^\\mu}. \\end{equation} With $e^\\mu$ denoting the reciprocal frame vectors we then define the vectors \\begin{equation} g_\\mu = \\hu^{-1}(e_\\mu) , \\qquad g^\\mu = \\ho(e^\\mu). \\end{equation} In terms of these the metric tensor is defined by \\begin{equation} g_{\\mu\\nu} = g_\\mu \\dt g_\\nu. \\end{equation} The $\\ho(a)$ field ensures that one only ever has to make `flat-space' contractions, which is an attractive feature of the gauge-theory approach.    The field strength corresponding to the $\\Om(a)$ gauge field is defined by \\begin{equation} \\liR(a\\wdg b) \\eqv a\\dt\\grad \\Om(b) - b \\dt \\grad \\Om(a) + \\Om(a) \\crs \\Om(b) \\end{equation} and is a linear function mapping bivectors to bivectors.  From this the covariant Riemann tensor is defined by \\begin{equation} \\clr(a\\wdg b) \\eqv \\liR(\\hu(a \\wdg b)). \\end{equation} We often write this in the form $\\clr(B)$, where $B$ is an arbitrary (constant) bivector argument.  The tensor components of the Riemann tensor are recovered by writing \\begin{equation} {R^\\mu}_{\\nu \\rho \\sigma} = (g^\\mu \\wdg g_\\nu) \\dt \\clr(g_\\sigma \\wdg g_\\rho). \\end{equation} The Ricci and Einstein tensors are defined from the Riemann tensor in the obvious way, \\begin{alignat}{2} \\mbox{Ricci Tensor:}& & \\quad  \\clr(b)&\\eqv \\da \\dt \\clr(a \\wdg b) \\\\ \\mbox{Ricci Scalar:}& & \\quad \\clr &\\eqv \\da \\dt \\clr(a) \\\\ \\mbox{Einstein Tensor:}& &\\quad \\clg(a)&\\eqv \\clr(a) - \\half a \\clr. \\end{alignat} Again, the tensor components of the Ricci and Einstein tensors are easily recovered.   \\subsection{Kerr--Schild fields}   We are interested in fields of the form \\begin{equation} \\ho(a) = a + a \\dt l \\, l \\label{ksanz} \\end{equation} where $l$ is a (flat-space) null vector, $l^2=0$.  This is the gauge theory analogue of the Kerr--Schild ansatz.  The function~\\eqref{ksanz} extends to act on multivectors as \\begin{equation} \\ho(A) = \\hu(A) = A + A \\dt l \\, l, \\end{equation} and we see immediately that $\\det(\\ho)=1$.  The following results are also useful: \\begin{gather} \\hu^{-1}(A) = \\ho^{-1}(A) = A - A \\dt l\\, l \\\\ \\lig(A) = A - 2 A \\dt l \\, l \\\\ \\ho(l) = l. \\end{gather} In terms of an orthonormal coordinate frame $\\gamdm$ we can write \\begin{equation} g_\\mu = \\gamdm - l_\\mu l \\end{equation} which confirms that the metric is given by \\begin{equation} g_{\\mu\\nu} = \\eta_{\\mu\\nu} - 2 l_\\mu l_\\nu, \\end{equation} where $\\eta_{\\mu\\nu}$ is the flat Minkowski metric tensor.   The $\\Om(a)$ field defined by~\\eqref{ksanz} has the simple form \\begin{align} \\Om(a) &= \\ho\\bigl( \\grad \\wdg (a \\dt l \\, l) \\bigr) \\nn \\\\ &= \\grad \\wdg (a \\dt l \\, l) - a \\dt l \\, v \\wdg l \\label{ksom} \\end{align} where \\begin{equation} v \\eqv l \\dt \\grad l. \\end{equation} It follows from the fact that $l$ is null that \\begin{equation} l \\dt v =0 \\end{equation} and \\begin{equation} \\Om(l) = 0. \\end{equation}   Following the route adopted by Chandrasekhar~\\cite[Section 57]{cha83}, we next form the quantity \\begin{align} l \\dt \\clr(l) &= l \\dt (\\da \\dt \\clr(a \\wdg l) ) \\nn \\\\ &= (l \\wdg \\da) \\dt \\liR(a \\wdg l) \\nn \\\\ &= (l \\wdg \\da) \\dt \\bigl(a \\dt \\dgrad \\dot{\\Om}(l) - l \\dt \\grad \\Om(a)\\bigr). \\end{align} Substituting equation~\\eqref{ksom} into the above we find that \\begin{align} l \\dt \\clr (l) &=  (l \\wdg \\da) \\dt \\bigl( - \\dgrad((a\\dt \\grad l) \\dt \\ldot) \\wdg l  - l \\dt \\grad \\grad \\wdg (a \\dt l \\, l) \\bigr) \\nn \\\\ &= \\da \\dt l \\, (a \\dt \\grad l) \\dt v - l \\dt \\grad (\\grad \\dt l \\, l + v) \\dt l \\nn \\\\ &= v^2 - (l \\dt \\grad v) \\dt l \\nn \\\\ &= 2 v^2. \\end{align} If we are looking solely for vacuum solutions, then we can conclude from this that $v$ must be null.  Since $v \\dt l=0$, it follows that $v$ must be parallel to $l$, \\begin{equation} v = \\phi l, \\label{leqn} \\end{equation} where $\\phi$ is a scalar field.  We will restrict attention to solutions for which this relation does hold, even if matter is present.  (This places a restriction on the form of matter distributions that we can consider.)  It follows from equation~\\eqref{leqn} that $\\Om(a)$ reduces to the simpler form \\begin{equation} \\Om(a) = \\grad \\wdg (a \\dt l \\, l). \\label{newOm} \\end{equation}   The Riemann tensor now splits into terms that are second-order and fourth-order in $l$.  The fourth-order contribution is \\begin{equation} \\clr_4 (a \\wdg b) = - \\Omdot \\bigl(((a \\wdg b) \\dt l \\, l) \\dt \\dgrad \\bigr) + \\Om(a) \\crs \\Om(b). \\label{Riem4} \\end{equation} After some rearrangement this can be brought to the form \\begin{equation} \\clr_4 (B) = \\qrt \\da \\dt \\db \\, (a \\dt \\grad l) lBl (b \\dt \\grad l) - \\qrt (a \\dt \\grad l) \\dt (b \\dt \\grad l) \\, \\da l B l \\db. \\label{clr4} \\end{equation} Both the contraction, $\\da \\cdot \\clr(a \\wedge b)$, and the protraction, $\\da \\wedge \\clr(a \\wedge b)$, of this contribution to the Riemann tensor vanish.  This can be seen from the result that \\begin{equation} \\da F_1 a \\wdg b F_2 = \\da F_1 (ab - a \\dt b) F_2 = -b F_1 F_2, \\end{equation} which holds for any two bivectors $F_1$ and $F_2$.  The presence of the null vector $l$ in the analogous terms in $\\clr_4(B)$ ensures that \\begin{equation} \\da \\clr_4 (a\\wdg b) = 0, \\end{equation} so that $\\clr_4(B)$ makes no contribution to the Ricci tensor.   The only part of $\\clr(B)$ that contributes to the Einstein tensor is therefore the second-order term \\begin{equation} \\clr_2(a \\wdg b) = a \\dt \\grad \\Om(b) - b \\dt \\grad \\Om(a). \\label{KSeq1} \\end{equation} Contracting this and setting the result to zero we find that the vacuum Einstein equations reduce to solving the equation \\begin{equation} \\clr(a) = \\grad \\dt \\Om(a) - a \\dt \\grad \\, \\db \\dt \\Om(b) = 0. \\end{equation}  The Ricci scalar and Einstein tensor are now straightforward to calculate: \\begin{equation} \\clr = -2 \\grad \\dt (\\da \\dt \\Om(a)) \\end{equation} and \\begin{equation} \\clg(a) = \\grad \\dt \\bigl( \\Om(a) - a \\wdg (\\db \\dt \\Om(b)) \\bigr). \\label{Einst} \\end{equation} The formulae for $\\Om(a)$~\\eqref{newOm} and $\\clg(a)$ are valid for any Kerr--Schild type solution for which $l\\dt\\grad l=\\phi l$.  For such fields the Einstein tensor~\\eqref{Einst} is a total divergence in Minkowski spacetime.  In general, the field equations will be satisfied everywhere except for some singular region over which the fields are discontinuous.  This singular region contains the source of the fields.  In this paper we assume that the entire solution to the Einstein equations is described by fields defined over a single Minkowski spacetime, so that the manifold has not been subjected to maximal extension.  In this case we can use Gauss' theorem straightforwardly to convert volume integrals over the source region to surface integrals and so learn how the source matter is distributed.  For the case of static fields, Virbhadra~\\cite{vir90} gave a formula which agrees with~\\eqref{Einst} for the timelike component $a=\\go$, but the fact that the expression is a total divergence was not exploited. ",
        "conclusions": "Many of the significant solutions to the Einstein equations can be represented in Kerr--Schild form and gauge-theoretic approach of~\\cite{DGL98-grav} is well suited to their analysis.  For all solutions of Kerr--Schild type where the null vector $l$ satisfies $l\\dt\\grad l =\\phi l$ the Einstein tensor is a total divergence in flat spacetime.  The structure of the sources generating the fields can therefore be elucidated by employing Gauss theorem to transform volume integrals to surface integrals.  This approach is fully justified within the gauge-theory formulation, since one only ever deals with fields defined over a flat spacetime.   For the case of the Schwarzschild, Reissner--Nordstrom and Vaidya solutions the gravitational fields are seen to result from a $\\del$-function point source of mass at the origin.  For the Reissner--Nordstrom solution the $\\del$-function point source is surrounded by a Coulomb field.  An unexpected bonus of this approach is that the infinite self-energy of the Coulomb field is removed by the gravitational field.  Similar techniques can be applied to Kinnersley's and Bonnor's work on accelerating and radiating masses~\\cite{kin69,bon94}, as will be discussed elsewhere.   Applied to more general stationary, vacuum solutions we find that the complex structure at the heart of vacuum Kerr--Schild fields is the same as the natural complex structure inherent in the Weyl tensor through its self-duality symmetry.  Further algebraic insights are obtained through the use of null vectors as idempotent elements, simplifying many of the derivations of the vacuum equations.  Both of these insights highlight the algebraic advantages of the spacetime algebra approach.  A further example of this is seen clearly in equation~\\eqref{defriem}, which gives a remarkably simple and compact expression for the Riemann tensor.   The application of Gauss' theorem to the Kerr solution reveals some surprising features of the singularity.  The ring of matter follows a lightlike trajectory and surrounds a disk of tension.  The tension distribution over the disk is precisely that predicted by special relativity.  The correct tension distribution was computed by Hamity~\\cite{ham76}, though he did not comment on its origin in terms of classical relativistic physics.  We find no evidence of either the negative surface energy density or the superluminal speeds claimed by Isreal~\\cite{isr70}.  Both Hamity and Isreal asserted that they used the same results for surface layers in general relativity, but neither gave detailed calculations, so the reason for Isreal's disagreement with our result is hard to pin down.   Almost all trajectories in the Kerr geometry finish up on the disk, rather than the ring.  Quite what happens when a particle encounters the $\\del$-function tension over the disk is unclear and can only really be understood using a quantum framework to study the effect of the disk on a wavepacket.  (A start on such an analysis in made in~\\cite{D00-kerr}.)  Assuming that all geodesics do terminate on the disk then any non-causal features of the Kerr solution are removed~\\cite{isr70}, which is a physically attractive feature of the picture presented here.  The fact that the tension membrane violates the weak energy condition raises a further interesting question --- how can it be formed from collapsing baryonic matter?  Furthermore, if baryonic matter cannot form the membrane, then what is the endpoint of the collapse process?  Answers to these questions will only emerge when realistic collapse scenarios are formulated, though these are notoriously difficult computations to perform.   The discussion in this paper implicitly rules out considering any extensions to the manifold, such as obtained by converting the Schwarzschild solution to Kruskal coordinates.  We therefore do not consider distinct universes connected by Schwarzschild `throat', with separate future and past singularities~\\cite{mis-grav,haw-large}, or the maximum analytic extension of the Reissner--Nordstrom geometry with infinite ladder of possible `universes' connected by wormholes~\\cite{haw-large,kauf-front}.  In such scenarios the applications of Gauss' theorem employed in this paper would not be valid.  While infinite ladders of connected universes remain popular with science fiction writers, there is no reason to believe they could ever form physically in any collapse process.  The descriptions presented here for both the Reissner--Nordstrom and Kerr solutions have a much more plausible physical feel to them, even if the final description of the singular region must ultimately involve quantum gravity.  One final speculation concerns the nature of the membrane supporting the Kerr ring singularity.  This bears a remarkable similarity to some of the structures encountered in string theory, and it would be of great interest to see if string theory can provide a quantum description of such a source.     \\appendix"
    },
    "0404/astro-ph0404239_arXiv.txt": {
        "abstract": "We discuss the mass assembly history both on cluster and galaxy scales and their impact on galaxy evolution.  On cluster scale, we introduce our on-going PISCES project on Subaru, which plans to target $\\sim$15 clusters at $0.4\\le z\\le 1.3$ using the unique wide-field (30$'$) optical camera Suprime-Cam and the spectrograph both in optical (FOCAS, 6$'$) and near-infrared (FMOS, 30$'$). The main objectives of this project are twofold: (1) Mapping out the large scale structures in and around the clusters on 10--14~Mpc scale to study the hierarchical growth of clusters through assembly of surrounding groups. (2) Investigating the environmental variation of galaxy properties along the structures to study the origin of the morphology-density and star formation-density relations. Some initial results are presented.  On galactic scale, we first present the stellar mass growth of cluster galaxies out to $z\\sim1.5$ based on the near-infrared imaging of distant clusters and show that the mass assembly process of galaxies is largely completed by $z\\sim1.5$ and is faster than the current semi-analytic models' predictions. We then focus on the faint end of the luminosity function at $z\\sim1$ based on the Subaru/XMM-Newton Deep Survey imaging data. We show the deficit of red galaxies below M$^*$+2 or 10$^{10}$~M$_{\\odot}$, which suggest less massive galaxies are either genuinely young or still vigorously forming stars in sharp contrast to the massive galaxies where mass is assembled and star formation is terminated long time ago. ",
        "introduction": "\\begin{figure} \\centerline{ \\psfig{file=cl0016_map.ps,angle=0,width=6.5cm} \\psfig{file=rxj0153_map.ps,angle=0,width=6.5cm} } \\caption{ The panoramic maps of CL0016+16 cluster ($z$=0.55; left panel) and RXJ0152.7$-$1357 ($z$=0.83; right panel). 10~arcminutes correspond to physical scales of 3.8 and 4.6~Mpc, respectively. Using photometric redshift technique based on multi-colour data ($BVRi'z'$ and $VRi'z'$, respectively), plotted are the photometric member candidates selected with redshift cuts of 0.50$\\le$$z$$\\le$0.58 and 0.78$\\le$$z$$\\le$0.86, respectively. Contours show local 2-D number density of galaxies at 1.5, 2, 3, 4, 5 $\\sigma$ above the mean density. Coordinates are shown relative to the centre of the main cluster. Large scale filamentary structures ($>$10Mpc) are seen in both clusters. } \\label{fig:map_photz} \\end{figure}  \\begin{figure} \\centerline{ \\psfig{file=cl0016_red.ps,angle=0,width=6.5cm} \\psfig{file=rxj0153_red.ps,angle=0,width=6.5cm} } \\caption{ The same as Fig.~\\ref{fig:map_photz}, but only red galaxies on the colour-magnitude sequence are shown rather than phot-$z$ selected galaxies. The red galaxies are selected on the basis of $VRi'$ colours for CL0016+16 and $Ri'z'$ colours for RXJ0152.7$-$1357, respectively, and some red colour slice cuts are applied to isolate the passively evolving galaxies at cluster redshifts. This technique gives narrower redshift slice ($\\Delta$$z$$\\sim$0.05) hence has less projection effect, but tends to be biased to the systems dominated by red populations. It is therefore complemental to the phot-$z$ slice technique used in Fig.~\\ref{fig:map_photz}. } \\label{fig:map_red} \\end{figure}  PISCES is a panoramic imaging survey of distant clusters using the Subaru wide-field optical camera Suprime-Cam which provides 34$'$$\\times$27$'$ field of view corresponding to a physical area of 16$\\times$13~Mpc$^2$ at $z\\sim1$. This long term project has started since 2003, and we aim to target $\\sim$15 X-ray selected distant clusters in total at $0.4\\lsim z\\lsim 1.3$, in good coordination with {\\it ACS/HST}, {\\it XMM}, and {\\it Chandra} observations. This unique project is currently underway and some preliminary results on the large scale structures over the entire Suprime-Cam fields are shown in Figs.~\\ref{fig:map_photz} and \\ref{fig:map_red}. These two rich clusters at $z$=0.55 and 0.83 were imaged in multi optical bands, and photometric redshifts (\\cite{k99}) have been applied to efficiently remove foreground/background contaminations and isolate the cluster member candidates (\\cf \\cite{k01}). Many substructures are then clearly seen around the main body of the clusters which tend to be aligned in filamentary structures extending to $>$10~Mpc scale across. Although these structures should be confirmed spectroscopically later on, these already provide good evidence for cluster scale assembly in the hierarchical Universe.  \\begin{figure} \\centerline{ \\psfig{file=cl0939_colden.ps,angle=0,width=6.5cm} } \\caption{ The variation in colour versus local galaxy density, for phot-$z$ selected cluster members brighter than $I=23.4$ in CL0939+47 cluster at $z=0.41$ (\\cite{k01}). The open circles and filled triangles show the galaxies brighter or fainter than $I=21.4$ ($M_V^{\\ast}$+2), respectively. The three red lines represent the loci of the 25, 50, and 75th percentile colours.  The local number density is calculated from the 10 nearest galaxies and we correct this for residual field contamination in the photometric members using the blank field data. } \\label{fig:colden} \\end{figure}  The large scale structures that we see around the clusters provide us the unique opportunities to look into the environmental effects on galaxies as they assemble to denser regions. Kodama \\etal\\, (2001) have presented the environmental dependence of galaxy colours along the filamentary structures around the CL0939+47 cluster at $z$=0.41.  They have shown that the galaxy colours change rather sharply at relatively low density regions such as galaxy groups along the filaments well outside of the cluster core where the galaxies have not yet passed the central region of the clusters yet (see also \\cite{gray04}; \\cite{treu03}). Together with the similar findings in the local Universe (\\cite{lewis02}; \\cite{gomez03}). the environmental effects that truncate star formation are not driven by the cluster specific mechanism such as ram-pressure stripping (\\cite{abadi99}) and but are found to be much wider spread into low density regions. We should therefore pay greater attention to galaxy groups as the key hierarchy for the environmental effects and try to identify what is happening on galaxies in this environment (see below). It is also important to extend this analysis to higher redshifts as the galaxy environment should change dramatically during the course of vigorous assembly, which is probably related to the appearance of morphology-density relation (\\cite{d80}).    Obviously, taking as many spectra as possible from the photometrically identified large scale structures is crucial to prove their reality, since our photometric approach may well suffer from the projection effect along the line of site as we go to lower density regions due to the broad phot-$z$ slice cuts that we apply. Importantly, \\oii\\, line and/or the 4000\\AA\\, break feature are detectable for our PISCES targets in the optical spectroscopy (such as FOCAS) and \\halpha\\, line comes to the FMOS window (0.9$\\mu$m$<$$\\lambda$$<$1.8$\\mu$m). Not only definitively removing the foreground/background contamination and identifying the physically associated real groups, spectroscopic redshifts of individual cluster members will also provide us two critical information: (1) Dynamical mass of the systems which can then be compared to lensing mass and the X-ray mass to address the dynamical state of the systems. (2) 3-D velocity structures, providing the recent and/or near future cluster-cluster/cluster-group merger histories (\\eg \\cite{czoske02}).  Also, \\oii\\, and \\halpha\\, lines will offer the measures of on-going star formation rate of galaxies.  (The latter is preferred since it is much less affected by dust extinction or metallicity variation (\\cite{kenn84})). Therefore we can directly identify when, where and on what timescale the star formation is truncated as the galaxies/groups fall into clusters along the filaments (\\eg \\cite{k04b}). Moreover, we will combine the information from other spectral indices such as Balmer lines and colours in order to resolve the recent star formation histories in galaxies in fine time scales. Different spectral features are sensitive to different stellar ages (\\cite{kenn84}; \\cite{balogh99}; \\cite{bia99}): the emission lines (\\oii\\, \\halpha) will measure the amplitude of on-going star formation (10$^7$ yrs), while the Balmer absorption line give the luminosity (mass) contribution from the stars formed immediately before the truncation (10$^{8-9}$ yrs) (\\cite{d92}; \\cite{couch87}) and the 4000\\AA\\, break and broad-band colours specify the features of longer-lived populations ($>$10$^9$ yrs) (\\cite{kb01}). Resolving the recent star formation histories in galaxies in the transition regions (groups) is the key to understanding the physical processes behind the truncation. In particular, the existence of strong nebular emissions would support the galaxy-galaxy mergers which trigger star-burst, and very strong Balmer absorption line (E+A or k+a) would follow in the post star-burst phase (\\cite{bia99}). If the star formation is more gradually truncated ($\\gsim$1Gyr) due to halo gas stripping (strangulation), we would not see any excess of E+A/k+a features (\\cite{balogh99}).  It is also important to investigate the morphologies of the galaxies in these groups. The key question here is whether the transformation of morphologies is driven by the same mechanisms as those responsible for the truncation of star formation (\\cite{bia99}; \\cite{treu03}). Furthermore, stellar mass function of galaxies (see \\S2) in groups compared to other environment will provide information on galaxy-galaxy merger in this hierarchy since the mergers increase the fraction of massive galaxies.  On the contrary, strangulation does not change mass, hence can be distinguished by this test.  \\vspace{-0.2cm} ",
        "conclusions": ""
    },
    "0404/astro-ph0404075_arXiv.txt": {
        "abstract": "I compare an aggressive ground-based gravitational microlensing survey for terrestrial planets to a space-based survey. The Ground-based survey assumes a global network of very wide field-of-view $\\sim 2$m telescopes that monitor fields in the central Galactic bulge. I find that such a space-based survey is $\\sim 100$ times more effective at detecting terrestrial planets in Earth-like orbits. The poor sensitivity of the ground-based surveys to low-mass planets is primarily due to the fact that the main sequence source stars are unresolved in ground-based images of the Galactic bulge. This gives rise to systematic photometry errors that preclude the detection of most of the planetary light curve deviations for low mass planets. ",
        "introduction": "Gravitational microlensing can be used to detect extra-solar planets orbiting distant stars \\citep{mao-pac,gould-loeb} with relatively large photometric signals as long as the angular planetary Einstein Ring radius is not much smaller than the angular size of the source star \\citep{bennett-rhie}. For giant source stars in the Galactic bulge, however, the signals of Earth-mass planets are largely washed out by their large angular size, but Galactic bulge main sequence stars are small enough to allow the detection of planets with masses as low as $0.1 M_\\oplus$.  This fact has led to suggestions that ground-based microlensing surveys might be able to measure the abundance of Earth-like planets \\citep{tytler-exnps,sackett97}. However, the initial estimates of the sensitivity of such a survey neglected the blending of main sequence source stars in the bulge that is illustrated in Fig.~\\ref{fig-image}. The brightness of the source stars was over-estimated, and it was not realized that the actual source stars were blended with several other stars of similar brightness in ground-based images. More realistic estimates of the sensitivity of the proposed ground-based extra-solar planet searches indicated very poor sensitivity to terrestrial planets \\citep{peale-gnd_v_space}, and suggested that only a space-based survey \\citep{gest-sim} could discover a significant number of Earth-like planets.  In this paper, we simulate the most capable ground-based microlensing survey that could reasonably be attempted: a network of three 2-m class wide field-of-view telescopes spanning the Globe in the Southern Hemisphere that are dedicated to the microlensing planet search survey for four years.  \\begin{figure}[!ht] \\plotfiddle{space_gnd_imagesC.ps}{2.60in}{0}{53}{53}{-160}{-108} \\caption{The difference between ground and space-based data for microlensing of a bulge main sequence star is illustrated with images of microlensing event MACHO-96-BLG-5. The two top panels are 50 min. R-band exposures with the CTIO 0.9m telescope taken in 1\" seeing at different microlensing magnifications, and the two images on the bottom have been constructed from HST frames. Ground-based photometry suggests that the lensed source star is bulge ``turn-off'' star 1-2 magnitudes brighter than the top of the bulge main sequence, but the HST image reveals that the actual source star is bulge G-dwarf on the main sequence. \\label{fig-image}}  \\end{figure} ",
        "conclusions": "I have carried out detailed simulations of the most capable ground-based microlensing terrestrial planet search program that could plausibly be attempted. The simulated survey employs three 2m class telescopes spanning the Globe in the Southern Hemisphere which observe the Galactic bulge whenever possible for a period of 4 years. Such a network is about 100 times less sensitive than a space-based survey which might only cost 5 times as much. So, without some dramatic improvement in ground-based crowded field photometry, it will not be possible to conduct a microlensing survey to determine the abundance of terrestrial planets in the inner Galaxy.  For a gravitational census of terrestrial planets, a space-based survey will be required \\citep{gest-sim,gest-spie}. As Fig.~\\ref{fig-n_vs_sep} indicates, such a mission would complete the survey for terrestrial planets that will be started by the Kepler mission \\citep{kepler} which is sensitive to Earth-like planets at separations $\\la 1\\,$AU using the transit method. A space-based microlensing survey would overlap Kepler with sensitivity at $\\sim 1\\,$AU, and extend sensitivity to terrestrial planets to all separations including free-floating planets, which may have been ejected from their parent stars during the planetary formation process."
    },
    "0404/astro-ph0404596_arXiv.txt": {
        "abstract": "The first IBIS galactic plane survey has provided a list of high energy emitting objects above 20 keV; these sources have been detected mostly in the crowded region of the Galactic Centre and partly along the Galactic Plane. In order to validate the detection procedure, to help in the identification process and to study the nature of these IBIS sources, this list has been cross correlated with the data archive of the PDS instrument on BeppoSAX, which operated in a similar energy band and with a similar sensitivity. We discover a number of associations whose detailed analysis will be particularly useful for the survey work. Also, thanks to the imaging capability of IBIS/ISGRI, objects which could not be studied by the PDS due to contamination from nearby sources can now be associated with a definite source or sources. ",
        "introduction": "The IBIS/ISGRI Survey (Bird et al. 2004) contains 123 high energy emitting objects detected with the unprecedented sensitivity of $\\sim$1 mCrab in the energy range 20-100 keV discovered by mosaicing all core program observations performed in the first year of the mission; this first catalogue contains 23 high mass and 53 low mass binary systems, 5 AGN, a few SNR/X-ray pulsar systems, a few isolated pulsars and a handful of other objects. Around 30 remain at the moment unidentified and are the main targets of follow up observations. Observations of these and other sources in the IBIS survey at X-ray wavelengths are useful in order to assess their nature and overall characteristics; in this sense the BeppoSAX/PDS archive is a powerful tool as it can provide information on any spectral and/or flux variation provided that the PDS observation was not contaminated by nearby sources. To this end, we have started a program to analyse all PDS observations which contain in the field of view a source detected by IBIS: the systematic search of the archive has provided a set of 68 objects which were targets of BeppoSAX observations and so have both MECS/PDS (2-100 keV band) data. In this case it is possible to reconsider the PDS data in the light of the IBIS images in order to exclude or evaluate any contamination present.  The PDS field of view is 1.3$^{\\circ}$ (FWHM), hexagonal in shape and with no imaging capability (Frontera et al. 1997). The response matrix of this instrument is triangular in both directions, with a flat top of 3' and a reduction in sensitivity of a factor of 2 at 38' from the centre up to zero response at 78' distance . The MECS has instead a field of view of 30' radius and so covers only about 10\\% of the PDS view. Therefore it is also possible that a source is not seen by the MECS but serendipitously observed by the PDS.  In fact, from our search, 21 IBIS survey sources were measured by the PDS in this mode.  For the remaining sources there is no BeppoSAX pointing.  Herein we present a few interesting cases where IBIS/PDS data can be used together in order to define the high energy characteristics of some of these objects which were serendipitously detected by Beppo/SAX. ",
        "conclusions": "In this paper we have demonstrated that the combination of BeppoSAX/PDS and ISGRI data can produce a clearer understanding of the nature of high energy emitters. Work is continuing on this correlation analysis and will improve as the IBIS survey dataset increases."
    },
    "0404/astro-ph0404243_arXiv.txt": {
        "abstract": "{High angular resolution Australia Telescope Compact Array (ATCA) observations of SiO `thermal' millimetre line emission towards the two oxygen-rich, low mass loss rate AGB stars  \\object{R Dor} and  \\object{L$^2$ Pup} are presented. In both cases the emission is resolved with an overall spherical symmetry. Detailed radiative transfer modelling of the SiO line emission has been performed, and the comparison between observations and models are conducted in the visibility plane, maximizing the sensitivity. The excitation analysis suggests that the abundance of SiO is as high as $4\\times 10^{-5}$ in the inner part of the wind, close to the predicted values from stellar atmosphere models.  Beyond a radius of $\\approx 1\\times 10^{15}$ cm the SiO abundance is significantly  lower, about  $3\\times 10^{-6}$, until it decreases strongly at a radius of about   $3\\times 10^{15}$ cm. This is consistent with a scenario where SiO first freezes out onto dust grains, and then eventually becomes photodissociated by the interstellar UV-radiation field. In these low expansion velocity sources  the turbulent broadening of the lines plays an important role in the line formation. Micro-turbulent velocity widths in the range $1.1-1.5$ km\\,s$^{-1}$ result in a very good reproduction of the observed line shapes even if the gas expansion velocity is kept constant. This, combined with the fact that the SiO and CO lines are well fitted using the same gas expansion velocity (to within $5-10$\\%), suggest that the envelope acceleration occurs close to the stellar photosphere, within $\\lesssim 20-30$ stellar radii. ",
        "introduction": "The intense winds that  low to intermediate mass stars develop in their final evolutionary stage, as they ascend the asymptotic giant branch (AGB), return enriched stellar material to the interstellar medium. In addition to significantly contributing to the chemical evolution of galaxies, the mass loss will dictate the time scale for the future evolution of the star towards the planetary nebula phase. Considering its importance little is known with certainty about the mechanism(s) behind the mass loss. Partly this is due to a lack of observational constraints, in particular close to the stellar photosphere where the wind is accelerated, and partly due to the complexity of the physical/chemical processes involved. Among other things, the formation of dust grains is thought to play an important role in these radiatively driven winds. Most of the information on this region comes from observations of infrared ro-vibrational molecular lines in absorption, mainly towards the carbon star \\object{IRC+10216} \\citep[e.g.,][]{Keady93,Winters00a}, infrared continuum emission from the circumstellar dust \\citep[e.g.,][]{Danchi94}, and SiO maser radio line emission from excited vibrational states \\citep[e.g.,][]{Cotton04}.  Information can also be gained from radio observations of suitable circumstellar molecular species towards stars having low to intermediate mass loss rates. A major survey of CO radio line emission from circumstellar envelopes (CSEs) around oxygen-rich AGB stars of different variability types were done by \\citet{Kerschbaum99} and \\citet{Olofsson02}. These data were modelled  in detail to derive stellar mass loss rates and terminal gas expansion velocities \\citep{Olofsson02}. Subsequently, a survey of `thermal' SiO radio line emission, meaning emission from the ground vibrational state rotational lines which are normally not (strongly) masering, was done and the data were interpreted using a detailed numerical radiative transfer modelling presented in \\citet{Delgado03b}.  An immediate conclusion from these large surveys is that the SiO and CO radio line profiles are different from each other. The SiO line profiles are narrower in the sense that the main fraction of the emission comes from a velocity range smaller by about $10-20$\\% than twice the expansion velocity determined from the CO data.  On the other hand, the SiO line profiles have weak wings, such that the total velocity width of its emission is very similar to that of the CO emission. Furthermore, it appears that the SiO line profiles change character with the mass loss rate, at low mass loss rates they are narrow with weak extended wings, while at high mass loss rates they become distinctly triangular.  These `peculiar' SiO line profiles have been interpreted as being due to the influence of gas acceleration in the region which produces most of the SiO line emission \\citep{Bujarrabal86}. However, as illustrated in \\citet{Delgado03b}, the SiO lines are usually strongly self-absorbed also for low mass loss rate objects and this produces narrower lines. There is some controversy in the literature over the scale length of the acceleration region. In the only previously published works where thermal SiO emission have been observed towards AGB stars using interferometry, \\citet{Sahai93} find no need for the slowly varying velocity fields introduced by \\citet{Lucas92}.  Other conclusions come from the detailed radiative transfer modelling of the SiO line data. This is in many respects a more difficult enterprise than the CO line modelling.  The SiO line emission predominantly comes from a region closer to the star than does the CO line emission, and this is a region where we have fewer observational constraints.  The SiO excitation is also normally far from thermal equilibrium with the gas kinetic temperature, and radiative excitation plays a major role.  Finally, there exists no detailed chemical model for calculating the radial SiO abundance distribution. \\citet{Delgado03b}  adopted the assumption that the gas-phase SiO abundance stays high only very close to the star, since further out the SiO molecules are adsorbed onto the grains.  Beyond this the abundance stays low until the molecules are eventually dissociated by the interstellar UV radiation.  This photodissociation radius, which is crucial to the modelling, was estimated using both SiO multi-line modelling and existing interferometer data \\citep{Lucas92, Sahai93}, but only for a few sources. The result of the radiative transfer modelling is a circumstellar SiO abundance that is roughly the same as that obtained from stellar atmosphere equilibrium chemistry for low mass loss rate objects, and which declines with mass loss rate reaching an abundance about two orders of magnitude lower for high mass loss rate objects.  Thus, there are strong indications that `thermal' SiO radio line emission is a useful probe of the formation and evolution of dust grains in a CSE, as well as of its dynamics. Hence, circumstellar SiO line emission potentially carries information on the properties of the region where the mass loss of AGB stars is initiated.  Presented here are results from high spatial resolution imaging of the two M-type semiregular AGB stars \\object{R Dor} and  \\object{L$^2$ Pup}. Both sources studied have low mass loss rates ($\\dot{M}\\lesssim 1\\times 10^{-7}$ M$_\\odot$\\,yr$^{-1}$) and low terminal velocities of their winds ($\\lesssim6$ km\\,s$^{-1}$). It has been suggested from hydrodynamical calculations \\citep{Winters00b,Winters02,Winters03} that the main driving mechanism behind these tenuous winds is stellar pulsation and that dust plays only a secondary role.  The observations, performed with the Australia Telescope Compact Array\\footnote{The Australia Telescope is funded by the Commonwealth of Australia for operation as a National Facility managed by CSIRO.} (ATCA), are presented in Sect.~2. ATCA is an array of six 22 m dishes operating from $1.4-26$ GHz and with an upgrade to $85-105$ GHz operation in progress.  Its location makes it a unique instrument to study molecular-line sources in the southern hemisphere at high angular resolution, and we have recently used it to conduct a study of circumstellar HCN emission from the carbon star R Scl \\citep{Wong04}. The analysis in Sect.~3 and comparison with envelope models in Sect.~4 are carried out in the $uv$-plane in order to maximize the sensitivity and resolution of the data. The modelling is followed by a discussion in Sect.~5 and conclusions are presented in Sect.~6.   \\begin{figure} \\centerline{\\includegraphics[width=7cm]{fig1.ps}} \\caption{Coverage of the visibility plane  for \\object{R Dor} and \\object{L$^2$ Pup} obtained after combining the three ATCA configurations used (EW367, 750B, and H214).} \\label{uvcov} \\end{figure} ",
        "conclusions": "High resolution, interferometric, millimetre observations at $\\approx 1\\arcsec$ resolution of SiO $v=0, J=2\\rightarrow 1$ line emission towards the two O-rich AGB stars \\object{L$^2$ Pup} and \\object{R Dor} have been performed. The emission is resolved, very centrally peaked, and suggests an overall spherical symmetry, even though there is an indication of a departure from spherical symmetry in the case of \\object{R Dor} at the arcsecond scale. A detailed excitation analysis was performed suggesting that the SiO abundance is very high ($\\approx 4\\times10^{-5}$) in the inner part ($\\lesssim 1\\times 10^{15}$ cm) of the circumstellar envelopes around both stars, consistent with predictions from LTE stellar atmosphere chemistry. For  \\object{R Dor} the interferometer data further suggests that there is an additional, more extended, component with a significantly lower SiO abundance. We interpret this as the result of effective adsorption of SiO onto dust grains. Such a conclusion is less clear in the case of \\object{L$^2$ Pup}.  A comparison of model and observed line profiles further suggests that micro-turbulent motions are of the order of $1-1.5$ km\\,s$^{-1}$. This is a significant fraction of the wind velocity in these slowly expanding winds ($\\lesssim 5$ km\\,s$^{-1}$), and it produces strong self absorption of the blue-shifted emission giving the SiO line profiles their characteristic shape. Additionally, the interferometer data provide constraints on the size of the region in which the wind is being accelerated. It is found that a constant velocity over the SiO emitting region can successfully account for the visibility amplitudes obtained at different velocity intervals. Furthermore, it is possible to model the line shapes of both SiO and CO emission using the same expansion velocity (within $5-10$\\%) for both our sources. This constrains the acceleration region to within $\\lesssim 20-30$ stellar radii.  We conclude that SiO line emission plays a very important role in the study of circumstellar envelopes, both with respect to its dynamics and the gas and grain chemistry. However, it should be remembered that both sources studied here have very low mass loss rates and expansion velocities and that a more representative sample of O-rich sources needs to be studied in order to solidify the conclusion reached here. This could also be used to investigate the role of chromospheric radiation in photodissociation of SiO molecules. Such an effort is currently underway, and it justifies the development of a more elaborate model to describe the abundance distribution of SiO."
    },
    "0404/astro-ph0404419_arXiv.txt": {
        "abstract": "Recent fits of cosmological parameters by the first year Wilkinson Microwave Anisotropy Probe (WMAP) measurement seem to favor a primordial scalar spectrum with a large varying index from blue to red. We use the inflationary flow equations to reconstruct large running-index inflaton potentials and comment on current status on the inflationary flow. We find previous negligence of higher order slow rolling contributions when using the flow equations would lead to  unprecise results.  \\\\ PACS number(s): 98.80.Cq, 11.10.Kk ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404133_arXiv.txt": {
        "abstract": "We present the first edition of the \\underline{\\bf S}DSS \\underline{\\bf H}{\\sc ii}-galaxies with \\underline{\\bf O}xygen abundances \\underline{\\bf C}atalog (SHOC), which is a listing of strong emission-line galaxies (ELGs) from the Sloan Digital Sky Survey (SDSS). Oxygen abundances have been obtained with the classic $T_{\\rm e}$-method. We describe the method exploiting the SDSS database to construct this sample. The selection procedures are described and discussed in detail, as well as some problems encountered in the process of deriving reliable emission line parameters. The method was applied to the SDSS Data Release 1 (DR1). We present 612 SDSS emission-line galaxies (624 separate SDSS targets in total), for which the oxygen abundances 12+$\\log$(O/H) have r.m.s. uncertainties $\\le\\,$0.20 dex. The subsample of 263 ELGs (272 separate SDSS targets) have an uncertainty $\\le\\,$0.10 dex, while 459 ELGs (470 separate SDSS targets) have an uncertainty $\\le$0.15 dex. The catalog includes the main parameters of all selected ELGs, the intensities and equivalent widths of hydrogen and oxygen emission lines, as well as oxygen abundances with their uncertainties. The information on the presence of Wolf-Rayet blue and/or red bumps in 109 galaxies is also included. With the use of combined $g,r,i$ SDSS images we performed visual morphological classification of all SHOC galaxies. 461 galaxies ($\\sim$75\\%) are classified as confident or probable blue compact galaxies (BCG/BCG?), 78 as irregular ones, 20 as low surface brightness galaxies (LSBG), 10 as obviously interacting and 43 as spiral galaxies. In creating the catalog, 30 narrow line AGN and 69 LINERs were also identified; these are also presented apart of the main catalog. We outline briefly the content of the catalog, and the prospects of its use for statistical studies of the star formation and chemical evolution issues. Some of these studies will be presented in the forthcoming paper. Finally, we show that the method presented by \\citet{Kniazev03a} for calculating O$^+$/H$^+$ using intensities of the [O\\,{\\sc ii}] $\\lambda$7320,7330 \\AA\\ lines for SDSS emission-line spectra in the absence of [O\\,{\\sc ii}] $\\lambda$3727 \\AA\\ line appears to yield reliable results over a wide range of studied oxygen abundances: $7.10 < 12+\\log$(O/H) $ < 8.5$. ",
        "introduction": "The heavy element abundances of gas-rich galaxies and their gas-mass fractions are the main parameters characterizing their global evolution \\citep[ e.g.,][]{Pagel97,Matteucci01}. These can be related to both the global properties of galaxies \\citep[e.g.,][]{Dalcanton97,GGH03} and the membership of galaxies in various elements of large-scale structure \\citep[e.g.,][]{Popescu96,Grogin00,Pustilnik02b,Vilchez03,Lee03}. The knowledge of the metallicity for large homogeneously selected galaxy samples would allow us to address various issues of galaxy chemical evolution on a good statistical basis. In particular, the possible difference in the chemical evolution rate in the various types of galaxy environments can be systematically examined. Having metallicities for galaxy samples at redshifts of, e.g., $z \\sim 0.3$ and in the nearby Universe, one can directly probe the chemical evolution of gas-rich galaxies over timescales of several Gyr. In addition, estimates of stellar mass from galaxy photometry and of neutral gas mass from H{\\sc i} measurements bring new opportunities to confront the predictions of modern chemical evolution models with the observed properties of a large sample of galaxies.  To address the question of metallicity distributions in gas-rich galaxies and to understand possible relations between metallicity and other galaxy parameters, one can use several large ELG samples such as the results of University of Michigan (UM), Tololo and Calan-Tololo \\citep{Smith76,McAlpine77,Salzer89}, Case \\citep{Pesch82,Salzer95,Ugryumov98}, Second Byurakan Survey  \\citep[SBS;][]{Markarian83,Izotov93}, Heidelberg Void \\citep{Popescu96}, KPNO International Spectral Survey \\citep[KISS;][]{Salzer00,Melbourne02}, Hamburg/SAO Survey for Emission-Line Galaxies \\citep[HSS-ELG;][]{Ugryumov99,Pustilnik00} and Hamburg/SAO Survey for Low Metallicity Galaxies \\citep[HSS-LM;][]{Ugryumov03}.  Of special interest for such samples are H{\\sc ii} galaxies and their most prominent representatives -- Blue Compact Galaxies (BCGs). BCGs are gas-rich objects with typical total masses lower than 10$^{10}~M$\\sunn, have low metallicities in the range 1/15$ \\leq$ Z/Z\\sunn\\ $\\leq$ 1/3, and form stars at noticeably non-stationary rates. Previously, samples of H{\\sc ii} galaxies with reliably known metallicities (i.e., derived with the $T_{\\rm e}$ method) were obtained through high S/N spectroscopy of strong-line ELGs, selected from the surveys cited above. However, well-selected gas-rich galaxy samples with reliable metallicity determinations are still quite small. Currently, it is possible to estimate that we have no more than $\\sim$200 galaxies selected from the different samples with measured classic $T_{\\rm e}$ method oxygen abundances with an accuracy better than 0.1 dex. This is related to the weakness of the key temperature-sensitive line [O\\,{\\sc iii}]$\\lambda$4363, used in the classic $T_{\\rm e}$ method to derive oxygen abundances with r.m.s. uncertainties of 0.01--0.1 dex. As well, many galaxies from these samples often have poor photometry, and, apart from the KISS survey, selection criteria are not well-defined in terms of apparent magnitude. Nevertheless, the accumulated data on low-mass galaxies give important clues about the metallicity distribution and indicate correlations with other galaxy global parameters.  In particular, these surveys have uncovered a significant number of extremely metal-poor galaxies (XMPGs)\\footnote{Another name for these galaxies is extremely metal-deficient galaxies (XMDs).} with $Z \\leq 1/20~Z$\\sunn. Some XMPGs are similar to the well-known I~Zw~18 \\citep{SS72} and SBS 0335--052 \\citep{Izotov90}, which are candidates for young galaxies in the nearby Universe. These are probably the best local analogs of young low-mass galaxies which formed at high redshifts. Despite the paucity of such galaxies their systematic study can advance significantly the understanding of the details of galaxy formation and their early evolution. Therefore, it is important to have an effective means of enlarging substantially the number of XMPGs.  Besides the great interest in understanding the details of star formation, massive-star (including Wolf-Rayet stars) evolution and their interaction with the interstellar medium at very low metallicities, there are several other important directions related to the studies of H{\\sc ii}/BCG metallicities in general. For example, with a larger ELG sample with abundances measured by the classic $T_{\\rm e}$ method we can improve the calibration of the empirical methods \\citep[e.g.,][]{Pagel,McGaugh91,Pilyugin01,Pilyugin03,Denicolo}, which provide a broad picture about the range of oxygen abundances for ELGs in general.  Therefore, it is natural to look for new opportunities offered by the Sloan Digital Sky Survey \\citep[SDSS;][]{York2000}. Owing to its homogeneity, area coverage, spectral resolution, and depth, the SDSS provides an excellent means of creating a large flux-limited sample of H{\\sc ii} galaxies with heavy element abundances derived with the classic $T_{\\rm e}$ method.  The SDSS consists of an imaging survey in five photometric bands \\citep{SDSS_phot,Gunn98,Hogg01}, as well as a follow-up spectroscopic survey of a magnitude-limited sample of galaxies \\citep[mainly field galaxies brighter than $r = 17\\fm77$;][]{Strauss02} and a color-selected sample of quasars \\citep{QSO02}. An automated image-processing system detects astronomical sources and measures their photometric and astrometric properties \\citep{Lupton01,EDR02,SDSS_phot1,Pier03} and identifies candidate galaxies and quasars for multi-fibre spectroscopy.  The samples of galaxy and quasar candidates include a substantial number of emission-line galaxies. The spectra are automatically reduced and wavelength- and flux-calibrated \\citep{EDR02,DR1}.  The SDSS spectral data have already been used in a number of galaxy studies \\citep[e.g.,][]{Bernardi03,Eisen02,Goto03,Kauffmann03,Kniazev03a,Stas03}. The paper presented here will extend the possibilities of using SDSS ELG spectra for the statistical studies of galaxy metallicities.  In this paper we describe the method used to extract from the SDSS database strong-line ELGs, which are suitable for determining oxygen abundances with the classic $T_{\\rm e}$ method, i.e., the temperature-sensitive [O$\\;${\\sc iii}] $\\lambda$4363~\\AA\\ line. The original galaxy sample is obtained from the SDSS DR1, which is briefly described in Section~\\ref{txt:Selection}. The application of the developed pipeline yields the list of ELGs with H{\\sc ii}-type spectra. In the same section the procedure of the ELG selection is described in detail. The method used to estimate the physical conditions in H{\\sc ii} regions of the galaxies studied and their element abundances is described in Section~\\ref{Method}. In the same section we outline a number of problems encountered while using the SDSS spectral data and the ways to resolve them. In Section~\\ref{Tests_OH} we check the quality of the oxygen abundance determination. The catalog of all selected ELGs along with their main parameters, emission line data, and the derived oxygen abundances is presented in Section~\\ref{Results}. The results are presented in Section~\\ref{Discussion}. We conclude with the key results of this paper in Section~\\ref{Summary}. We adopt here the Hubble constant $H_\\mathrm{0}$ = 75~\\kms~Mpc$^{-1}$. ",
        "conclusions": " \\begin{itemize} \\item SDSS spectra permit accurate oxygen abundance determinations over the range $7.1 \\la 12+\\log({\\rm O/H}) \\la 8.5$. \\item The method for calculating O$^+$/H$^+$ using intensities of the [O\\,{\\sc ii}] $\\lambda$7320,7330 \\AA\\ lines appears to yield reliable results over a wide range of oxygen abundances. \\item A large number of strong-line ELGs with measurable oxygen abundances and detectable WR populations is selected from the SDSS DR1 database. \\item A large majority of strong-line ELGs with detected [O~{\\sc iii}]$\\lambda$4363 are H{\\sc ii} galaxies with a broad range of $r$-band luminosities, corresponding to absolute $r$ magnitudes $-22 \\la M_r \\la -12$. \\end{itemize}  We plan to produce regular updates of the SHOC catalog of strong-line ELGs with measured oxygen abundances, based on subsequent Data Releases from the SDSS."
    },
    "0404/astro-ph0404305_arXiv.txt": {
        "abstract": "In this talk I discuss the role of proto-globular clusters as the dominant sources of radiation that reionised hydrogen in the intergalactic medium (IGM) at redshift $z \\sim 6$.  Observations at lower redshift indicate that only a small fraction, \\fesc, of hydrogen ionising radiation emitted from massive stars can escape unabsorbed by the galaxy into the IGM.  High redshift galaxies are expected to be more compact and gas rich than present day galaxies, consequently \\fesc from their disks or spheroids might have been very small. But if the sites of star formation in the galaxies are off-centre and if the star formation efficiency of the proto-clusters is high, then the mean \\fesc calculated for these objects only, is expected to be close to unity.  Here I argue that this mode of star formation is consistent with several models for globular clusters formation.  Using simple arguments based on the observed number of globular cluster systems in the local universe and assuming that the oldest globular clusters formed before reionisation and had \\fesc$\\sim 1$, I show that they produced enough ionising photons to reionise the IGM at $z \\sim 6$. I also emphasise that globular cluster formation might have been the dominant mode of star formation at redshifts from 6 to 12. ",
        "introduction": "In this talk, using simple arguments, I emphasise the important cosmological role of globular cluster (GC) formation at high-redshift. I show that the formation of GCs may have been the dominant mode of star formation near the epoch of reionisation and have contributed significantly to it. The material presented in this talk is based on published work by \\cite{RicottiS:00} and \\cite{Ricotti:02}.  Observation of Ly$\\alpha$ absorption systems toward high-redshift quasars \\citep{Becker:01} indicate that the redshift of reionisation of the intergalactic medium (IGM) is $z_{\\rm rei} \\sim 6$. The recent result from the WMAP satellite \\citep{Kogut:03} of an early epoch of reionisation will not be addressed in this talk. The reader can refer to \\cite{RicottiOI:03, RicottiO:03} if interested in this topic.  A key ingredient in determining the effectiveness by which galaxies photoionise the surrounding IGM is the parameter \\fesc, defined here as the mean fraction of ionising photons escaping from galaxy halos into the IGM.  Cosmological simulations and semi-analytical models of IGM reionisation by stellar sources find that, in order to reionise the IGM by $z =6 - 7$, the escape fraction from galaxies must be relatively large: \\fesc$\\simgt 10$\\% assuming a Salpeter initial mass function (IMF) and the standard $\\Lambda$CDM cosmological model.  The assumption of a universal star formation efficiency (SFE) is consistent with the observed values of the star formation rate (SFR) at $0<z<5$ and total star budget at $z =0$. However, the assumption of a constant \\fesc does not seem to be consistent with observations.  An escape fraction \\fesc$\\sim 1$ is required for reionisation at $z \\sim 6$ but the ionising background at $z \\sim 3$ is consistent with \\fesc$\\simlt 10$\\% \\citep{Bianchi:01}.  Small values of \\fesc at $z \\simlt 3$ are also supported by direct observations of the Lyc emission from Lyman-break and starburst galaxies.  \\cite{Giallongo:02} find an upper limit \\fesc$< 16$\\% at $z \\sim 3$ \\citep[but see][]{Steidel:01} and observations of low-redshift starbursts are consistent with \\fesc upper limits ranging from a few percent up to 10\\% \\citep{Hurwitz:97,Deharveng:01}.  Theoretical models \\citep[\\eg,][]{Dove:00} for the radiative transfer of ionising radiation through the disk layer of spiral galaxies similar to the Milky Way find \\fesc$\\sim 6$\\%. At high redshift the mean value of \\fesc is expected to decrease almost exponentially with increasing redshift \\citep{RicottiS:00, Wood:00}; at $z > 6$, \\fesc$\\simlt 0.1-1$\\% even assuming star formation rates typical of starburst galaxies.  The majority of photons that escape the halo come from the most luminous OB associations located in the outermost parts of the galaxy.  Indeed \\cite{RicottiS:00} have shown that changing the luminosity function of the OB association and the density distribution of the stars has major effects on \\fesc (see their Figs. 8 and 9).  A star formation mode, in which very luminous OB associations form in the outer parts of galaxy halos, may explain the large \\fesc required for reionisation.  Globular clusters are possible observable relics of such a star formation mode.  Their age is compatible with formation at reionisation or earlier.  Because of their large star density they survived tidal destruction and represent the most luminous tail of the luminosity distribution of old OB associations.  In \\S~\\ref{sec:mods} I show that several models for the formation of proto-GCs imply an \\fesc$\\sim 1$.  I will also show that the total amount of stars in GCs observed today is sufficient to reionise the universe at $z \\sim 6$ if their \\fesc$\\sim 1$. This conclusion is reinforced if the GCs we observe today are only a fraction, $1/f_{di}$, of primordial GCs as a consequence of mass segregation and tidal stripping.  In \\S~\\ref{sec:rev} I briefly review a few observational properties and in \\S~\\ref{sec:mods} formation theories of GCs that motivate the assumption of \\fesc$\\sim 1$; in \\S~\\ref{sec:meth} I discuss the model assumptions in light of GC observations and present the results.  In \\S~\\ref{sec:disc} I present my conclusions. ",
        "conclusions": "\\label{sec:disc}  The observed Lyman break galaxies at $z \\sim 3$ are probably the most luminous starburst galaxies of a population that produced the bulk of the stars in our universe. Their formation epoch corresponds to the assembly of the bulges of spirals and ellipticals. Nevertheless the observed upper limit on \\fesc from Lyman break galaxies, \\fesc$\\simlt 10$\\%, may be insufficient to reionise the IGM according to numerical simulations.  \\cite{Ferguson:02}, using different arguments based on the presence of an older stellar population, also noticed that the radiation emitted from Lyman break galaxies at $z>3$ was insufficient to reionise the IGM assuming a continuous star formation mode.  I propose that GCs during their formation may have produced enough ionising photons to reionise the IGM.  Assuming $f_{di}=2$ (\\ie, during their evolution GCs have lost half of their original mass), I find a stellar mass fraction in GCs, $\\omega^f_{gc} \\approx 0.1$\\%, small compared to the total stellar budget $\\omega^f_* \\sim 10$\\% at $z=0$.  But GCs are about 12-13 Gyr old and, if they formed between $5<z<7$ (in about 0.5 Gyr), the expected total $\\omega^f_*$ formed during this time period is about $\\omega^f_* \\sim 1$\\%, only 10 times larger than $\\omega^f_{gc}$. Assuming $f_{di}=20$, expected from the results of N-body simulations, I find $\\omega^f_{gc} \\approx 1$\\%, suggesting that GC formation is an important mode of star formation at high-redshift.  The star formation mode required to explain the formation of GCs suggests an \\fesc$\\sim 1$ from these objects.  This is because the mean \\fesc is dominated by the most luminous OB associations and GCs are extremely luminous, emitting $S \\sim 3 \\times 10^{53}$ s$^{-1}$ ionising photons in bursts lasting only 4 Myrs. Moreover, according to many models, GCs form in the hot, collisionally-ionised galaxy halo, from which all the ionising radiation emitted can escape into the IGM.  Therefore it is not too surprising, if GCs started forming before $z = 6$, that their expected contribution to reionisation is significant. I find that the number of ionising photons per baryon emitted in a Hubble time at $z=6$ by GCs is ${\\cal N}^{gc}_{ph}=(5.1^{+4.3}_{-3.2})f_{di}/\\Delta t_{gc}>1$, therefore sufficient to reionise the IGM even if we assume $f_{di}=1$. Here, $\\Delta t_{gc} \\sim 0.5-2$ is the period of formation of the bulk of old GCs in Gyrs. Using simple calculations based on Press-Schechter formalism [see Fig.~\\ref{fig:sfr}(right)] I find that, if galaxies have \\fesc$ \\simlt 5$\\%, GC contribution to reionisation is important. If GCs formed by thermal instability in the halo of $T_{vir} \\sim 10^5$ K galaxies (case (iii)), the ionising sources have a large bias (\\ie, they form in rare peaks of the initial density field). Therefore, the mean size of intergalactic \\HII regions before overlap is large and reionisation is inhomogeneous on large scales."
    },
    "0404/hep-th0404168_arXiv.txt": {
        "abstract": "{Many compactifications of higher--dimensional supersymmetric theories have approximate vacuum degeneracy. The associated moduli fields are stabilized by non--perturbative effects which break supersymmetry. We show that at finite temperature the effective potential of the dilaton acquires a negative linear term. This destabilizes all moduli fields at sufficiently high temperature. We compute the corresponding critical temperature which is determined by the scale of supersymmetry breaking, the $\\beta$--function associated with gaugino condensation and the curvature of the K\\\"ahler potential, $T_{\\rm crit} \\sim \\sqrt{m_{3/2} M_\\mathrm{P}}\\, (3/\\beta)^{3/4}\\,K''^{-1/4}$.  For realistic models we find $T_{\\rm crit} \\sim 10^{11}$--$10^{12}\\,$GeV, which provides an upper bound on the temperature of the early universe. In contrast to other cosmological constraints, this upper bound cannot be circumvented by late--time entropy production. }  \\clearpage ",
        "introduction": "Compactifications of higher--dimensional supersymmetric theories generically contain moduli fields, which are related to approximate vacuum degeneracy.  In many models these fields acquire masses through condensation of fermion pairs \\cite{Nilles:1982ik}, which breaks supersymmetry. Generically, gaugino condensation models suffer from the dilaton `run--away' problem \\cite{Dine:1985he}, which can be solved, for example, by multiple gaugino condensates \\cite{Krasnikov:jj} or non--perturbative string corrections \\cite{Shenker:1990uf,Banks:1994sg}.  Moduli play an important role in the effective low energy theory. Their values determine geometry of the compactified space as well as gauge and Yukawa couplings. Their masses, determined by supersymmetry breaking, are much smaller than the compactification scale. Hence, moduli can have important effects at low energies. Cosmologically, they can cause the `moduli problem' \\cite{cfx83, deCarlos:1993jw}, in particular their oscillations may dominate the energy density during nucleosynthesis, which is in conflict with the successful BBN predictions. For an exponentially steep dilaton potential, like the one generated by gaugino condensation, there is also the problem that during the cosmological evolution the dilaton ($S$) may not settle in the shallow minimum at Re $S\\sim 2$, but rather overshoot and run away to infinity \\cite{Brustein:nk}. These problems can be cured in several ways (cf.~\\cite{Dine:2000ds}).  In this paper we shall discuss a new cosmological implication of the dilaton dynamics, the existence of a critical temperature $T_{\\rm crit}$ which represents an upper bound on allowed temperatures in the early universe.  If exceeded, the dilaton will run to the minimum at infinity, which corresponds to the unphysical case of vanishing gauge couplings.  The existence of a critical temperature is a consequence of a negative linear term in the dilaton effective potential which is generated by finite--temperature effects in gauge theories \\cite{Buchmuller:2003is}. This shifts the dilaton field to larger values and leads to smaller gauge couplings at high temperature. As we shall see, this effect eventually destabilizes the dilaton, and subsequently all moduli, at sufficiently high temperatures. In the following we shall calculate the critical temperature $T_{\\rm crit}$ beyond which the physically required minimum at Re $S\\sim 2$ disappears.  There can be additional temperature--dependent contributions to the dilaton effective potential coming from the dilaton coupling to other scalar fields \\cite{Huey:2000jx}. These contributions are model dependent and usually have a destabilizing effect on the dilaton, at least in heterotic string models \\cite{Barreiro:2000pf}. Our results for the critical temperature can therefore be understood as conservative upper bounds on the allowed temperatures in the early universe.  The paper is organized as follows. In Sec.~\\ref{sec:gathT} we review the dependence of the free energy on the gauge coupling in $\\mathrm{SU}(N)$ gauge theories. As we shall see, one--loop corrections already yield the qualitative behaviour of the full theory.  In Sec.~\\ref{sec:Tcrit} we study the dilaton potential at finite temperature and derive the critical temperature $T_{\\rm crit}$ for the most common models of dilaton stabilization. Sec.~\\ref{sec:cosmology} is then devoted to the discussion of cosmological implications, the generality of the obtained results is discussed in Sec.~5, and the appendix gives some details on entropy production in dilaton decays. ",
        "conclusions": "At finite temperature the effective potential of the dilaton acquires locally a negative linear term. As we have seen, this important fact is established beyond perturbation theory by lattice gauge theory results.  As a consequence, at sufficiently high temperatures the dilaton $S$, and subsequently all other moduli fields, are destabilized and the system is driven to the unphysical ground state with vanishing gauge coupling.  We have calculated the corresponding critical temperature $T_{\\rm crit}$ which is larger than the scale of supersymmetry breaking, $M_\\mathrm{SUSY} = \\sqrt{M_{\\rm P} m_{3/2}} = \\mathcal{O}(10^{10}\\,{\\rm GeV})$, but significantly smaller than the scale of gaugino condensation, $\\Lambda = [d\\,\\exp(-3S/(2\\beta))]^{1/3} M_\\mathrm{P} = 10^{13}$--$10^{14}\\,\\mathrm{GeV}$.  This is the main result of our paper.  Our result is based on the well understood thermodynamics of the observable sector. In contrast, the temperature of gaugino condensation can place a bound on the temperature of the early universe only under the additional assumption that the hidden sector is thermalized.  The upper bound on the temperature in the radiation dominated phase of the early universe, $T < T_{\\rm crit} \\sim 10^{11}$--$10^{12}\\,$GeV, has important cosmological implications.  In particular, it severely constrains baryogenesis mechanisms and inflation scenarios. Models requiring or predicting $T>T_\\mathrm{crit}$ are incompatible with dilaton stabilization. In contrast to other cosmological constraints, this upper bound cannot be circumvented by late--time entropy production.   We have also discussed more model dependent cosmological constraints. Even if $T<T_\\mathrm{crit}$, the $S$--modulus problem restricts the allowed temperature of thermal leptogenesis and makes the corresponding upper bound on light neutrino masses more stringent. Furthermore, depending on the assumed coupling between dilaton and inflaton, stronger bounds can apply to the reheating temperature.    Our discussion has been based on the assumption that moduli are stabilized by non--perturbative effects which break supersymmetry. Thus the barrier separating the realistic vacuum from the unphysical one with zero gauge couplings is related to supersymmetry breaking. Recently, an interesting class of string compactifications has been discussed where fluxes lead to moduli stabilization and supersymmetry breaking (see, for example, \\cite{gkp02,kkx03,ggx03}). Realistic, metastable de Sitter vacua also require non--perturbative contributions to the superpotential from instanton effects or gaugino condensation \\cite{kkx03}. It remains to be seen how much fluxes can modify the critical temperature in realistic string compactifications.      \\subsection*"
    },
    "0404/astro-ph0404311_arXiv.txt": {
        "abstract": "The ORPHEUS dark matter detector is operating at our underground laboratory in Bern (70~m.w.e.). The detector relies on measuring the magnetic flux variation produced by weakly interacting massive particles (WIMPs) as they heat micron--sized superheated superconducting tin granules (SSG) and induce superconducting--to--normal phase transitions. In an initial phase, 0.45~kg of tin granules in a segmented detector volume have been used. In this paper a general description of the experimental set-up, overall performance of the detector, and first results are presented. ",
        "introduction": "According to the latest cosmic microwave background measurements \\cite{Spergel} combined with other cosmological observations \\cite{SNIa,SNIa2,2dFGRS,SDSS,lensing}, the matter density of the universe comprises 27\\% (for a Hubble constant $H_{\\rm o} = 71$~km~s$^{-1}$~Mpc$^{-1}$) of the critical density  while the remaining 73\\% is in the form of a cosmological constant of an unknown nature. With only $\\approx$~4\\% of the matter density in the form of baryons, the rest should be of some exotic non--baryonic nature. This non--luminous and non--relativistic matter, known as cold dark matter, constitutes an essential ingredient in theories of structure formation in the universe. From the class of candidates for cold dark matter generically known as weakly interacting massive particles (WIMPs), the lightest supersymmetric particle, the neutralino~\\cite{Jungman}, seems to be the most promising. The upper bound for the mass of the neutralino in the constrained minimal supersymmetric extension of the Standard Model is 500~GeV~\\cite{Ellis}, its lower bound of 46~GeV comes from $e^+ e^-$ collider experiments~\\cite{LEP}.  The ORPHEUS project is a direct search experiment looking for cold dark matter aiming at detecting the small energy deposited by a neutralino after being elastically scattered off a target nucleus~\\cite{Ablanalp}. The detector consists of a homogeneous mixture of superheated superconducting tin granules (SSG) embedded in a dielectric filling material and exposed to an external magnetic field. The type~I superconducting granules are kept slightly below the boundary of their superconducting--to--normal phase transition in a metastable state. The recoil energy released by a particle interacting with a granule causes a temperature increase inversely proportional to the specific heat of the granule. Due to the small value of the specific heat at low temperatures, the deposited energy is enough to make a granule normal--conducting. A 1~keV energy deposition in a granule of 35~$\\mu$m diameter at a bath temperature of 110~mK increases the temperature of the granule by typically 100~mK. The induced flux change ({\\em flip}) due to the disappearance of the Meissner--Ochsenfeld effect is measured by an RLC circuit. A general description of SSG detectors is given in Refs.~\\cite{Drukier,Pretzl,Pretzl2}.  This paper is organised as follows: section 2 describes the experimental setup, the shielding, and the data acquisition system; section 3 describes the data taking and its reduction procedure as well as the detector response to muons, section 4 outlines the main results of this experiment and presents the inferred exclusion plot. Finally, section 5 discusses possible improvements of the detector capabilities and outlook. An appendix briefly describes how the expected rates in our detector are computed. ",
        "conclusions": "Due to the relatively high background of our experiment the excluded region in Fig.~\\ref{exclus} lies above those of other experiments. In the same figure, projected exclusion plots are shown under several assumptions: a) same detector characteristics of the present setup, but a background level of 1~cnt~kg$^{-1}$~d$^{-1}$ (attainable at a deeper underground site and with better shielding), b) 1~cnt~kg$^{-1}$~d$^{-1}$, smaller granules ($\\phi = 10$~$\\mu$m) and a narrower relative spread in the superheating distribution (FWHM = 1~\\%). Narrower superheating distributions enhance the sensitive mass of the detector and improve at the same time the background rejection capability. The use of smaller granules allows smaller energy depositions to be detected which makes the detector also more sensitive to smaller WIMP masses.  A hint that narrower superheating distribution could be obtained, has been shown by us in Ref.~\\cite{cyl} (FWHM $\\approx$ 6~\\%), where a regular array of cylinders, produced with an evaporation method, was used. Also, thermal treatment of the granules with a laser beam has shown to have some effect in reducing the spread of the superheating curve \\cite{spimp}. Detecting single flip signals of smaller granules would require more and smaller pickup coils or a different readout system. The use of superconducting quantum interference devices (SQUID) could offer such a possibility. In the past, already granules of 20$\\mu$m in diameter were measured in a large size prototype  \\cite{squid}."
    },
    "0404/astro-ph0404127_arXiv.txt": {
        "abstract": "Using N-body simulations we show that the process of formation of the brightest cluster galaxy through dissipationless galactic cannibalism can affect the inner cluster dark matter density profile.  In particular, we use as realistic test case the dynamical evolution of the galaxy cluster C0337-2522 at redshift $z=0.59$, hosting in its centre a group of five elliptical galaxies which are likely to be the progenitor of a central giant elliptical.  After the formation of the brightest cluster galaxy, the inner cluster dark matter density profile is significantly flatter (logarithmic slope $0.49 \\lsim \\beta \\lsim 0.90$) than the original cusp ($\\beta=1$), as a consequence of dynamical friction heating of the massive galaxies against the diffuse cluster dark matter. In our simulations we have assumed that the cluster galaxies are made of stars only. We also show that the presence of galactic dark matter haloes can steepen the cluster central density profile.  We conclude that galactic cannibalism could be a viable physical mechanism to reconcile -- at least at the cluster scale -- the flat dark matter haloes inferred observationally in some galaxy clusters with the steep haloes predicted by cosmological simulations. ",
        "introduction": "Cold Dark Matter (CDM) cosmological simulations predict that the inner density profile of dark matter (DM) haloes is characterized by a cusp: $\\rhodm (r)~\\propto~r^{-\\beta}$, with logarithmic slope $\\beta \\sim 1-1.5$ for $r \\to 0$: the exact value of the slope and its universality are at the centre of a lively debate. For example, several authors (Navarro, Frenk \\& White 1996, NFW; Moore et al. 1998; Ghigna et al.\\ 2000; Navarro et al.\\ 2004) claim that the resulting profiles are universal, in the sense that $\\beta$ is independent of the halo mass, while other authors suggest that the inner profile depends significantly on the slope of the power spectrum, so that the higher is the mass of the halo the steeper is the cusp (Subramanian, Cen \\& Ostriker 2000; Ricotti 2003, and references therein).  These predictions have been extensively tested by observations both in galaxies (e.g., Salucci \\& Burkert 2000; van den Bosch et al. 2000; de Blok et al.\\ 2001; Simon et al.\\ 2003) and in clusters of galaxies (e.g., Smith et al. 2001; Sand, Treu \\& Ellis 2002; Kelson et al.\\ 2002; Gavazzi et al. 2003; Lewis, Buote \\& Stocke 2003; Sand et al. 2004). The derived values of $\\beta$ range from $\\sim0$ to $\\sim1.3$. The observational results -- shallow slopes and intrinsic scatter -- could represent a serious challenge to the standard cosmological model, and exotic scenarios have been also proposed to solve this problem (e.g. Spergel \\& Steinhardt 2000).  In any case, the role of baryons must be better understood before we are forced to reject a successful paradigm. In fact, it is well known that on scales of the order of kiloparsecs baryons are important and their evolution may affect substantially the DM distribution (see Binney 2004 for a discussion). For example, in the so-called ``adiabatic contraction'' approximation (Blumenthal et al. 1986; Gnedin et al. 2004), baryon dissipation effectively steepens the inner density profile of the host halo (Mo, Mao \\& White 1998; Kochanek \\& White 2001), thus exacerbating the contrast between theory and observations.   Here, motivated by recent observational studies of galaxy clusters based on a joint lensing and dynamical analysis -- in which Sand et al.  (2002, 2004) report inner logarithmic slopes in the range\\footnote{The uncertainties on $\\beta$ could be higher than those given by Sand et al. (2004) if the ellipticity of the cluster mass distribution is higher than estimated by the authors (cf. Dalal \\& Keeton 2004; Bartelmann \\& Meneghetti 2004).  Additional observational work is currently under way to clarify those issues and measure the distribution of DM inner slopes for larger samples of clusters.}  $\\beta=0.38-0.99$ -- we focus on the cluster-sized DM haloes, which are considered to be less affected by baryon dissipation, with respect to galaxy-sized haloes.  We note that the spiraling in of massive galaxies in clusters could flatten the cluster DM distribution, by heating it through dynamical friction (similar to the galaxy-globular cluster interaction at smaller scales; see Bertin, Liseikina \\& Pegoraro 2003).  Recently, El-Zant et al. (2004), using N-body simulations, found that dynamical friction heating due to galaxies, modeled as rigid baryons clumps, is indeed effective in flattening the inner DM profile in clusters. Similar conclusions, in a different context, were also reached by Ma \\& Boylan-Kolchin (2004), who explored with N-body simulations the effects of the dynamical evolution of (deformable) DM sub-haloes on the DM distribution of their host halo. On the other hand, the mass infall associated with the spiraling in of the galaxies towards the cluster centre deepens the central potential well. This dissipationless contraction, in contrast with dynamical friction heating, can result in shrinking the cluster DM distribution and steepening its inner density profile (see, e.g., Barnes \\& White 1984; Jesseit, Naab \\& Burkert 2002): the final density profile will be the result of the two competing effects (see also Section~4).  Here we investigate the effects of dynamical friction on the cluster DM distribution in a realistic scenario, by simulating the time evolution of the galaxy cluster C0337-2522 at redshift $z=0.59$ (ROSAT Deep Cluster Survey; Rosati et al. 1998). This cluster is characterized by the presence of five bright elliptical galaxies (Es) located within $\\sim30$ kpc from the centre, in projection. Nipoti et al. (2003b, hereafter N03), using N-body simulations, showed that the five galaxies are very likely to merge and form a brightest cluster galaxy (BCG) by $z=0$, in accordance with the predictions of the galactic cannibalism scenario (Ostriker \\& Tremaine 1975; see also Merritt 1983; Dubinski 1998).  In the present work we consider the evolution of the underlying cluster DM distribution, with the aid of additional higher resolution simulations (see also Nipoti et al. 2003c).  In our models the single galaxies, as well as the cluster DM, are represented by N-body distributions: this is particularly important, because tidal stripping could remove mass from the individual galaxies, thus altering the infall process and its effects on the cluster DM halo. ",
        "conclusions": "In this paper we explored, with the aid of N-body simulations, the effects on the inner cluster DM density profile of the formation of a BCG, through dissipationless multiple merging of pre-existing galaxies. The initial conditions are designed to reproduce the observed galaxy cluster C0337-2522 ($z=0.59$). Although we focus on a single well observed system as a case study, we stress that the results of the simulations could be generalized to poor galaxy clusters where the central giant elliptical was formed as a consequence of galactic cannibalism, with negligible dissipation.  We found that the inner cluster DM profile is sensitive to the dynamical evolution of cluster galaxies, and in particular to the formation of a BCG.  In our models, where the initial condition are realistic and the galaxies are deformable, we find final slopes in the range $0.49 \\lsim \\beta \\lsim 0.90$, with average $\\avbeta=0.71$.  The qualitative behaviour of our simulations is in agreement with what found by El-Zant et al. (2004), though they obtained a quite small value of the final inner slope ($\\beta \\simeq 0.35$), which is likely to be due to the use of rigid clumps, combined with a high fraction of luminous (clumped) matter ($M_{\\rm lum}/\\Mcltot \\simeq 0.06$). Dynamical friction heating in a cluster with original inner density profile $\\rhodm \\propto r^{-1}$ can thus produce a flatter final inner slope for the cluster DM. Recent measures of the DM profiles of observed clusters can be interpreted taking into account this mechanism.  As outlined in the Introduction, our results are based on the assumption that cluster galaxies are essentially made of stars. However, in a scenario in which the contribution of galactic DM haloes to galaxy masses is important, the inspiraling galaxies might bring a significant amount of DM in the centre of the cluster.  For instance, if we consider, as extreme case, that the massive galaxies in our simulations are made of DM only, the final DM cusp would be steeper than $r^{-1}$, because the DM infall is dominant over dynamical friction heating (cf. solid curve in Fig.~2).  A similar scenario has been recently considered by Ma \\& Boylan-Kolchin (2004), who studied the interplay between a DM halo and its sub-haloes (see also Zhang et al. 2002).  Ma \\& Boylan-Kolchin show with N-body simulations that the dynamical evolution of sub-haloes, depending on their mass and concentration, can either steepen or flatten the halo cusp. Thus, the determination of the amount and distribution of DM in cluster galaxies is very relevant also to the problem of the cluster cusp.  As a consequence, the discrepancies found between the observationally determined inner DM profiles and the predictions of CDM models must not necessarily represent a failure of CDM, at least at the cluster scale.  Instead, they could reflect our poor understanding of the formation and evolution of galaxies, where the physics of baryons is important.  These processes have to be taken into account when predicting the properties of DM haloes of clusters, on scales of few tens of kiloparsecs or smaller. At the same time, our simulations -- when interpreted in the context of structure formation (i.e., galaxies considered as sub-haloes of DM) -- indicate that is quite remarkable to obtain a universal DM profile as a result of successive aggregations\\footnote{We note that a non--homology effect of successive dissipationless mergings was already reported by Nipoti et al. (2003a).}, due to the requirement that dynamical friction heating, potential well contraction, and mass increase all balance. This detailed balance should be further studied in cosmological simulations."
    },
    "0404/astro-ph0404582_arXiv.txt": {
        "abstract": "We report the discovery of non-thermal pulsed X-ray/soft $\\gamma$-ray emission up to $\\sim 150$ keV from the anomalous X-ray pulsar AXP \\axp\\ located near the centre of supernova remnant Kes 73 using RXTE PCA and HEXTE data. The morphology of the double-peaked pulse profile changes rapidly with energy from 2 keV up to $\\sim 8$ keV, above which the pulse shape remains more or less stable. The pulsed spectrum is very hard, its shape above 10 keV can be described well by a power law with a photon index of $0.94\\pm 0.16$. \\axp\\ is the first AXP for which such very-hard pulsed emission has been detected, which points to an origin in the magnetosphere of a magnetar. We have also derived the total emission spectrum from the Kes73/\\axp\\ complex for the $\\sim$ 2-270 keV energy range using RXTE HEXTE and XMM-Newton PN data. A comparison of the total emission from the complex with the pulsed+DC emission from \\axp\\ as derived from Chandra ACIS CC-mode data \\citep{morii03} leaves little room for emission from Kes 73 at energies near 7 keV or above. This suggests that the HEXTE spectrum above $\\sim$ 10 keV, satisfactorily described by a power law with index $1.47 \\pm 0.05$, is dominated by emission from \\axp. In that case the pulsed fraction for energies above 10 keV would increase from about 25\\%  near  10 keV to 100\\%  near 100 keV. The origin of this DC-component extending up to $\\sim$ 100 keV is probably magnetospheric and could be a manifestation of pulsed emission which is ``on\" for all phases. ",
        "introduction": "The quest whether anomalous X-ray pulsars (AXPs) and soft gamma-ray repeaters (SGRs) are both manisfestations of isolated neutron stars with ultra-strong magnetic fields ($\\sim$ $10^{14}-10^{15}$ G), so-called ``magnetars\" \\citep{thompson96}, has recently been decided. Notably, the first detection of SGR-like bursts from the direction of AXP 1E 1048.1-5937 \\citep{gavriil02}, followed by the detection of SGR-like bursts from AXP 1E 2259+586 \\citep{kaspi03} conclusively unified AXPs and SGRs \\citep[see also][]{gavriil04}. Similarly, \\citet{kulkarni03} showed that the originally classical SGR 0526-66 was recently found to behave like an AXP. This was predicted uniquely by the magnetar model \\citep{thompson96}. For recent reviews on AXPs and SGRs see \\citet{mereghetti02} and \\citet{woods04}.  In this paper we present detailed timing and spectral characteristics of AXP \\axp, located in the centre of the supernova remnant (SNR) Kes 73 (G27.4+0.0), which is at a kinematic distance between 6 and 7.5 kpc \\citep{sanbonmatsu92}. The X-ray pulsations of the source were discovered with ASCA \\citep{vasisht97} and a phase-connected timing solution was published by \\citet{gotthelf02}, who analyzed observations with  the Rossi X-ray Timing Explorer (RXTE) spanning 2 years. A linear ephemeris appeared to be consistent with the pulse periods measured over 15 years with Ginga, ASCA, RXTE and BeppoSAX. The measured constant, long-term spin-down of \\axp\\ as well as the phase-connected timing with ``timing-noise-like\" residuals supported a magnetar identification.  \\citet{morii03} reported detailed results for  \\axp\\ from Chandra ACIS CC-mode observations. They could for the first time discriminate the compact object from the surrounding SNR Kes 73.  Like other AXPs, the phase integrated spectrum (pulsed and DC emission) was well fitted with a power law (photon index 2.0 $\\pm$ 0.3) plus blackbody model (kT$_{\\hbox{\\tiny \\rm BB}}$= 0.44 $\\pm$ 0.02 keV). The photon index is the flattest among AXPs. It should be noted, however, that a two-blackbody fit rendered a similarly good fit. They also reported that the pulse profile is double-peaked, with the second pulse exhibiting a harder spectrum between 3 and 7 keV.  \\citet{molkov04} published recently a source catalog of 28 sources detected by INTEGRAL/IBIS between 18 and 60 keV in a survey of the Sagittarius Arm tangent region. The source in this catalog with the hardest spectrum was identified with  \\axp\\ and SNR Kes 73. Since AXPs are ``known\" to be soft-spectrum sources, the SNR seemed at first glance the most likely counterpart. This identification stimulated us to analyse archival RXTE PCA and HEXTE data to search for a pulsed signature from \\axp, particularly in the hard X-ray range covered by HEXTE. We also analyzed archival XMM-Newton data in an attempt to unravel the contributions from Kes 73 and the pulsed and unpulsed emissions from \\axp, combining the informations from our RXTE PCA/HEXTE and XMM-Newton analyses, and the Chandra results on \\axp\\ published by \\citet{morii03} . ",
        "conclusions": "We have presented the surprising discovery of very hard, photon index 0.94 $\\pm$ 0.16, pulsed hard X-ray emission extending up to energies of $\\sim$ 150 keV from the AXP \\axp, obtained using archival monitoring observations by the PCA and HEXTE aboard RXTE. Such hard, non-thermal emission can only originate within the magnetosphere of a neutron star/magnetar. In fact, the total pulsed X-ray spectrum of \\axp\\ is very reminiscent to that of young spin-down powered pulsars. Particularly the spectrum of the Vela pulsar with a blackbody spectrum at lower X-ray energies and a hard power-law extension to higher energies looks very similar. These hard power-law  spectral extensions with photon indices  of $\\sim 1$ are also found for the young (radio weak/quiet) spin-down powered  pulsars PSR B1509-58, PSR B1846-0258, PSR J1811-1925, PSR J1930+1852, all associated with SNRs , like  \\axp, and the old millisecond pulsars PSR J0218+4232, PSR B1821-24 and PSR B1937+21.  Such hard X-ray emission from AXPs has hardly been considered in the modelling of AXPs. However, \\citet{cheng01} studied the production of high-energy gamma-ray radiation in the outer magnetospheres of AXPs. They argue that due to the strong field of a magnetar, the gamma-ray emission rooted at the polar caps will be quenched. However, far away from the pulsar surface, i.e. in the outer gaps, the gamma-rays could be emitted because the local magnetic field will drop below the critical quantum limit. The production process would be curvature radiation by the acceleration of electrons/positrons, resulting from the collisions between high-energy photons from the outer gap and the soft X-rays originating from the stellar surface. They applied their model calculations also to the case of \\axp, and predict an integral gamma-ray photon flux for energies above 100 MeV of $\\sim$ 8.7 $\\times$ 10$^{-9}$ (d/6 kpc)$^{-2}$ cm$^{-2}$ s$^{-1}$, which is below the EGRET sensitivity but above the sensitivity of GLAST. The very-hard pulsed spectrum we measured with HEXTE might be consistent with this interpretation.  We studied the energy dependence of the pulse profile morphology and confirm the findings by \\citet{morii03}, namely the pulse profile is double-peaked, and the second pulse has a harder spectrum for energies below $\\sim$ 7 keV. However, this trend does not continue above $\\sim 10$ keV. Within the statistics, both pulses remain visible at about the same strengths up to $\\sim$ 150 keV. Significant variations with energy of pulse profile morphologies below $\\sim 8$ keV have been reported earlier for three other AXPs (4U 0142+61, 1RXS 1708-4009 and 1E 2259.1+586) by \\citet{gavriil02a}.  The set of spectra shown in Fig. \\ref{high_energy_spc} summarizes our present knowledge on the contributions from the different components constituting the total emission from the direction of Kes 73 and AXP \\axp, i.e. the pulsed and DC components of \\axp, and the emission from SNR Kes 73. In addition to the total pulsed spectrum (PCA/HEXTE) and the pulsed+DC+Kes-73 spectrum (XMM EPIC PN/HEXTE) derived in this work, we show the INTEGRAL IBIS ISGRI flux values for energies 18--120 keV \\citep{molkov04}, and the Chandra ACIS spectrum \\citep{morii03} for the pulsed+DC spectrum of \\axp. We note that the INTEGRAL flux points are consistent with our total HEXTE spectrum (one should realize that these preliminary INTEGRAL IBIS ISGRI flux values have still systematic uncertainties of $\\sim$ 10\\%).  For energies up to about 7 keV we have information on the relative contributions: 1) the pulsed fraction of \\axp\\ is $\\sim$ 25\\%, 2) there is very little margin for emission from SNR Kes 73 around 7 keV. This suggests that the HEXTE spectrum above 10 keV is dominated by emission from \\axp\\ (pulsed+DC), and above $\\sim 100$  keV the spectrum is totally due to pulsed emission. This would mean that also the DC emission from \\axp\\ extends up to $\\sim 100$  keV! An alternative interpretation is that the relative contribution from Kes 73 increases again above 10 keV due to a new hard component. Although the latter seems unlikely, higher spatial-resolution measurements at energies above 10 keV would be required to exclude this possibility.  If we assume that the contribution from Kes 73 is small or negligible above 10 keV, the challenge is now to explain the hard DC emission from \\axp. Most plausible might be an origin in the magnetosphere, namely for the situation that the pulsed emission is ``on\" for all phases. In the competing polar cap \\citep[e.g. as discussed for AXPs by][] {zhang00} and outer gap \\citep[see e.g.][]{cheng01} scenarios explaining the high-energy emission from spin-down powered pulsars there are geometries  for which this is indeed possible. Allowing for a small contribution from Kes 73, e.g. up to about 30 keV, and knowing  that pulsar spectra depend on pulse phase, detailed model calculations are required to investigate this interpretation further."
    },
    "0404/astro-ph0404061_arXiv.txt": {
        "abstract": "Continuing our census of late-type dwarfs in the Solar Neighbourhood, we present BVRI photometry and optical spectroscopy of 800 mid-type M dwarfs drawn from the NLTT proper motion catalogue. The targets are taken from both our own cross-referencing of the NLTT catalogue and the 2MASS Second Incremental release, and from the revised NLTT compiled by Salim \\& Gould (2003). All are identified as nearby-star candidates based on their location in the (m$_r$, (m$_r$-K$_S$)) diagram. Three hundred stars discussed here have previous astrometric, photometric or spectroscopic observations. We present new BVRI photometry for 101 stars, together with low resolution spectroscopy of a further 400 dwarfs. In total, we find that 241 stars are within 20 parsecs of the Sun, while a further 70 lie within 1$\\sigma$ of our distance limit. \\\\ Combining the present results with previous analyses, we have quantitative observations for 1910 of the 1913 candidates in our NLTT nearby-star samples. Eight hundred and fifteen of those stars have distance estimates of 20 parsecs or less, including 312 additions to the local census. \\\\ With our NLTT follow-up observations essentially complete, we have searched the literature for K and early-type M dwarfs within the sampling volume covered by the 2MASS Second Release. Comparing the resultant 20-parsec census against predicted numbers, derived from the 8-parsec luminosity function, shows an overall deficit of $\\sim20\\%$ for stellar systems and $\\sim35\\%$ for individual stars. Almost all are likely to be fainter than M$_J$=7, and at least half are probably as-yet undiscovered companions of known nearby stars. Our results suggest that there are relatively few missing {\\sl systems} at the lowest luminosities, M$_J > 8.5$. We discuss possible means of identifying the missing stars. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404531_arXiv.txt": {
        "abstract": "Interpretation of the \\ion{Fe}{2}(UV)/\\ion{Mg}{2} emission ratios from quasars has a major cosmological motivation. Both Fe and Mg are produced by short-lived massive stars. In addition, Fe is produced by accreting white dwarf supernovae somewhat after star formation begins. Therefore, we expect that the Fe/Mg ratio will gradually decrease with redshift. We have used data from the Sloan Digital Sky Survey to explore the dependence of the \\ion{Fe}{2}(UV)/\\ion{Mg}{2}  ratio on redshift and on luminosity in the redshift range of $0.75< z< 2.20$, and we have used predictions from our 830-level model for the \\ion{Fe}{2} atom in photoionization calculations to interpret our findings.  We have split the quasars into several groups based upon the value of their \\ion{Fe}{2}(UV)/\\ion{Mg}{2}  emission ratios, and then checked to see how the fraction of quasars in each group varies with the increase of redshift. We next examined the luminosity dependence of the \\ion{Fe}{2}(UV)/\\ion{Mg}{2} ratio,  and we found that beyond a threshold of \\ion{Fe}{2}(UV)/\\ion{Mg}{2}~=~5, and $M_{2500} < -25\\rm\\ mag$, the \\ion{Fe}{2}(UV)/\\ion{Mg}{2} ratio increases with luminosity, as predicted by our model.  We interpret our observed variation of the \\ion{Fe}{2}(UV)/\\ion{Mg}{2} ratio with redshift as a result of the correlation of redshift with luminosity in a magnitude limited quasar sample. ",
        "introduction": "Recently, the number of observational efforts focused on  \\ion{Fe}{2}(UV)/\\ion{Mg}{2} emission ratio measurements has increased due to the potential use of these ratios to trace star formation history.   Both Fe and Mg are produced by short lived massive stars (SN type II), and additional Fe is produced by accreting white dwarf supernovae (SN type Ia) sometime later. Thus, one expects that the Fe/Mg ratio is low at high redshift, and gradually increases with decreasing redshift, with the increase starting at a redshift of about 4, corresponding to an age for the Universe of 1.5 Gyr (Hamann and Ferland 1993; Yoshii et al. 1998), or even as early as a redshift of 6 if star formation in the extreme environment of a quasar central region begins at $z \\ge 10$ (Matteucci \\& Recchi 2001).   Although many independent ways to measure \\ion{Fe}{2}(UV)/\\ion{Mg}{2} ratios observed in the UV to IR range have been developed and applied to different data sets, the common conclusion is that the \\ion{Fe}{2}(UV)/\\ion{Mg}{2} ratio shows a large scatter at all redshifts, and little evolution with redshift. (Kwan \\& Krolik 1981; Wills et al. 1985; Kinney et al. 1991; Kawara et al. 1996; Thompson et al. 1999; Iwamuro et al.  2002; Dietrich et al. 2002; Freudling et al. 2003; Barth et al. 2003; Dietrich et al. 2003; and Maiolino et al. 2003). Furthermore, it is usually assumed that there is a linear dependence between the \\ion{Fe}{2}(UV)/\\ion{Mg}{2} and Fe/Mg ratios because the ionization potentials of Mg (7.65 eV) and Fe (7.87 eV) are nearly the same.  In our previous study, we showed that the \\ion{Fe}{2}(UV)/\\ion{Mg}{2} ratios are more sensitive to other physical properties of the emitting region  than to abundance (Verner et al. 2003; 2004). If the physical conditions are different in the Broad Line Regions (hereafter, BLR) of different quasars, the resulting scatter of \\ion{Fe}{2}(UV)/\\ion{Mg}{2} ratios obscures any dependence on abundance. Thus, it is important to explain the origin of the observed scatter before attempting to derive the Fe/Mg relative abundance in quasars.  Although the prominent emission lines observed in quasars  generally allow us to evaluate the metalicity of galactic nuclei, to link specifically the \\ion{Fe}{2} emission with Fe abundance is not a simple task. Compared to many other ions, the \\ion{Fe}{2}  ion has a very rich spectrum due to its half-filled 3d-shell. As a result the \\ion{Fe}{2}(UV) band from 2200$-$3000\\AA\\  [hereafter, \\ion{Fe}{2}(UV))] can contain hundreds of strong lines. The \\ion{Fe}{2}(UV)/\\ion{Mg}{2} emission ratio is therefore heavily affected by many \\ion{Fe}{2} lines but by only  the two \\ion{Mg}{2} doublet lines at ~2800\\AA. Nevertheless, in order to measure the Fe/Mg abundance ratio, obtaining the \\ion{Fe}{2}(UV)/\\ion{Mg}{2} ratio is unavoidable.  As quasar spectra (line intensities and continuum) are heavily affected by \\ion{Fe}{2} emission, it is possible to learn more about BLRs in quasars by studying how the \\ion{Fe}{2} originates. To achieve such a goal we have constructed an 830-level model for the \\ion{Fe}{2} atom and investigated how \\ion{Fe}{2}(UV)/\\ion{Mg}{2} ratios vary with changes in hydrogen density, microturbulence and abundance. The model also predicts that the \\ion{Fe}{2}(UV)/\\ion{Mg}{2} ratio strongly depends upon the ionizing flux of the central source.  In this Letter we have investigated if there is any dependence of \\ion{Fe}{2}(UV)/\\ion{Mg}{2} ratios with redshift. For the first time our approach combines model predictions with measurements [provided by Iwamuro (2004)] of the \\ion{Fe}{2}(UV)/\\ion{Mg}{2} ratios of quasars in the extended redshift range, $0.75 < z < 2.20$, for quasars in the Sloan Digital Sky Survey (hereafter, SDSS)\\footnote{http ://www.sdss.org/}. ",
        "conclusions": "While the extensive quantitative  comparisons between \\ion{Fe}{2}(UV)/\\ion{Mg}{2} ratio measurements and model predictions  are needed, it is clear that our model predictions are in a general agreement with our observational results. Although it has been suggested to use the \\ion{Fe}{2}(UV)/\\ion{Mg}{2} ratio to trace evolution, the physical processes forming lines in the  \\ion{Fe}{2}(UV) band  means that the \\ion{Fe}{2}(UV)/\\ion{Mg}{2} ratio tells us more about the central ionizing source than about abundances. Thus, we can only convert an \\ion{Fe}{2}(UV)/\\ion{Mg}{2} ratio into an Fe/Mg ratio when  a measurement of the \\ion{Fe}{2}(Opt) band has also been obtained.  Our analysis demonstrates that  the observed \\ion{Fe}{2}  spectra can be produced by a wide range of  ionizing flux regimes. The statistical approach applied to physically distinguished  \\ion{Fe}{2}(UV)/\\ion{Mg}{2} ratio groups will provide a more efficient way to improve our understanding of  BLR in quasars than the universal template approach.  Due to complexity and richness of the \\ion{Fe}{2} spectra in quasars, only the comparison between model predictions and  extensive observations make it possible to provide accurate measurements of changes of physical conditions and abundances in quasars with redshift. It is possible that the change we see is due to not only luminosity but to the hydrogen density changes in the BLR with redshift as well. Wide wavelength coverage in the observations is needed for a successful understanding.   We find an increase in the \\ion{Fe}{2}(UV)/\\ion{Mg}{2} ratio with redshift, while evolutionary models predict a decrease. The \\ion{Fe}{2}(UV)/\\ion{Mg}{2} ratio is not sensitive to abundance changes (see Fig. 4, Verner et al. 2003), but it is strongly affected by luminosity. We have found that the relative numbers of quasars with a high \\ion{Fe}{2}(UV)/\\ion{Mg}{2} ratio increases with luminosity, as predicted by our model.   We conclude that the apparent change in the \\ion{Fe}{2}(UV)/\\ion{Mg}{2} ratio with redshift is simply a result of the correlation of redshift with luminosity in the magnitude limited quasar sample, and is not produced by a change in abundance."
    },
    "0404/astro-ph0404194_arXiv.txt": {
        "abstract": "{We report on a 30\\,ksec {\\em XMM-Newton} observation of the central region of the Cha\\,I star forming cloud. The field includes a substantial fraction of the known pre-main sequence population of Cha\\,I South, including all thirteen known very-low mass H$\\alpha$ emitters. We detect two bona-fide brown dwarfs (spectral types M7.5 and M8) and seven H$\\alpha$ emitting objects near the hydrogen burning mass limit, including six of seven earlier detections by {\\em ROSAT}. Three objects classified as Cha\\,I candidate members according to their NIR photometry are revealed by {\\em XMM-Newton}, providing further evidence for them being truely young stars. A total of $11$ new X-ray sources without known optical/IR counterpart may comprise further as yet unrecognized faint cloud members. Spectral analysis of the X-ray bright stars % shows that previous X-ray studies in Cha\\,I have underestimated the X-ray luminosities, as a result of simplified assumptions on the spectral shape. In particular, the extinction is variable over the field, such that the choice of a uniform value for the column density is inappropriate. We establish that the X-ray saturation level for the late-type stars in Cha\\,I is located near $L_{\\rm x}/L_{\\rm bol} \\sim 10^{-2.5}$, with a possible decline to $L_{\\rm x}/L_{\\rm bol} \\sim 10^{-3}$ for the lowest mass stars. A group of strongly absorbed stars with unusual hard X-ray emission is clustered around HD\\,97048, a HAeBe star and the only confirmed intermediate-mass star in the field. While the X-ray properties of HD\\,97048 are indistinguishable from its lower-mass neighbors, another presumably A-type star (identified as such based on NIR photometry) stands out as the softest X-ray emitter in the whole sample. This suggests that various X-ray emission mechanisms may be at work in intermediate-mass pre-main sequence stars. We find that X-ray luminosity follows a tight correlation with age, effective temperature, and mass. No dramatic changes in these correlations are seen at the substellar boundary, suggesting that the same dynamo mechanism operates in both low-mass stars and brown dwarfs, at least at young ages. The variability of the lowest-mass objects is also similar to that of higher-mass T Tauri stars. X-ray flares are seen on about $1/10$th of the Cha\\,I members in the field. ",
        "introduction": "\\label{sect:intro}  The Chamaeleon cloud complex is one of the most nearby star forming complexes, composed of three major clouds. Their isolated position at high galactic latitude ($b \\approx -15^\\circ$), resulting in both low foreground extinction and low contamination with background objects, makes them attractive targets for the study of the formation of low-mass stars.  Cha\\,I hosts the largest number of known low-mass pre-main sequence (PMS) stars and is one of the best-studied of these regions (see e.g. \\cite{Schwartz77.1}, \\cite{Gauvin92.1}, \\cite{Prusti92.1}, \\cite{Hartigan93.1}). A substantial number of the Cha\\,I cloud members have first been identified as X-ray sources in pointed {\\em ROSAT} observations (\\cite{Feigelson93.1}; henceforth F93), and later -- with help of optical observations -- been confirmed as PMS stars (\\cite{Lawson96.1}; henceforth LFH96). In contrast to most of the previously known Cha\\,I members, the so-called classical T Tauri stars (cTTS), these latter ones belong to the class of weak-line T Tauri stars (wTTS), i.e. they show only weak H$\\alpha$ emission, presumably because accretion has ceased after the dispersal of the circumstellar disk. LFH96 estimate that the number ratio of wTTS to cTTS in Cha\\,I is $\\geq 2$. Both populations of TTS stars are mixed in the Hertzsprung-Russell diagram (HRD), indicating a wide range of disk lifetimes.  In the last few years a large number of faint new candidate members of Cha\\,I have been proposed based on their near-infrared (NIR) colors (\\cite{Cambresy98.1}, \\cite{Oasa99.1}, \\cite{Persi00.1}, \\cite{Kenyon01.1}, \\cite{Gomez01.1}, \\cite{Carpenter02.1}). Several of these candidates have been confirmed to be low-mass stars on basis of their NIR spectra (\\cite{Gomez03.1}).  The masses of the presently known Cha\\,I members reach down into the substellar regime, including 13 very-low mass (VLM) objects at or below the hydrogen burning mass limit discovered in an H$\\alpha$ survey by \\citey{Comeron99.1}. ChaH$\\alpha$\\,1 (spectral type M7.5; \\cite{Comeron00.1}; henceforth CNK00) was the first brown dwarf to be detected in X-rays in a {\\em ROSAT} pointed observation (\\cite{Neuhaeuser98.1}). However, {\\em ROSAT} observations were hampered by low sensitivity and low spatial resolution, and thus unable to constrain the X-ray properties of the VLM H$\\alpha$ sources.  The X-ray emission of VLM stars and brown dwarfs is poorly constrained. In lack of sensitive observations it is unclear how far the solar-stellar connection reaches into the VLM regime. In particular the influence of mass, temperature, and age on the activity of the coolest stars and the substellar objects remains obscure. The only bona-fide field brown dwarf detected in X-rays so far, LP\\,944-20, is of intermediate age ($500$\\,Myrs), and was revealed with {\\em Chandra} only during a flare (\\cite{Rutledge00.1}). A subsequent deep {\\em XMM-Newton} observation was not able to recover the source (\\cite{Martin02.1}) suggesting that this object exhibits only episodic outbursts of activity. The X-ray properties of VLM stars and brown dwarfs on the PMS may be different, because at young ages the atmospheres are hotter, such that more ions are present enhancing the coupling between matter and magnetic field (\\cite{Mohanty02.1}).  Recent deep {\\em XMM-Newton} and {\\em Chandra} observations centered on nearby regions of star formation such as $\\rho$\\,Oph (\\cite{Imanishi01.1}), IC\\,348 (\\cite{Preibisch02.1}) and the Orion Nebular Cluster (\\cite{Feigelson02.1}) started to open up the X-ray window to the brown dwarf regime. However, these observations provided at most a few dozen counts per source and/or concern a poorly characterized population of VLM objects. Cha\\,I is one of the nearest star forming regions with the best-studied group of young brown dwarfs known to date, providing the highest sensitivity yet for the detection of the lowest mass stars and brown dwarfs at young ages.  In this paper we present the {\\em XMM-Newton} observation of the central region of the Cha\\,I South cloud. The observation is described in Sect.~\\ref{sect:obs_and_data}, where we also outline the steps of the data reduction. The nature of the X-ray sources is discussed in Sect.~\\ref{sect:nature}. A detailed spectral analysis is performed for the brighter X-ray sources, and an analysis based on hardness ratios for the fainter X-ray sources (Sect.~\\ref{sect:spec}). In Sect.~\\ref{sect:lcs} we present our variability study. The results on individual stars and groups of stars are summarized in Sect.~\\ref{sect:indiv}. Finally, we provide a comparison with earlier {\\em ROSAT} observations of the same field (Sect.~\\ref{sect:rosat}), and we examine correlations between X-ray emission and stellar parameters (Sect.~\\ref{sect:corr}). A summary of our findings is given in Sect.~\\ref{sect:summary}. ",
        "conclusions": "\\label{sect:summary}  In a 30\\,ksec {\\em XMM-Newton} observation of the Cha\\,I South cloud we detected the intermediate-mass HAeBe star HD\\,97048, all known TTS members except Sz\\,23 (confused with VW\\,Cha) and the IR\\,Nebula, most of the VLM H$\\alpha$ objects near or below the substellar limit, and three photometric Cha\\,I candidates two of which are probably late-type stars and one is likely to be an A-type star. A possibly sub-Jupiter mass object recently proposed as a Cha\\,I candidate by \\citey{Comeron04.2} is in the central part of the {\\em XMM-Newton} field, but not detected.  Our X-ray detection of candidate young stars identified by means of NIR photometry is an important indication for them being true members of the star forming region. We argue on basis of their X-ray properties that some additional new cloud members may be among the unidentified X-ray sources. The nature of these objects as well as that of the detected NIR candidate members will be revealed in optical/IR follow-up studies.  We performed a detailed spectral analysis of the brighter half of the Cha\\,I members in the {\\em XMM-Newton} field. The spectra are described by thermal emission from a hot, optically thin plasma, and a photo-absorption term taking account of interstellar and/or circumstellar extinction. Absorbing column densities derived from {\\em Chandra} and {\\em XMM-Newton} spectra were used by \\citey{Vuong03.1} to constrain the gas-to-dust extinction relation. Their analysis of X-ray data from various star forming regions showed that in $\\rho$\\,Oph the measured $N_{\\rm H,X}$ are systematically lower than expected from $A_{\\rm J}$ assuming the standard value of $N_{\\rm H}/A_{\\rm V} = (1.8-2.2)\\,10^{21}\\,{\\rm cm^{-2}}$ per magnitude; see references in \\citey{Vuong03.1}. Their sample in Cha\\,I comprised only 4 reasonably X-ray bright stars in the northern cloud observed with {\\em Chandra}, that span a small range in optical extinction. We point out that low absorption is not a general characteristic of the Cha\\,I cloud. In fact, the {\\em XMM-Newton} FOV contains some objects with $A_{\\rm V} > 10$\\,mag. But the statistics in their X-ray spectra are too small to derive a well-constrained $N_{\\rm H}$. We consider the spectral fits with column density fixed on the value expected from the canonical $N_{\\rm H} - A_{\\rm V}$ relation more reliable, and feel that the interpretation of the $N_{\\rm H} - A_{\\rm V}$ relation based on present-day X-ray observations demands caution.  The X-ray temperatures of the coronal sources in Cha\\,I are similar to those found by \\citey{Favata03.1} in the Taurus cloud L\\,1551, but lower than those found in recent X-ray observations of IC\\,348 (\\cite{Preibisch02.1}), NGC\\,1333 (\\cite{Getman02.1}), and the Orion Nebula Cloud (\\cite{Feigelson02.1}). This can probably be attributed to the use of different instruments and model assumptions. The latter three studies are based on {\\em Chandra} and adopted a 1-T approach. It is a well-known fact that 1-T models do not represent a valid description of the temperature structure in stellar coronae. Nevertheless, they are often used to describe low-resolution X-ray spectra with low numbers of counts. We find that 1-T models have a tendency to underestimate $L_{\\rm x}$ because they underestimate absorption. For this reason we chose to describe all Cha\\,I members which are bright enough for spectral analysis by a 2-T model. Not surprisingly then our spectral analysis brought forth systematically higher X-ray luminosities as compared to the previous estimates derived from {\\em ROSAT} observations. The X-ray luminosities of the low-mass stars in Cha\\,I have now been re-adjusted, and reveal a saturation level near $10^{-2.5...-3}$ (versus $\\sim 10^{-4}$ suggested by {\\em ROSAT}). On basis of this conclusion is it impractical to engage in a study of long-term variability based on X-ray luminosities. We have, however, no signs for variability exceeding a factor of $\\sim 2$ within the last $11$\\,yrs.  The confirmation of X-ray emission from all but one of the {\\em ROSAT} detected VLM ChaH$\\alpha$ objects provides important support for the reliability of the {\\em ROSAT} source detection process. ChaH$\\alpha$\\,7 is the first M8 brown dwarf in Cha\\,I detected in X-rays. In addition, the higher sensitivity and continuous data stream of {\\em XMM-Newton} has allowed for the first time to examine the spectral characteristics and the time variability of the latest type stars and brown dwarfs in Cha\\,I. We find no evidence for a dramatic change in the X-ray properties (such as X-ray temperature and X-ray luminosity) at the substellar limit. The oldest VLM objects are undetected presumably due to the decline of X-ray luminosity with age, and not to an effect of the atmospheric temperature. In terms of variability the lowest mass Cha\\,I members behave similar to higher-mass TTS in the cloud. In particular they are shown to undergo flares."
    },
    "0404/astro-ph0404477_arXiv.txt": {
        "abstract": "We explore linear redshift distortions in wide angle surveys from the point of view of symmetries. We show that the redshift space two-point correlation function can be expanded into tripolar spherical harmonics of zero total angular momentum $S_{l_1 l_2 l_3}(\\hat x_1, \\hat x_2, \\hat x)$. The coefficients of the expansion $B_{l_1 l_2 l_3}$ are analogous to the $C_l$'s of the angular power spectrum, and express the anisotropy of the redshift space correlation function. Moreover, only a handful of $B_{l_1 l_2 l_3}$ are non-zero: the resulting formulae reveal a hidden simplicity comparable to distant observer limit. The $B_{l_1 l_2 l_3}$ depend on spherical Bessel moments of the power spectrum and $f = \\Omega^{0.6}/b$. In the plane parallel limit, the results of \\cite{Kaiser1987} and \\cite{Hamilton1993} are recovered. The general formalism is used to derive useful new expressions. We present a particularly simple trigonometric polynomial expansion, which is arguably the most compact expression of wide angle redshift distortions. These formulae are suitable to inversion due to the orthogonality of the basis functions. An alternative Legendre polynomial expansion was obtained as well. This can be shown to be equivalent to the results of \\cite{SzalayEtal1998}. The simplicity of the underlying theory will admit similar calculations for higher order statistics as well. ",
        "introduction": "It has been known for decades that the two-point correlation function, or power spectrum, measured in redshift surveys is distorted by the peculiar velocities of galaxies. The anisotropy of the correlation function was demonstrated by \\cite{DavisPeebles1983,Peebles1980}. In a seminal work, \\cite{Kaiser1987} demonstrated that in the plane parallel limit there is a simple transformation between the redshift space and real space density contrast. This  results in an anisotropic enhancement of the power spectrum by $(1+f\\mu^2)^2$, where $\\mu$ is the cosine of the angle between the line of sight and the wave-vector. This simple formula has become the starting point of many extensions, which have used expansion into Legendre polynomials \\citep[e.g.,][]{Hamilton1993,HamiltonCulhane1996}, or numerical methods \\citep{ZaroubiHoffman1996}. Most analyses assume a small opening angle \\citep{ColeEtal1995}, i.e. they stay essentially in the distant observer limit. Others works used a expansion with formally infinite number of coefficients \\citep{HeavensTaylor1995}. Numerous galaxy redshift surveys have been successfully analyzed with such methods, most notably the PSCz \\cite{TadrosEtal1999}, 2dF \\citep{PeacockEtal2001,HawkinsEtal2003, TegmarkEtal2002}, and SDSS \\citep{ZehaviEtal2002}. For a review of methods in the above spirit and the corresponding applications, see \\cite{Hamilton1998}.  To address the needs of wide angle redshift surveys, full treatment of wide angle distortions have been given by \\cite{SzalayEtal1998}, where they identify the coordinates in which the expression of the redshift space two-point correlation function is compact, and most importantly finite. They have also argued, that the power spectrum will necessarily have an infinite expansion, as it arises from the convolution of the density field with a non-compact kernel. They concluded that correlation functions are more convenient for redshift space analyses then power spectra. Their results is suitable and has been used for ``forward'' analyses, such as the Karhunen-Loeve method, in which the correlation function is predicted and contrasted with data. Applications to the SDSS are presented most recently by \\cite{PopeEtal2004}, \\citep[see also][]{TegmarEtal2003a,TegmarEtal2003b}.  In this paper we analyze the symmetries of redshift space distortions. The next section shows that zero angular momentum tripolar functions form a natural basis to expand the redshift space correlation function, and that only a surprisingly small set of expansion coefficients are will be non-zero. In section 3 we present the most important properties of the basis functions, and the connection with the Kaiser-Hamilton limit. Section 4 employs the general theory to obtain compact expressions for the redshift distortions using conveniently chosen variables. In the final section we present discussions, and conclusions. ",
        "conclusions": ""
    },
    "0404/astro-ph0404157_arXiv.txt": {
        "abstract": "s{ The Universe became fully reionized, and observable optically, at a time corresponding to redshift $z \\sim 6.5$, so it is only by studying the HI and molecular absorption lines against higher-redshift, radio-loud sources that one can hope to make detailed studies of the earliest stages of galaxy formation. At present no targets for such studies are known. In these proceedings we describe a survey which is underway to find radio-loud quasars at $z > 6.5$, and present broad-band SEDs of our most promising candidates. } ",
        "introduction": "The epoch of reionization has now been discovered as a protracted period reaching from $z \\sim 20 \\rightarrow 6.5$ \\cite{kog,beck}. However, prior to $z \\sim 6.5$ galaxy formation was already well underway (e.g. [3]). It is essentially impossible to study this `grey age' at optical wavelengths, but great progress can be expected if radio and millimetre telescopes can be targeted on quasars observed within the reionization epoch.  Radio-loud targets allow absorption studies that can probe the evolving neutral and molecular content of the high-$z$ Universe \\cite{carilli}, and radio HI absorption is the {\\it only} way of probing the neutral gas which goes on to form stars. We could begin these studies with current facilities (e.g.\\ the GBT and GMRT), and with the next generation of large radio telescopes, such as the LOFAR and the SKA, we will easily be able reach depths of lower luminosity radio sources and still detect 21~cm absorption. Unfortunately, there are currently no known $z > 6.5$ radio-loud objects.  This is because such objects are rare, $\\ll 1$ per cent of the radio population. Interest in pursuing them was dampened by the claim of a much sharper cut-off in their redshift distribution\\cite{shaver} than earlier work\\cite{dp90} had suggested. Jarvis \\& Rawlings\\cite{jr00} and Jarvis et al.\\cite{Jea01} have re-examined all the evidence concerning this redshift cutoff, obtaining results strongly favouring a fairly gradual decline with redshift. ",
        "conclusions": ""
    },
    "0404/astro-ph0404361_arXiv.txt": {
        "abstract": "We study  the evolution  of a system  composed by a  1.4~\\msun neutron star and a  normal, solar composition star of 2~\\msun  in orbit with a period  of 1 day.   Calculations were  performed employing  the binary hydro code presented in Benvenuto \\& De Vito (2003) that handle the mass  transfer rate in a  fully implicit way. Now  we included the main standard  physical ingredients together  with diffusion processes and  a proper  outer boundary  condition.  We  have assumed  fully non~-~conservative mass transfer episodes.  In order to study the interplay of mass loss episodes and diffusion we considered evolutionary sequences with  and without diffusion in which all Roche  lobe overflows (RLOFs)  produce mass transfer.   Another two sequences  in which thermonuclearly-driven RLOFs are not allowed to  drive mass transfer have been computed with and without diffusion. To our notice, this study represents the first binary evolution calculations in which diffusion is considered.  The system  produces a  helium white dwarf  of $\\sim 0.21$~\\msun  in an orbit with a  period of  $\\sim 4.3$~days for the  four cases.  We find that mass transfer episodes induced by hydrogen thermonuclear flashes drive a  tiny  amount of  mass  transfer.   As  diffusion produces stronger flashes, the amount of hydrogen  - rich matter transferred is slightly higher than in models without diffusion.  We find that diffusion is the main agent  in determining the evolutionary timescale  of low mass white dwarfs even in presence of mass transfer episodes. ",
        "introduction": "\\label{sec:intro}  At present,  it is a well  established fact that low  mass white dwarf (WD)  stars should  be formed  during  the evolution  in close  binary systems  (CBSs).  These  objects are  expected to  have a  helium rich interior  simply because  they have  a  mass below  the threshold  for helium  ignition  of about  $0.45~M_\\odot$.  If  they  were formed  as consequence  of single  star  evolution,  we would  have  to wait  for timescales far  in excess of the  present age of the  Universe to find some of them.  Formation  of  helium WDs  in  CBSs  was  first investigated long ago by Kippenhahn, Kohl  \\& Weigert (1967) and Kippenhahn,  Thomas \\& Weigert (1968). They found that these  objects are formed during the evolution of low mass CBSs and that the cooling evolution is suddenly stopped by thermonuclear flashes  that are able to  swell the star  up to produce further Roche lobe overflows (RLOFs).  Since sometime ago, low mass WDs have been discovered as companions to millisecond pulsars  (MSP). This fact sparkled interest  in helium WDs in order to  investigate the deep physical links  between both members of  a given  pair.    In  particular,  it  represents  an  attractive possibility to infer characteristics of the neutron star behaving as a MSP  by  studying the  WD  in detail.  Studies  devoted  to helium  WD properties are those  of Alberts, et al. (1996);  Althaus \\& Benvenuto (1997); Benvenuto  \\&  Althaus  (1998);  Hansen  \\&  Phinney  (1998a); Driebe, et al. (1998); Driebe, et al. (1999); Sch{\\\"o}nberner, Driebe, \\&   Bl{\\\"o}cker  (2000);  Althaus   \\&  Benvenuto   (2000);  Althaus, Serenelli, \\& Benvenuto (2001abc); Serenelli, et al. (2001); Rohrmann, et al. (2002);  Serenelli, et  al.  (2002). Sarna,  Ergma, \\&  Ger{\\v s}kevit{\\v s}-Antipova  (2000) considered the problem in  the frame of detailed binary evolution  calculations. More recently, Podsiadlowski, Rappaport, \\& Pfahl (2002) have   also  computed   the  evolution   of some  CBS configurations that give rise to the formation of helium WDs. Also, Nelson \\& Rappaport (2003) have explored in detail the evolutionary scenarios of binary systems with initial periods shorter than the bifurcation one leading to the formation of ultra-compact binaries with periods shorter than an hour. On the opposite, in this paper we shall deal with a system with an initial period larger than the bifurcation one leading to wider binaries.  Remarkably, the first WD found as companion of a MSP in globular clusters has been detected by Edmonds et al. (2001). Among recent observations  of low mass WD companion  to MSPs we should quote  those by  van Kerkwijk  et al.  (2000) who  have detect  the WD companion  of  the  binary  MSP  PSR  B1855+09,  whose  mass  is  know accurately  from  measurements of  the  Shapiro  delay  of the  pulsar signal, $M_{WD}= 0.258^{+0.028}_{-0.016}$~\\msun. The orbital period of this binary  MSP is 12.3 days.  More recently, Bassa,  van Kerkwijk \\& Kulkarni (2003) have  found a faint bluish counterpart  for the binary MSP  PSR J02018+4232.   The spectra  confirm that  the companion  is a helium WD and, in spite  that observations are of insufficient quality to  put  a strong  constrain  on the  surface  gravity,  the best  fit indicates  a  low  $\\log{g}$   value  and  hence  low  mass  ($\\approx 0.2$~\\msun).  On the other  hand, independently, Ferraro et al. (2003) and Bassa et  al. (2003) have identified the  optical binary companion to  the MSP  PSR  J1911-5958A, located  in  the halo  of the  galactic globular cluster  NGC 6752. This object  turned out to be  a blue star whose position  in the color-magnitude diagram is  consistent with the cooling  sequence  of   a  low-mass  ($\\approx~0.17-0.20$~\\msun),  low metalicity  helium WD  at the  cluster distance.  This is  the second helium WD with a mass in this range that has been found to orbit a MSP in a galactic globular clusters.  Also, Sigurdsson, et al. (2003) have detected two companions for the  pulsar B 1620-26, one of stellar mass and  one of planetary  mass. The  color and  magnitude of  the stellar companion    indicate   a   WD    of   $0.34\\pm0.04$~\\msun    of   age $4.8~\\times~10^{8}$~y. For previous detections of this kind of objects we refer the reader to the paper by Hansen \\& Phinney (1998b).  From a  theoretical point of view,  it was soon realized  that the key ingredient of WD models is the hydrogen mass fraction in the star. Consequently, this called for a detailed treatment of the outer layers of the  star.  Iben \\&  MacDonald (1985) demonstrated the  relevance of diffusion in the  evolution of intermediate mass CO  WDs while Iben \\& Tutukov (1986) found it to be also important in low mass WDs.  More recently, Althaus, et al.  (2001abc) revisited the problem of the formation of helium  WDs. In doing so, they  mimicked binary evolution by abstracting mass  to a 1~\\msun object on the RGB.  The main goal of these papers was to investigate in detail the role of diffusion during the evolution as a pre-WD object. They allowed gravitational settling, chemical  and thermal  diffusion  to operate.  However,  they did  not considered  the  possibility  of   any  mass  transfer  episode  after detachment from the  RGB.  Perhaps the main result  of Althaus, et al. (2001abc) was the finding that  for models with diffusion there exists a  threshold mass  value  $M_{th}$ above  which  the object  undergoes several  thermonuclear  flashes  in  which  a large  fraction  of  the hydrogen present in  the star is burnt out.  Consequently, as the star enters on the  final cooling track it evolves  fast, reaching very low luminosities  on   a  timescale  comparable   with  the  age   of  the Universe. Quite contrarily,  in models without diffusion, evolutionary timescales  are much  longer, making  it difficult  to  reconcile with observations.  For WDs belonging  to CBSs  in companion with MSP,  WD ages should be  comparable to the  characteristic age of pulsars  $\\tau_{PSR}= P/2\\dot{P}$  (for  a pulsar  of period  $P$ with  period derivative  $\\dot{P}$ that  had an initial  period  $P_0$  such  that  $P_0 << P$  and  braking  index $n=3$). This should  be so, because it is  generally accepted that the MSP is recycled  by accretion from its normal  companion.  However, it was found that  the WD was much dimmer than predicted  by models without diffusion,  which  should  be  interpreted  as consequence  of  a  faster evolution.  This  has been the  case of the companion of PSR B1855+09.  For  objects with  masses  below $M_{th}$ no thermonuclear  flash occurs and the star  does not suffer from another  RLOF.  Consequently, it retains a  thick hydrogen layer, able to support nuclear burning, forcing the WD to remain bright for a very long time. This is the case of the companion of PSR J1012+5307.  It  is the aim  of the  present paper  to revisit  the problem  of the formation  and evolution  of helium  WDs  in CBSs  by performing  full binary computations considering diffusion  starting with models on the main sequence all  the way down to stages of  evolution of the remnant as a very cool  WD. To our notice this is the  first time such a study is carried out. In this way we largely generalize the previous studies from our group  on this  topic.  In doing  so, we  have preferred  to concentrate  on  a  particular  binary system,  deferring  a  detailed exploration  of the  huge  parameter space  (masses, orbital  periods, chemical compositions, etc.)  to  future publications. To be specific, we have chosen  to study  a  CBS composed  by a  2~\\msun normal  star together with a neutron star with a ``canonical'' mass of 1.4~\\msun on an initial orbit  of 1  day of  period. We assumed solar chemical composition with $Z=0.02$   for    which $M_{th}~\\approx~0.19$~\\msun (Althaus et al. 2001b).  In order to  explore the role and interplay of  mass loss episodes and diffusion  we  have   constructed  four complete  evolutionary calculation:  \\begin{itemize} \\item Case A: Diffusion, all RLOF operate (including flash-induced RLOF) \\item  Case B:  Diffusion, without flash-induced RLOF \\item Case C: No diffusion, all RLOF operate \\item Case D: No diffusion, without flash-induced RLOF \\end{itemize}  Regarding mass  transfer episodes we have  chosen to study  the case of fully non conservative conditions, i.e., those in which all the matter transferred from  the primary  star is lost  from the  system carrying away all  its intrinsic angular  momentum. We do  so in order  get the strongest possible RLOFs which, in turn, will produce the largest mass transfer episodes. In  this sense, we shall get an  upper limit to the effects of  RLOFs on  the whole evolution  of the star,  in particular regarding the ages of very cool WDs.  The  reminder  of the  paper  is  organized  as follows.   In  Section \\ref{sec:numerical} we describe our code paying special attention to  the  changes we  implemented  in  the  scheme for  computing  mass transfer episodes. Then, in  Section \\ref{sec:calcula} we describe the evolutionary results for the  four cases considered here.  Finally, in \\ref{sec:discu}  we discuss  the implicances  of our  calculations and summarize the main conclusions of this work. ",
        "conclusions": "\\label{sec:discu}  In  this  work we  have  computed the  evolution  of  a binary  system composed by a neutron star  with a ``canonical mass'' of 1.4~\\msun and a normal, population I main sequence star of 2~\\msun in orbit with a 1 day period.   We have performed the calculations  employing an updated version of the code presented in  Benvenuto \\& De Vito (2003) in which we have included the  main standard physical ingredients together with diffusion  processes.   Also a  proper  outer  boundary condition  was incorporated following Ritter (1988) (see \\S \\ref{sec:numerical}).  In  order to  explore  the role  of  mass transfer  episodes from  the primary star and its interplay  with diffusion we have considered four situations:  diffusion,  all RLOF operate (Case  A); diffusion, no flash~-~induced RLOF operates (Case  B); no diffusion, all RLOF  operate (Case C); and no  diffusion, no flash~-~induced RLOF operates   (Case D).  See  introduction (\\S~\\ref{sec:intro}) of further details.  To our  notice, these calculations represent the  first detailed study of binary  evolution considering diffusion.  In this sense,  this work represents  a  natural  generalization  of the  results  presented  by Althaus, Serenelli \\& Benvenuto  (2001a) in which binary evolution processes was  mimicked by  forcing a 1~\\msun  star on the  red giant branch  to undergo  an appropriate  mass  loss rate.   Now the  proper inclusion of the specific processes that govern binary evolution offer us a more  physically sound description of the  formation of low mass, helium white dwarfs (WDs). In  particular, now we have the possibility of  connecting  stellar  structure  and  evolution  with  the  orbital parameters  of  the systems  allowing  for  a  deeper comparison  with observations.  From the results presented in  the previous sections, it is clear that diffusion  is  far more  important  in  determining  the timescale  of evolution of the stars than  mass transfer episodes during flash~-~induced RLOFs. This is so especially  when the object reaches the  final cooling track. We found  that timescales are  almost insensitive  to the occurrence of flash~-~induced RLOF  episodes (see Fig.~\\ref{fig:lumi}). This constitutes the  main result of the present work."
    },
    "0404/astro-ph0404227_arXiv.txt": {
        "abstract": "We report initial results for spectroscopic observations of candidates of Lyman Break Galaxies (LBGs) at $z\\sim5$ in a region centered on the Hubble Deep Field-North by using the Faint Object Camera and Spectrograph attached to the Subaru Telescope. Eight objects with $I_C\\leq25.0$ mag, including one AGN, are confirmed to be at $4.5<z<5.2$. The rest-frame UV spectra of seven LBGs commonly show no or weak Ly$\\alpha$ emission line (rest-frame equivalent width of $0-10$\\AA) and relatively strong low-ionization interstellar metal absorption lines of \\ion{Si}{2} $\\lambda$1260, \\ion{O}{1}+\\ion{Si}{2} $\\lambda$1303, and \\ion{C}{2} $\\lambda$1334 (mean rest-frame equivalent widths of them are $-1.2 \\sim -5.1 $\\AA). These properties are significantly different from those of the mean rest-frame UV spectrum of LBGs at $z\\sim3$, but are quite similar to those of subgroups of LBGs at $z\\sim3$ with no or weak Ly$\\alpha$ emission. The weakness of Ly$\\alpha$ emission and strong low-ionization interstellar metal absorption lines may indicate that these LBGs at $z\\sim5$ are chemically evolved to some degree and have a dusty environment. Since the fraction of such LBGs at $z\\sim5$ in our sample is larger than that at $z\\sim3$, we may witness some sign of evolution of LBGs from $z\\sim5$ to $z\\sim3$, though the present sample size is very small. It is also possible, however, that the brighter LBGs tend to show no or weak Ly$\\alpha$ emission, because our spectroscopic sample is bright (brighter than $L^{\\ast}$) among LBGs at $z\\sim5$. More observations are required to establish spectroscopic nature of LBGs at $z\\sim5$. ",
        "introduction": "In order to understand formation and evolution of galaxies, systematic searches for galaxies in the early universe and detailed studies of the detected galaxies are necessary. In this decade, a number of galaxies detected at high redshift ($z>2\\sim3$) has been rapidly increasing. The use of photometric redshift selection from deep broad-band images, especially so-called \"Lyman break method\" (e.g., \\citet{Ste92}; Steidel, Pettini \\& Hamilton 1995) gives the largest sample of galaxies at $z\\sim3$ \\citep[Lyman Break Galaxies; LBGs, e.g.,][]{Ste03}. Extensive studies  of them have been revealing individual and statistical nature of star-forming galaxies at $z\\sim3$ (e.g., \\citet{Ste99,Ste98,Pet01}; Papovich, Dickinson, \\& Ferguson 2001).  At the same time, their optical follow-up spectroscopy has also been made extensively \\citep[e.g.,][]{Ste96a,Ste96b,Ste99}, and revealed rest-frame UV spectral features of LBGs in addition to their redshift information. Using about 800 spectra of LBGs at $z\\sim3$, \\citet{Shap03} found that about three fourth of them show significant Ly$\\alpha$ emission line and the remainders show only Ly$\\alpha$ absorption. The LBGs with weaker Ly$\\alpha$ emission have stronger low-ionization interstellar metal absorption lines and redder UV continua. These trend suggest that the LBGs at $z\\sim3$ with weak Ly$\\alpha$ emission are more metal enriched and dusty, i.e., chemically evolved than those with strong Ly$\\alpha$ emission.  How about properties of LBGs at higher redshift? Are there any signs of evolution of LBGs compared with LBGs at $z\\sim3$? To answer the question and obtain clues to understand formation and evolution of galaxies in the early universe, we made a systematic search for LBGs at $z\\sim5$ ($\\sim$1Gyr earlier to $z=3$). We carried out wide (effectively $\\sim600$ arcmin$^2$) and deep broad-band ($V, I_C,$ and $z'$) imaging observations toward an area centered on the Hubble Deep Field-North \\citep[HDF-N;][]{Wil96} with Suprime-Cam \\citep{Miya} attached to the Subaru telescope. Thanks to plenty of redshift information of galaxies in and around the HDF-N, we could set suitable color criteria on the two-color ($V-I_C$ and $I_C-z'$) diagram to effectively select  galaxies at $4.5\\lesssim z \\lesssim5.5$ avoiding foreground contamination, and obtained $\\sim300$ LBG candidates at $z\\sim5$ with $23.5$ mag $<I_C\\leq26.0$ mag \\citep{Iwa03}. This is one of the largest LBG samples at $z\\sim5$ systematically surveyed. Using this sample, we statistically derived a rest-frame UV luminosity function (UVLF) of LBGs at $z\\sim5$, which shows only a slight change from $z\\sim5$ to $z\\sim3$. We also found that a spatial distribution of LBGs at $z\\sim5$ shows a large void-like structure and a larger clustering amplitude than that at $z\\sim3$ (in preparation). The presence of a similar void-like structure is seen in \\citet{Bre04}, and a large clustering amplitude of LBGs at $z\\sim5$ may also suggest the presence of the larger bias in galaxy formation at $z\\sim5$ than at $z\\sim3$ \\citep{Ouchi03c}.  In order to confirm redshifts of the LBG candidates and to study spectral features and relations with other properties of them, we started follow-up optical spectroscopy of our LBG candidates at $z\\sim5$. Although similar systematic searches for $z=5\\sim6$ galaxies have been carried out (e.g., \\citet{Ouchi03b,Bou03}; Yan, Windhorst, \\& Cohen 2003; Stanway, Bunker, \\& McMahon 2003), a number of spectra of LBGs at $z\\sim5$ is  still extremely small. Moreover most of the spectra obtained so far show only Ly$\\alpha$ emission, and the continuum spectral feature and relations with other properties observed in LBGs at $z\\sim3$ are still quite unknown. In this paper, we present initial results of the optical spectroscopy of a part of our sample. Observations and data reduction are described in section 2, and the results are presented in section 3. In section 4, we compare the results  with spectral features obtained for LBGs at $z\\sim3$, and discuss possible origins of the similarity and difference seen in spectral features. The magnitude system is based on \\rm{AB} magnitude unless otherwise noted. ",
        "conclusions": "\\citet{Shap03} divided spectroscopically confirmed $\\sim$800 LBGs at $z\\sim3$  into four subgroups (each group consists of $\\sim$200 LBGs) based on  rest-frame Ly$\\alpha$ equivalent width, and made composite spectra for each group. They found that about three fourth of the $\\sim$ 800 LBGs at $z\\sim3$ show Ly$\\alpha$ emission lines, and about two third of them show strong Ly$\\alpha$ emission (EW$_{\\rm rest} > 10$\\AA) \\footnote{\\citet{Shap03} measured the equivalent width of Ly$\\alpha$ line by summing up both emission and absorption. In our case, we measured the equivalent width of Ly$\\alpha$ only for the emission part mainly because it is hard to distinguish intrinsic absorption and intergalactic absorption. Thus the value of rest-frame equivalent width of Ly$\\alpha$ by \\citet{Shap03} is the lower limit when it is compared with our equivalent width.}. They also found that LBGs with weaker Ly$\\alpha$ emission have stronger LIS absorption lines, redder UV continuum slopes, and larger $E(B-V)$ values. The weak Ly$\\alpha$ emission, strong LIS absorption lines, the red UV color, and the large $E(B-V)$ are considered to originate in dusty environment. In the LBGs with no or weak Ly$\\alpha$ emission, star formation may occur earlier and may be chemically more evolved than those with strong Ly$\\alpha$ emission. Although Shapley et al. (2001) pointed out a possibility that  the LBGs with weak Ly$\\alpha$  emission are younger than those with strong emission from the SED fitting, this might be reconciled with the chemically evolved nature if the ages derived by SED fitting are affected by the most recent star formation occurred in the LBGs.  The spectra of our seven LBGs (expect for one AGN) at $z\\sim5$ and thus the composite spectrum of them is significantly different from the composite spectrum of LBGs at $z\\sim3$ as shown in Figure 2. The Ly$\\alpha$ emission is much weaker, and the continuum depression in the wavelength region shorter than the redshifted Ly$\\alpha$ emission is much larger than those at $z\\sim3$. The measured equivalent width of Ly$\\alpha$ emission is $4.5$\\AA\\ at $z\\sim5$ while $15.1$\\AA\\ at $z\\sim3$. The LIS absorption lines are stronger in the spectrum for $z\\sim5$ than in that for $z\\sim3$; measured equivalent widths of \\ion{Si}{2} $\\lambda$1260, \\ion{O}{1}+\\ion{Si}{2} $\\lambda$1303, and \\ion{C}{2} $\\lambda$1334 are $-2.8$\\AA, $-2.3$\\AA, and $-2.4$\\AA, respectively for $z\\sim5$, while $-1.7$\\AA, $-2.3$\\AA, and $-1.5$\\AA, respectively for $z\\sim3$. However, the spectra of the seven LBGs at $z\\sim5$ fairly resemble to subpopulations of LBGs at $z\\sim3$; the composite spectra of LBGs at $z\\sim3$ with no or weak Ly$\\alpha$ emission \\citep[Group 1 and Group 2 by][]{Shap03} are quite similar to the spectra of our LBGs at $z\\sim5$. The average rest-frame equivalent widths of the three LIS absorption lines of the seven LBGs at $z\\sim5$ is $-2.8$\\AA\\, being very close to $-2.5$\\AA\\ for Group 1 (no Ly$\\alpha$ emission). The average value of EW$_{\\rm rest}$(LIS) corresponds to metallicity of 12$+$log(O/H)$\\sim8.0$, if we assume that the relation obtained in the local universe by \\citet{Hek98} can be applied to the high redshift LBGs, which is not certain at this moment. These results suggest that these LBGs at $z\\sim5$ are chemically evolved to some degree.  All our LBGs confirmed to be at $z\\sim5$ show no or weak Ly$\\alpha$ emission with relatively strong LIS absorption lines; a fraction of LBGs with strong Ly$\\alpha$ emission is very small, though the sample size is still small. We may be witnessing some sign of evolution of LBGs from $z\\sim5$ to $z\\sim3$. The lack of strong Ly$\\alpha$ emission as well as the presence of strong LIS absorption at $z\\sim5$ are likely to be due to their dusty and chemically evolved environment (though the escape of Ly$\\alpha$ photons may not be related in a simple way to the metallicity of the galaxy \\citep[e.g.,][]{kunth}) and to the presence of more neutral hydrogen in and/or around a galaxy than that at $z\\sim3$.  However, it is worth emphasizing that the LBGs we observed are relatively brighter ones among those at $z\\sim5$ ($I_C = 25.0$ mag corresponds to $M^{\\ast}$ of UVLF at $z\\sim5$ \\citep{Iwa03}.). There is a possibility that the strength of Ly$\\alpha$ emission depends on the magnitude (i.e., UV continuum). \\citet{Shap03} found for LBGs at $z\\sim3$ that the average UV magnitude is fainter for the LBGs with stronger Ly$\\alpha$ emission. They also found that the average rest-frame equivalent width of the Ly$\\alpha$ emission line of faint LBGs is larger than that of bright LBGs among the subgroup with strong Ly$\\alpha$ emission of EW$_{\\rm rest} \\geq$ 20\\AA. This trend could also be the case at $z\\sim5$. In fact, \\citet{Lehn02}, who made a similar search for LBGs at $z\\sim5$ using $R$, $I$, and $z$ band deep images, found that all of their spectroscopically confirmed six objects to be at $z\\sim5$ show very strong (EW$_{\\rm rest}>30$\\AA) Ly$\\alpha$ emission lines. The $I$ magnitudes of them are $\\sim1$ magnitude fainter than those of our seven LBGs at $z\\sim5$. In addition, \\citet{Ouchi03a} found that the number density of Ly$\\alpha$ emitters (LAEs) at $z=4.86$ against to LBGs at $z\\sim5$ rapidly decreases with increasing UV continuum light ($i'\\lesssim25$ mag). These observational results suggest that the brighter LBGs tend to show the weaker Ly$\\alpha$ emission also at $z\\sim5$. It is also worth noting here that the LBGs with secure redshifts tend to show rather compact morphology. Thus the weakness of Ly$\\alpha$ emission and the strong LIS absorption lines may also relate to their morphological property, which may also link to an evolutionary stage of galaxies.  If we interpret that the weakness of the Ly$\\alpha$ emission is caused by the dusty environment in the LBGs, it might be possible that the brightest LBGs are the most chemically evolved ones at the epoch. It is known that brighter LBGs have a larger correlation length at $z\\sim3$ \\citep{gd01}, suggesting that they are associated with more massive dark halos and star formation may occur with biased manner in the earlier epoch as compared with less clustered fainter LBGs. The similar result has been obtained for LBGs at $z\\sim4$ by \\citet{Ouchi03c}. We also found the same clustering segregation with magnitude in our sample of LBGs at $z\\sim5$ (in preparation). Thus the brighter LBGs at $z\\sim5$ may be associated with even more massive dark halos and have experienced more biased star formation, resulting in more dusty environment as compared with many of LBGs at $z\\sim3$ with a strong Ly$\\alpha$ emission.  Another possible reason of the difference of Ly$\\alpha$ appearance is the effect of a strong clustering of LAEs. \\citet{Ouchi03a} and \\citet{Shima} found that there is an overdensity region ($\\sim10^{\\prime}\\times10^{\\prime}$) of LAEs (EW$_{\\rm rest} > 14$\\AA) at $z\\sim4.86$ from their wide-field deep narrow-band observations. The observed field (44 arcmin$^2$) of \\citet{Lehn02} may happen to fall on an overdensity region of Ly $\\alpha$ emitters; their rest-frame equivalent widths of Ly$\\alpha$ emission ($>30-50$\\AA) is larger than the detection threshold of the LAE selection. While our field of view for the spectroscopy  (total of $\\sim$85 arcmin$^2$) may happen to point to a low density region of LAEs and we could not observe LBGs with strong Ly$\\alpha$ emission. However, since we selected spectroscopic targets from regions where the surface density of LBG candidates is relatively high, this may be unlikely provided that the distribution of LBGs broadly coincides with that of LAEs at the same epoch.  To summarize,  the results presented here may show some sign of evolution in spectroscopic feature from $z\\sim5$ to $\\sim3$, or the presence of luminosity dependence of nature. However, our sample size is too small to reach any significant conclusions. Further spectroscopic observations of LBGs at $z\\sim5$ over wider field and magnitude range are necessary to reveal spectroscopic nature and discuss relationship with evolution of LBG population."
    },
    "0404/astro-ph0404541_arXiv.txt": {
        "abstract": "We report the independent discovery of a new extrasolar transiting planet around OGLE-TR-113, a candidate star from the Optical Gravitational Lensing Experiment. Small radial-velocity variations have been detected based on observations conducted with the MIKE spectrograph on the Magellan~I (Baade) telescope at the Las Campanas Observatory (Chile) during 2003. We have also carried out a light-curve analysis incorporating new photometry and realistic physical parameters for the star.  OGLE-TR-113b has an orbital period of only 1.43 days, a mass of $1.08\\pm0.28$~M$_{\\rm Jup}$, and a radius of $1.09\\pm0.10$~R$_{\\rm Jup}$. Similar parameters have been obtained very recently in an independent study by Bouchy et al., from observations taken a year later.  The orbital period of OGLE-TR-113b, and also that of the previously announced planet OGLE-TR-56b ($P_{\\rm orb}=1.21$~days) ---the first two found photometrically--- are much shorter than the apparent cutoff of close-in giant planets at 3-4-day periods found from high-precision radial velocities surveys. Along with a third case reported by Bouchy et al.\\ (OGLE-TR-132b, $P_{\\rm orb} = 1.69$~days), these objects appear to form a new class of ``very hot Jupiters'' that pose very interesting questions for theoretical study. ",
        "introduction": "In recent years the field of extrasolar planet research has seen significant developments in the ability to discover and measure these objects using a variety of techniques. High-precision Doppler searches \\citep[e.g.,][]{Fischer:03, Naef:04} have yielded the vast majority of the discoveries, and measurements in at least one case have even been made astrometrically \\citep{Benedict:02}. Transit searches had their first success with HD~209458 \\citep{Henry:00, Charbonneau:00}, a bright star ($V = 7.65$) that was known previously to harbor a planet in a 3.5-day period orbit from its radial-velocity signature. Numerous photometric programs are monitoring large samples of stars looking for small dips in the brightness of the central object at the $\\sim$1\\% level \\citep[see][]{Horne:03}, which might indicate a planet-size object crossing in front of the star.  These studies are very important for the additional information they bring to bear on the nature of the companion, namely, the inclination angle of the orbit ($\\sim$90\\arcdeg) and the absolute radius of the planet.  The inclination angle complements the spectroscopic information and allows a direct determination of the mass.  Dozens of transiting planet candidates among faint stars have already been reported by several teams including OGLE \\citep{Udalski:02a}, EXPLORE \\citep{Mallen-Ornelas:03}, MACHO \\citep{Drake:04}, and others. Multiple efforts are underway to follow-up on these candidates, a necessary step given the high incidence of false positive detections, particularly among fainter stars in crowded fields.  The first case to be confirmed was that of OGLE-TR-56, a star with $V = 16.6$ located in the direction of the Galactic center \\citep{Udalski:02b, Konacki:03a, Torres:04}. The very short orbital period of this planet (only 1.21~days) makes it extremely interesting, and has provided theorists the opportunity to explore the effects of strong irradiation from the central star as well as evaporation \\citep[e.g.,][]{Burrows:03, Baraffe:03, Baraffe:04}.  In this paper we report the detection of a Doppler signature induced by a giant planet orbiting OGLE-TR-113, another faint transit candidate ($I = 14.42$) in the constellation of Carina reported recently by the OGLE project \\citep{Udalski:02c}. This star shows periodic dips in brightness of about 3\\%, and has a photometric period of 1.43~days.  OGLE-TR-113 was originally identified as a very promising candidate from our low-res\\-o\\-lu\\-tion spectroscopic observations conducted in 2002 \\citep[see][]{Konacki:03b}. This reconnaissance showed it to be a star of late spectral type with no obvious velocity variations at the level of a few \\kms, which would have otherwise disqualified it for implying a stellar companion. Subsequently it was placed on our program for high-resolution follow-up, and the observatios were carried out in early 2003. As this paper was being prepared we learned of a very recent independent detection of radial velocity variations in OGLE-TR-113 by \\cite{Bouchy:04}, based on observations taken in 2004. That study found yet another case of a very short-period transiting planet (OGLE-TR-132, $P_{\\rm orb} = 1.69$~d), which brings the number of such objects to three. It appears, therefore, that they form a new class of ``very hot Jupiters'' not previously seen in high-precision radial-velocity surveys. ",
        "conclusions": "From the combination of the spectroscopic and photometric solutions we have derived the key physical parameters of the planet. OGLE-TR-113b has an orbital period of only 1.43 days, a mass of $1.08 \\pm 0.28$~M$_{\\rm Jup}$, and a radius of $1.09 \\pm 0.10$~R$_{\\rm Jup}$ (Table~\\ref{tab:results}).  These values are consistent at the 1-$\\sigma$ level with the determinations by \\cite{Bouchy:04}.  Perhaps the most interesting parameter in this case is the very short orbital period.  OGLE-TR-56b, OGLE-TR-113b, and also the recently announced OGLE-TR-132b \\citep[][$P_{\\rm orb} = 1.69$~days]{Bouchy:04} all have orbital periods much shorter than the apparent cutoff of close-in giant planets at around 3-day periods, determined from the radial velocities surveys. Thus, OGLE-TR-56b (the first of these discoveries) can no longer be considered an oddity among the extrasolar planets, and it appears these new cases point toward an extremely interesting new class of ``very hot'' Jupiters.  It is worth pointing out that these three short-period planets are the result of just the first two campaigns conducted by the OGLE team, in relatively small fields toward the Galactic center and Carina. If these or similar surveys were to continue producing candidates at the current rate for a period of operation similar to that of the Doppler surveys, it is not unreasonable to expect that the number of very hot Jupiters could increase significantly and even exceed the number of 3-4-day period planets from the radial velocity searches. The \\emph{frequency} of occurrence of very hot Jupiters, however, appears to be much lower than that of the shortest-period Doppler planets, as discussed by \\cite{Bouchy:04}. Thus, the apparent inconsistency with the lack of any Doppler discoveries having periods as short as those of the OGLE planets may simply be due to a combination of their lower rate of occurrence and the much higher sensitivity to these objects in the photometric surveys. The latter is the result of the relatively short duration of the OGLE photometric campaigns (a few weeks) and the increased probability of transits from geometry, such that the chance of finding longer-period transiting planets actually falls off dramatically beyond $P_{\\rm orb}$ of 3-4 days. The extreme conditions of proximity to the parent stars in these very hot Jupiters opens up the possibility of very interesting theoretical studies into their structure and evolution, as well as migration scenarios."
    },
    "0404/astro-ph0404294_arXiv.txt": {
        "abstract": "{We present a high-resolution millimeter study of the very young Class~0 protostar IRAM~04191+1522 in the Taurus molecular cloud. N$_2$H$^+$(1-0) observations with the IRAM Plateau de Bure Interferometer and 30m telescope demonstrate that the molecular ion N$_2$H$^+$ disappears from the gas phase in the inner part of the protostellar envelope ($r < 1600$ AU, $n_{\\mbox{\\tiny {H$_2$}}} > 5 \\times 10^5$ cm$^{-3}$). This result departs from the predictions of current chemical models. It suggests either that N$_2$ is more depleted than the models predict, owing to a higher binding energy on polar ice or an enhanced grain chemistry transforming N$_2$ to less volatile species, or that strong deuterium fractionation enhances N$_2$D$^+$ to the detriment of N$_2$H$^+$. ",
        "introduction": "\\label{sec_intro}  Understanding the onset of gravitational collapse in dense cloud cores requires detailed (sub)millimeter studies of the structure of prestellar condensations and young protostars \\citep*[e.g.][]{Andre00}. However, recent observations have shown that molecules such as CO and CS deplete onto grain surfaces in the inner parts of dense cores \\citep[e.g.][]{Bacmann02,Tafalla02}. Since this depletion phenomenon is thought to affect many other species at high densities, studying the kinematics of pre/protostellar cores with future high-resolution instruments such as ALMA may be difficult and requires the identification of the best tracers of dense gas. N$_2$H$^+$ has been put forward as such a good tracer since it is much less sensitive to depletion effects than other species observed with single-dish telescopes \\citep[e.g.][]{Tafalla02}, even though \\citet{Bergin02} reported a slight decrease of its abundance toward the prestellar core B68.  Prime targets for observational studies of protostellar collapse are very young Class~0 objects such as IRAM~04191+1522 -- hereafter IRAM~04191 for short -- in Taurus \\citep[see][ -- hereafter \\citeauthor*{Andre99}]{Andre99}. This protostar features a prominent ($\\sim 0.5-1.5\\, M_\\odot$) envelope and a powerful bipolar outflow, but lacks a sizeable accretion disk. \\citet[][ -- hereafter \\citeauthor*{Belloche02}]{Belloche02} showed that the envelope is undergoing both extended infall motions and fast, differential rotation. They proposed that the rapidly rotating inner envelope ($r <$ 3500 AU) corresponds to a magnetically supercritical core decoupling from an environment still supported by magnetic fields and strongly affected by magnetic braking.  Here, we report new observations of IRAM~04191 carried out with the Plateau de Bure interferometer (PdBI) in the N$_2$H$^+$(1-0) line in an effort to probe the inner structure of the envelope. We discuss the results of these high-resolution observations which show that N$_2$H$^+$ disappears from the gas phase in the inner part of the envelope. ",
        "conclusions": ""
    },
    "0404/astro-ph0404407_arXiv.txt": {
        "abstract": "We discuss physical experiments achievable via the monitoring of stellar dynamics near the massive black hole at the Galactic center with a diffraction-limited, next generation, extremely large telescope (ELT). Given the likely observational capabilities of an ELT and what is currently known about the stellar environment at the Galactic Center, we synthesize plausible samples of stellar orbits around the black hole. We use the Markov Chain Monte Carlo method to evaluate the constraints that the monitoring of these orbits will place on the matter content within the dynamical sphere of influence of the black hole. We express our results as functions of the number $N$ of stars with detectable orbital motions and the astrometric precision $\\delta \\theta$ and spectroscopic precision $\\delta v$ at which the stellar proper motions and radial velocities are monitored. Our results are easily scaled to different telescope sizes and precisions. For $N=100$, $\\delta \\theta = 0.5 \\trm{ mas}$, and $\\delta v = 10 \\trm{ km s}^{-1}$---a conservative estimate of the capabilities of a 30 meter telescope---we find that if the extended matter distribution enclosed by the orbits at 0.01 pc has a mass greater than $\\sim 10^3 M_\\odot$, it will produce measurable deviations from Keplerian motion. Thus, if the concentration of dark matter at the Galactic Center matches theoretical predictions, its influence on the orbits will be detectable. We also estimate the constraints that will be placed on the mass of the black hole and on the distance to the Galactic Center, and find that both will be measured to better than $\\sim 0.1\\%$.  We discuss the significance of knowing the distance to within a few parsecs and the importance of this parameter for understanding the structure of the Galaxy. We demonstrate that the lowest-order relativistic effects, such as the prograde precession, will be detectable if $\\delta \\theta \\la 0.5 \\trm{ mas}$. Barring the favorable discovery of a star on a highly compact, eccentric orbit, the higher-order effects, including the frame dragging due to the spin of the black hole, will require $\\delta \\theta \\la 0.05 \\trm{ mas}$. Finally, we calculate the rate at which monitored stars experience detectable nearby encounters with background stars.  The encounters probe the mass function of stellar remnants that accumulate near the black hole.  We find that $\\sim 30$ such encounters will be detected over a ten year baseline for $\\delta \\theta = 0.5 \\trm{ mas}$. ",
        "introduction": "\\label{sec:intro}  Observational programs with ten meter class telescopes, including the W.~M.~Keck Observatory and the Very Large Telescope (VLT), have yielded a wealth of information on the stellar content inside the sphere of influence of the massive black hole at the Galactic center (GC; \\citealt{Ghez:98,Gezari:02,Hornstein:02,Figer:03,Genzel:03a,Ghez:03b,Schoedel:03}). The black hole is located at the center of a compact stellar cluster that has been the target of observational surveys for a decade (e.g.,~\\citealt{Krabbe:95,Figer:00,Gezari:02}).  Near-infrared monitoring with speckle and adaptive optics techniques has recently enabled complete orbital reconstruction of several stellar sources orbiting the black hole \\citep{Eckart:02,Schoedel:02,Schoedel:03,Ghez:03b}.  Sources have been monitored with astrometric errors of a few milli-arcseconds \\citep{Ghez:03a,Schoedel:03}, and radial velocity errors $< 50 \\trm{ km s}^{-1}$ \\citep{Eisenhauer:03, Ghez:03a}, allowing the detection of the accelerated proper motions of $\\sim 10$ stars. One of these stars has an orbital period of only $\\sim 15 \\trm{ yr}$ \\citep{Ghez:03b, Schoedel:03}.  The presence of a dark mass at the center of the Galaxy could in principle be inferred from the static nature of the radio source \\sgra located at the center of the stellar cluster \\citep{Backer:99,Reid:99}. Nevertheless, it is the stars with the shortest orbital periods that have provided unequivocal proof of the existence of a massive black hole and a measurement of its mass of $\\sim4\\times10^6 M_\\odot$ \\citep{Ghez:03b, Schoedel:03}.  Since, for a fixed angular scale, the orbital periods are proportional to $R_0^{3/2} M_{\\rm bh}^{-1/2}$ and the radial velocities are proportional to $R_0^{-1/2}M_{\\rm bh}^{1/2}$ where $R_0$ is the heliocentric distance to the black hole and $M_{\\rm bh}$ is its mass, the two parameters are not degenerate and can be determined independently \\citep{Eisenhauer:03}.  In spite of the quality of elementary data available about the black hole and the bright stellar sources, the matter content in the vicinity of the black hole remains unknown.  The observed stellar sources may represent only a fraction of the total matter content. Since the radial diffusion time $\\sim 10^{8-9}\\textrm{ yr}$ is shorter than the age of the bulge, a large number of massive compact remnants ($5-10M_\\odot$ black holes) could have segregated into, and may dominate the matter density inside the dynamical sphere of influence of the black hole \\citep{Morris:93,Miralda:00}.  Furthermore, adiabatic growth of the massive black hole could have compressed a pre-existing distribution of cold dark matter (CDM) \\citep{Ipser:87,Quinlan:95,Gondolo:99} and stars \\citep{Peebles:72,Young:80} into a dense ``spike'' .  A variety of dynamical processes, however, are capable of destroying such a spike \\citep{Ullio:01,Merritt:02,Gnedin:02,Merritt:03}. A sustained CDM spike would have implications for the detection of annihilation radiation for the CDM models in which the CDM consists of weakly-interacting massive particles (WIMPs).  The most complete catalogue of stars in the central parsecs was compiled by \\citet{Genzel:00} and \\citet{Schoedel:03}.  In a survey of the stellar sources, \\citet{Genzel:03a} infer a spatial number density of $n(r)\\propto r^{-1.4}$ over the radial range $0.004<r<0.4\\textrm{ pc}$.  Their sample was $50\\%$ complete for stars brighter than $K\\sim 18$, where completeness is defined as the percentage of stars in the field of view that are detectable and thus included in the sample. Expressed in terms of stellar mass, the sample is $50\\%$ complete for masses $m\\gtrsim 3M_\\odot$, assuming stars on the main sequence, a distance to the GC of 8.0 kpc \\citep{Reid:93} and $K$-band extinction of 3.3 mag \\citep{Rieke:89}. A picture is emerging in which the brightest stars in the Central Cluster ($<0.03$ pc) are young, main-sequence stars with apparent magnitudes $K>13$ and masses $10-15M_\\odot$.  The stars outside $0.03$ pc appear to be spectroscopically and kinematically distinct.  They span a larger range of magnitudes $K\\gtrsim 10$ and contain $\\sim40$ mass-losing Wolf-Rayet stars (e.g.,~\\citealt{Genzel:03a} and R.~Genzel, private communication).  Unlike the Central Cluster, these stars appear to belong to twin, misaligned stellar disks \\citep{Levin:03,Genzel:03a}.  The formation of the observed young stars with $\\times100$ larger specific binding energies relative to the black hole than that of the nearest observed accumulation of molecular gas (e.g.,~\\citealt{Jackson:93}) presents a challenge to star formation theories and is a persistent puzzle (e.g.,~\\citealt{Morris:93,Ghez:03a, Genzel:03a}).  A number of mechanisms for the formation and migration of stars in the tidal field of the massive black hole have been proposed \\citep{Gerhard:01,Gould:03,Hansen:03,Levin:03,Kim:03,Milosavljevic:04}. While the mechanisms have important implications, they are also each deficient in at least one way.  There is a dearth of giants in the GC region \\citep{Eckart:95}. Recently, \\citet{Figer:03} measured the radial velocities of 85 cool, normal giant stars with projected distances from the central region between $0.1-1$ pc. They find nearly complete deficiency of giants with large radial velocities ($V_{\\rm rad}>200\\textrm{ km s}^{-1}$). Since a star in a circular orbit at a distance of $0.1$ pc from the black hole has velocity $\\sim400\\textrm{ km s}^{-1}$, the absence of any such stars with comparable radial velocities indicates that the observed giants are indeed limited to the region outside the central $\\sim 0.5$ pc.  While the measured stellar density profile of the Galactic bulge is consistent with that of a singular isothermal sphere \\citep{Becklin:68}, the profile in the central parsec is not well known, especially for the lower-mass stellar populations.  Assuming relaxation that is driven by two-body processes, \\citet{Bahcall:76} showed that the equilibrium phase space distribution for a population of equal mass stars is a power law in density $\\rho\\propto r^{-7/4}$. For a multimass distribution the lighter stars are less centrally concentrated, resulting in a power-law profile that ranges from $r^{-3/2}$ for the least massive species to $r^{-7/4}$ for the most \\citep{Bahcall:77,Murphy:91}.  A coeval family of stars in the central region has reached equilibrium only if it is older than the relaxation time \\bea t_E & \\sim & \\frac{\\sigma^3}{G^2 m_{\\star} \\rho \\ln \\Lambda} \\nonumber \\\\ & \\approx & 2 \\times 10^8 \\textrm{ yr} \\left(\\frac{r}{1\\textrm{ pc}}\\right)^{1/4} \\left(\\frac{M_{\\rm bh}}{4 \\times 10^6 M_\\odot}\\right)^{3/2} \\nonumber \\\\ &  &          \\times \\left(\\frac{m_\\star}{10 M_\\odot}\\right)^{-1} \\left(\\frac{\\rho_{\\rm 1 pc}}{2 \\times 10^5 M_\\odot \\textrm{ pc}^{-3}}\\right)^{-1} \\left(\\frac{\\ln \\Lambda}{10}\\right)^{-1} \\eea where $\\sigma$ is the local linear stellar velocity dispersion, $m_\\star$ is the mass of a typical field star, $\\rho$ is the local stellar density, and $\\ln \\Lambda$ is the Coulomb logarithm.  Since the main sequence lifetime of stars more massive than $\\sim 2 M_\\odot$ is shorter than $t_E$, young massive stars in the GC are not relaxed; their distribution is primarily a reflection of their formative conditions. While lower mass dwarf stars are sufficiently old to be relaxed in the central potential, their distribution in the innermost region could be affected by an abundance of stellar mass black holes ($5 - 10 M_\\odot$). As products of normal stellar evolution, stellar mass black holes sink in the potential of the massive black hole \\citep{Morris:93,Miralda:00} and displace the less massive stars and remnants.  Speckle imaging and more recently adaptive optics with the Keck and VLT have provided several milliarcsecond astrometry, enabling the detection of proper motions within the inner 0.5 pc and accelerated proper motions of $\\sim 10$ stars within the inner 0.05 pc. Radial velocities with spectroscopic precisions of $\\delta v \\sim 30 \\trm{ km s}^{-1}$ have also been obtained for the star S0-2, which has been monitored for over 70\\% of its orbit including pericenter passage at $\\sim130\\textrm{ AU}$ from the black hole. These observations have enabled the black hole mass and GC distance to be measured to within $\\sim 10\\%$ \\citep{Ghez:03b, Schoedel:03}.  Here we examine the extent to which one can probe the GC potential by monitoring stars with a diffraction-limited, next generation, extremely large telescope (ELT). As compared with current 10 m class telescopes, the finer angular resolution of an ELT enables the orbital motions of many more stars to be detected, each at greater astrometric precision, $\\delta \\theta$, and spectroscopic precision, $\\delta v$. Given the range of possible sizes of future telescope and given the uncertainties in the ultimate capabilities of a specific telescope class (e.g., 30 meter telescopes) we choose to express our results not as functions of the ELT aperture but rather as functions of $\\delta \\theta$, $\\delta v$, and the number $N$ of stars with detectable orbital motions. We take $\\delta \\theta = 0.5 \\trm{ mas}$ and $\\delta v = 10 \\trm{ km s}^{-1}$ as our fiducial model, corresponding to a conservative estimate of the capabilities of a telescope with a $D = 30 \\trm{ m}$ aperture. We show that $N$ scales with telescope aperture as $N \\simeq 100 (D /30 \\trm{ m})^2$. We demonstrate that with an ELT one can measure the density profile of a dark matter spike and those general relativistic effects that scale as $(v/c)^2$, where $v$ is the speed of a star and $c$ is the speed of light. Furthermore, we show that the distance to the GC will be measured to remarkable precision. This will help place tight constraints on models of the overall Galactic structure. We also show that with an ELT one can detect the gravitational interactions between monitored stars and the background massive stellar remnants that accumulate near the central black hole. Such interactions may probe the mass function of the stellar mass black holes thought to dominate the matter density in the region.  The paper is organized as follows. In \\S~\\ref{sec:observations} we calculate the number of stars with accelerated proper motions that can be monitored with a given ELT based on its astrometric, spectroscopic, and confusion limits. We also describe a realistic monitoring program and demonstrate that confusion with the infrared emission from \\sgra is unlikely to affect an ELT's ability to measure stellar motions. In \\S~\\ref{sec:orbitmodel} we model the orbital data and estimate the magnitude of various non-Keplerian effects including Newtonian retrograde precession due to extended matter, relativistic prograde precession, precession induced by the coupling of orbits to the spin of the black hole, and the Roemer time delay.  In \\S~\\ref{sec:scattering} we consider the effect of stellar interactions on the motion of the monitored stars. Specifically, we estimate the rate at which discrete stellar encounters result in detectable changes of orbital motions. In \\S~\\ref{sec:method} we discuss a method for generating mock ELT orbital data and describe a computational technique for estimating uncertainties in the orbital parameters. The results of our calculations are given in \\S~\\ref{sec:results}. Finally, in \\S~\\ref{sec:discussion} we discuss astrophysical applications of the proposed observations. ",
        "conclusions": "\\label{sec:conclusions}  We have examined a variety of experiments that can be achieved through the infrared monitoring with an ELT of stars within a few thousand AU of the GC.  The astrometric limit of a 30 meter ELT is conservatively 0.5 mas and possibly as high as 0.1 mas. By comparison, the astrometric limit of current observations is $1 - 2 \\trm{ mas}$.  The greater point-source sensitivity and spectral resolution of an ELT enables the measurement of radial velocities with errors $\\la 10 \\trm{ km s}^{-1}$.  At present, of the $\\sim 10$ stars with measured accelerated proper motions, spectral lines have been detected only in S0-2, with radial velocity uncertainties of $\\sim 30 \\trm{ km s}^{-1}$. Measuring the radial velocities of stars breaks the degeneracy between mass and distance and thus yields a direct measurement of the distance to the GC. If the spectra of fainter stars can be obtained, the detection of deep molecular lines will improve upon the velocity estimates by an additional factor $\\times10$. The solar type stars that will be detectable with an ELT may therefore yield radial velocity uncertainties considerably smaller than $10 \\trm{ km s}^{-1}$.  A 30 meter ELT will be able to detect stars down to a $K$-band magnitude of $K \\sim 22$, approximately four magnitudes fainter than currently possible. Due to confusion, it will be difficult to detect still fainter stars. Using measurements of the $K$-band luminosity function within the inner $1\\arcsec$ of the GC, we estimate that such an ELT will detect the accelerated motion of $\\sim 100$ stars with semi-major axes in the range $200 \\la a \\la 3000 \\trm{ AU}$. Current observations are limited to the detection of $\\sim 10$ stars, all with $a \\ga 1000 \\trm{ AU}$. We find that the number of stars with detectable accelerated motion scales with the aperture of an ELT as $N \\simeq 100 (D / 30 \\trm{ m})^2$.  Given the observational capabilities of an ELT and the likely, albeit at low masses largely uncertain, stellar environment at the GC, we constructed a plausible sample of stellar orbits. The model includes the dynamical contribution of an extended distribution of dark matter around the black hole that is composed of stellar remnants and CDM. We find that for measurements at the precision obtainable with an ELT the uncertainty in the model parameters scale with the measurement errors $\\sigma$ (i.e., $\\delta \\theta$, $\\delta v$) and the number of monitored stars $N$ as roughly $\\sigma / N^{1/2}$. Thus, while we focus on the capabilities of a diffraction limited 30 meter ELT with $\\delta \\theta = 0.5 \\trm{ mas}$ and $\\delta v = 10 \\trm{ km s}^{-1}$, our results can be used to determine the capabilities of an ELT with different specifications. For example, a 100 meter ELT will detect $\\sim 10 \\times$ as many stars so that if it has astrometric and spectroscopic errors that are smaller by a factor of five, the measurement accuracy in the parameters will improve by a factor of approximately ten.  We find that with a 30 meter ELT the parameters $M_{\\rm bh}$ and $R_0$ will be measured to an accuracy better than $0.1\\%$. Determining $R_0$ to within a few parsecs will significantly constrain models of the Galactic structure as it aids the precise measurement of the dark matter halo shape.  While current observations of stellar proper motions are compatible with Keplerian motion, a number of dynamical effects produce significant deviations, including the Newtonian retrograde precession, the relativistic prograde precession, frame dragging due to the black hole spin, and interstellar interactions involving nearby encounters. All but the frame dragging effect produce non-Keplerian motions that are detectable with a 30 meter ELT.  Unfortunately, the spin of the massive black hole at the GC will probably be out of reach to kinematic studies unless an astrometric precision of $\\sim 0.05 \\trm{ mas}$ is achieved.  The presence of an extended distribution of matter results in a Newtonian retrograde precession due to differences in the amount of mass enclosed within an orbit's pericenter and apocenter. We considered extended matter density profiles consistent with current observations of the stellar distribution at the GC. We modeled the distribution as a power-law profile normalized such that $M_{\\rm ext}(r < 0.01 \\trm{ pc}) = 6000 M_\\odot$ and with slope $\\gamma = 1.5$ or 2. Standard models of dark matter clustering about a massive black hole predict similar profiles. An orbit monitoring program with a 30 meter ELT will constrain the mass and slope of such profiles to $\\sim 30\\%$ accuracy. Thus, monitoring orbits with an ELT provides a probe of the extended matter distribution within $\\sim 10^4$ Schwarzschild radii of the massive black hole at the GC.  We also calculated the rate at which the monitored stars experience detectable deflections due to stellar gravitational scattering encounters with background compact remnants. We considered a detection threshold set by the minimum detectable change in the velocity of a monitored star.  For a density cusp dominated by $\\sim 10 M_\\odot$ black holes, $\\sim 30$ nearby stellar encounters will be detected by a 30 meter ELT over a ten year observing baseline.  This will confirm the presence of a cusp of compact remnants at the GC and enable the measurement of the remnants' masses."
    },
    "0404/hep-ph0404132_arXiv.txt": {
        "abstract": "We study the possible signals of a pion string associated with the QCD chiral phase transition in LHC Pb - Pb collision at energy $ \\sqrt{s}=5.5$ TeV. In terms of the Kibble-Zurek mechanism we discuss the production and evolution of the pion string. The pion string is not topologically stable, it decays into neutral pions and sigma mesons which in turn decay into pions. Our results show that all the neutral pions from the pion string are distributed at the low momentum and the ratio of neutral to charged pions from the pion string violates the isospin symmetry. For the momentum spectra of the total pions, the signal from the sigma particle decay which is from the pion string will be affected by the large decay width of the sigma significantly. \\newline PACS number(s): 11.27.+d, 25.75.-q, 98.80Cq ",
        "introduction": "The formation of topological defects in phase transition is a very generic phenomenon in physics. It can be studied experimentally in different condensed matter systems. It is generally believed that the early evolution of Universe undergoes a sequence of phase transitions, and the produced topological defects in these phase transitions may have observable consequences to the properties of the universe today. For example, the cosmic strings have been suggested as one possible source for the primordial density perturbations that give rise to the large-scale structure of the universe and the temperature fluctuations of the cosmic microwave background (CMB) radiation\\cite{Rajantie03}\\cite{Villenkin00}. In particular, in Ref.\\cite{Brandenberger99} the effect of the pion string on the primordial magnetic field generation in the early universe has been considered and its cosmological significance is pointed out. However, in this paper, we will turn from cosmology to laboratory experiments and attempt to study the possible signals of the pion string in the heavy ion collision experiments which have many similarities with the nonequilibrium phenomena that also take place in heavy ion collisions experiments\\cite{Rajagopal93}\\cite{Rajagopal95}.  It is difficult to make experimental tests of our ideas about the formation and evolution of topological defects in cosmology directly. What we can do is to look for analogous processes in experimentally accessible condensed-matter systems. Fortunately, topological defects are formed at phase transitions in certain condensed matter systems such as super-fluids and superconductors, this phenomenon is theoretically very similar to its cosmological counterpart, and we can use this analogy to do \"cosmology experiments\"\\cite{Zurek85}. On a more fundamental level, these same experiments can be used to test our understanding of non-equilibrium dynamics of quantum field theories\\cite{Volovik03}.  In relativistic nucleus-nucleus collisions, some phenomena which like that happens in the Big Bang have been observed, this is called the Little Bang. Thus as hadron momentum spectra correlations provide strong evidence for the existence of the Little Bang: thermal hadron radiation with T=90-100MeV and strong 3-dimensional (Hubble like) expansion with transverse flow velocities 0.5-0.55c \\cite{Heinz99}. So the study of Little Bang from relativistic nucleus-nucleus collisions may construct a bridge between high energy particle physics and the cosmology. On the other hand, we are not yet in a position to give an evidence that quark-gluon plasma(QGP) has really been produced. It turns out to be a difficult task to figure out the theoretical picture of QGP even in equilibrium. Beyond the trivial level of the trees and the non-equilibrium properties of the QGP are essentially unknown. There is no unique signal of QGP in the understanding of nucleus-nucleus collisions so far. As pointed out by Rajantie \\cite {Rajan02}, since the heavy-ion collisions experiments are so complicated that the reliable and accurate theoretical calculations are needed in order to confront the experimental results, but our present understanding of the theory is too rudimentary for that.  Then the insight provided by condensed systems experiments is therefore likely to be extremely useful. In particular, it is believed that at a certain value of the beam energy, the QGP which is produced in the collision cools through a second-order transition point. Pion strings as well as other topological and non-topological strings are expected to be produced\\cite{Zhang98}\\cite{Balachandran02}. An early study on the effects of these strings in the case of heavy ion collisions and in the early universe has been performed in Ref.\\cite{Balachandran02}, in their paper, they speculate that formation and subsequent evolution of the network of these string defects can give rise to inhomogeneous distribution of baryons and also the energy density by using the Kibble mechanism. In this paper, we extend their works on the formation and evolution of strings and discuss the possible signal of pion strings during chiral phase transition in the high ion collisions.  So far the theoretical scenario what can be applied to study the formation of topological defects in systems with global symmetry is the Kibble-Zurek mechanism\\cite{Zurek85}\\cite{Kibble76}. Moreover, according to Pisarski and Wilczek \\cite{Pisarski84} the chiral phase transition is expected to be of the second order for two massless flavors, it is customary then in this paper to apply the Kibble-Zurek mechanism in order to study the formation and evolution of the pion string during Pb-Pb central collisions at the LHC with energy $\\sqrt{s}=5.5 TeV$.  The remainder of this paper is organized as follows. We give a brief review of the pion string in the linear sigma model in Sec.II. It is shown how to use the Kibble-Zurek mechanism to consider the formation of the pion string at LHC in Sec.III. We discuss the evolution and decay of the pion string and their possible observational consequences at LHC in Sec. IV. Sec. V is reserved for summary and discussion. ",
        "conclusions": "We have investigated the effects of the pion string in the experiment of the heavy ion collisions. Following the Kibble-Zurek mechanism pion strings are expected to be formed in LHC Pb - Pb collision at energy $\\sqrt{s}=5.5$ TeV, then decay after the freezing out time into pions. These pions are mostly distributed in two separated low momentum regimes. These effects are expected to be observable and differ from predictions of other models\\cite{Zhuan00}\\cite{Stephanov99}. The pion enhancement in a small window around the nonthermal momentum $p_0\\simeq 21.1$MeV for neutral pions. While for other nonthermal pions, the situation is completely different. If we ignore the large decay width of the sigma and take the sigma as a stable particle, there are the pion enhancement in a small window around the nonthermal momentum $p_0\\simeq 143$MeV, but actually the sigma has the large decay width, then the peak will smoothen out and almost tend to disappear. So it is difficult to detect this part of nonthermal pions in experiment. All the resultant pion spectrum depend strongly on how long the pion string can survive below the freeze out time $t_f$.  In this paper, we have made the assumption that pion strings can survive after the decoupling time, then both the neutral pions and sigma particles emitted from pion strings do not get chance to be thermalized. On the other hand, we do not exclude the situation in which part (even all) of pion strings will decay into pions and sigma particles when the time is very close to the decoupling time, and such produced pions from pion strings will be thermalized by the final state interactions and the peak in the pion spectra due to the pion string decay will disappear partly (or completely). However, even though this situation is happened, there still have possible signals of the pion string produced. Then pion string decay can lead to experimentally observable anomlies which are very similar to the DCC(Disoriented Chiral Condensate) decay. It is the ratio of neutral to charged pions, $r=\\frac{n_0}{n_0+n_{ch}}$, here $n_0$ is the number of neutral pions while $n_{ch}$ is corresponding to charged pions, which is different from what is naively expected ($\\frac{1}{3}$) if there are pion strings produced during chiral phase transition. Also enhancement of sigma which in turn decay into pions around the decoupling time will also result in nonthermal pions, though thses nonthermal pions is difficult to be detected because of the large decay width of the sigma.  Therefore, in order to obtain more reliable conclusions, we need to know more detail of the process of the pion string decay, the sigma decay and behaviors of the fire-ball at freeze out time in heavy ion collisions."
    },
    "0404/hep-ph0404256_arXiv.txt": {
        "abstract": "PeV neutrinos produce particle showers when they interact with the atomic nuclei in ice. We briefly describe characteristics of these showers and the radio Cherenkov signal produced by the showers.  We study pulses from electromagnetic (em), hadronic, and combined em-hadronic showers and propose extrapolations to EeV energies. ",
        "introduction": "Ultrahigh energy (UHE) neutrinos have energies of a PeV and higher. Although these neutrinos have not been observed, a number of models predict their existence. Another reason to believe in UHE neutrinos is the existence of UHE cosmic rays which have been observed by a number of detectors. Any models that account for these cosmic rays also predict existence of UHE neutrinos. Ultrahigh energy neutrino detection will have a great impact on fundamental physics and astrophysics. Because of their ultrahigh energies, they will help us explore fundamental interactions at energies well beyond the reach of man-made accelerators. Also, being neutral and presumably weekly interacting, these neutrinos, unlike charged particles, will bring us directional information about the source. Being weakly interacting they can help us explore regions of astrophysical objects which otherwise are opaque.  The electrons and hadrons from the neutrino-nucleon interactions will initiate particle showers as, unlike muons and taus, they dump their energy very quickly in matter through electromagnetic and strong interactions. These showers produce radio and optical radiation which can be detected\\cite{zhs,amz,ketal1,ketal2} in detectors like RICE\\cite{ketal1,ketal2} and ICECUBE\\cite{icecube}. Radio ice Cherenkov experiment (RICE)\\cite{ketal1,ketal2}, located at the south pole, can detect these showers from the radio signal they produce.  Here we give a brief introduction to the theory of radio emission from the showers\\cite{retal}. Particle showers in matter, unlike air showers, travel only a few 10's of meters and dump all their energy in the target material. The number of particles and the spacial size of the shower depends on the energy of the primary particle. For ultrahigh energy primaries, the shower, at a given time at shower maximum (which occurs at a few meters from the primary vertex), has millions of particles located in a few cubic centimeters. As these particles are moving faster than the speed of light in the given material, the charged particles will emit Cherenkov radiation at all wavelengths. Whether this radiation will be coherent depends on its wavelength, the net charge in the shower, and the shower size (the volume in which most of the shower particles are contained at a given time). For a typical shower size (a few cubic centimeters) one expects to get a coherent signal at Cherenkov angle due to constructive interference of the radio radiation from different parts of the shower. Hence one gets a strong radio signal from the shower which makes the radio detection of the showers (RICE) a very attractive technique as compared to the optical detection (ICECUBE\\cite{icecube}).  For radio detection of the showers, one needs a radio transparent material, with a huge volume (to have detector effective volume large enough, of the order of a cubic kilometer, to detect tiny fluxes of UHE neutrinos), in a region with as low as possible background radio signal. The last requirement makes a cold (to reduce background radio emission from atoms due to their vibrations at finite temperature) and isolated area an ideal environment for radio detection. Cold ice is remarkably transparent to radio transmission, with attenuation length of more than a kilometer.  RICE\\cite{ketal1,ketal2}, founded in 1995, is located 100m below surface, above AMANDA\\cite{hundertmark}, at the South pole. The instrumented volume is (200m)$^{3}$ which consists of 20 radio dipole antennas optimized to detect 0.5GHz radio signal. The effective volume depends on signal strength which in turn depends on the energy of the primary. For UHE showers, RICE effective volume is of the order of a km$^{3}$% . RICE expects to upgrade its volume by an order of magnitude in the coming years. ",
        "conclusions": "Ultrahigh energy neutrinos are a window to explore fundamental physics and astrophysics. They have not been observed/identified yet. However, a number of experiments have observed UHE cosmic rays and almost all the models that account for these cosmic rays also predict the existence of UHE neutrinos. These neutrinos can produce particle showers when they interact with matter. These showers produce coherent radio signal while their optical signal is incoherent. The strength of the radio signal depends on the net charge in the shower which in turn depends on the number of particles in the shower which increases with neutrino energy. This makes radio detection technique very powerful at UHE energies (as the coherent radio signal is much larger than the optical signal at these energies). RICE expects to detect the radio signal from these UHE initiated neutrino showers. So far RICE has produced strict upper bounds on some of the UHE neutrino flux models.  Simulation results show that hadronic showers, as compared to EM showers for energies below a PeV, are not as efficient in producing radio signal. However, as one goes to ultrahigh energies, hadronic and EM showers become equally efficient in producing radio signals . The signal rises linearly with energy at ultrahigh energies.  UHE neutrino detection situation is promising. AMANDA and RICE are taking data. RICE expects upgrade in the coming years which will increase its volume by an order of magnitude. ICECUBE and ANITA are funded and in development stages. A number of other projects are underway."
    },
    "0404/astro-ph0404139_arXiv.txt": {
        "abstract": "Analytic approximations for synchrotron, synchrotron self-Compton (SSC), and external Compton (EC) processes are used to constrain model parameters for knot and hot-spot emission in extended jets of radio galaxies.  Equipartition formulas are derived that relate the Doppler factor $\\delta$ and comoving magnetic field $B$ assuming a nonthermal synchrotron origin of the radio emission, and synchrotron, SSC and EC origins of the X-ray emission.  Expressions are also derived for $\\delta$ and $B$ that minimize the total jet powers of the emitting region in synchrotron, SSC and EC models for the X-ray emission. The results are applied to knot WK7.8 of PKS 0637-752. Predictions to test two-component synchrotron and EC models are made for {\\it Chandra} and {\\it GLAST}. ",
        "introduction": "The ability of the {\\it Chandra X-ray Observatory} to resolve knots and hot spots in radio jets has opened a new chapter in jet research \\citep{sta04}. The radio emission in the extended jets on multi-kpc -- Mpc size scales is almost certainly nonthermal synchrotron radiation, but the origin of the X-ray emission is controversial. In many knots and hot spots, the spectral energy distribution (SED) at X-ray energies is a smooth extension of the radio and optical fluxes, so that a synchrotron origin of the X-ray emission is implied. In other cases, the X-ray flux exceeds the level implied by smoothly extending the radio/optical SED. But even in these cases, a one-component synchrotron interpretation may be possible \\citep{da02}, and a two-component synchrotron model can be preferred on energetic and spectral grounds \\citep{ad04}.  Besides the synchrotron mechanism, two other nonthermal processes are often considered to account for the X-ray fluxes observed from the knots and hot spots of radio jets, namely the synchrotron self-Compton (SSC) and the external Compton (EC) processes \\citep{hk02}. The target photons for the EC model can be CMBR photons \\citep{tav00}, or nuclear jet radiation \\citep{bru01}.  The EC model involving CMBR target photons is the currently favored interpretation for quasar X-ray knots and hot spots where the X-ray spectrum is not a smooth extension of the radio/optical spectrum \\citep{cgc01,sam04}. In this model, the X-ray emission from knots such as WK7.8 of PKS 0637-752 is argued to be due to CMB photons that are Compton-upscattered by nonthermal electrons from kpc-scale emitting regions in bulk relativistic motion at distances up to several hundred kpc from the central engine.  In a recent paper \\citep{ad04}, we have addressed difficulties of the X-ray EC model to explain the opposite behaviors of the X-ray and radio spatial profiles. This model requires large energies, particularly in debeamed cases where the observer is outside the Doppler beaming cone.  Here we concentrate on radiation from the extended X-ray jets, and provide equations suitable for observers to interpret multiwavelength X-ray data of knot and hot-spot emission with synchrotron, EC and SSC models,  and evaluate jet powers. Application to knot WK7.8 of PKS 0637-752 is used to illustrate the results. ",
        "conclusions": "We have derived the minimum jet power $L_{j,min}$ for synchrotron, EC, and SSC models of the knots and hot spot X-ray emission. The two-component synchrotron model does not connect the radio and X-ray fluxes, so that $L_{j,min}$ depends only on the product $\\delta B$. Thus moderate values of $\\delta$ are possible. For the parameters of knot WK7.8 with $\\gamma_1 = 30$ and $k_{pe} = m_p/(\\gamma_1 m_e)$, $\\hat y =  560(1+k_{pe})^{0.26}/(\\gamma_1^{0.16} r_{kpc}^{0.79})$, implying a minimum jet power of $7\\times 10^{46}$ ergs s$^{-1}$.  This is also equal to $L_{j,min}$ for the synchrotron/EC model, which however only holds for specific values of $\\delta=27$ and $B=36\\,\\mu$G (see Fig.\\ 2). This is because most of the energy is contained in electrons with $\\gamma \\sim \\gamma_1$, which is restricted to low values in the EC model. Much larger values of $\\gamma_1$ are allowed in the two-component synchrotron model, so that even smaller jet powers are possible in this model.  The large value of $\\delta = 27$ for the minimum jet power in the EC model implies small and improbable observing angles $\\theta \\leq \\delta^{-1}\\lesssim 2^\\circ$ with a deprojected length of $\\approx 2$ Mpc of the jet in PKS 0637-752. For larger observing angles corresponding to $\\delta \\lesssim 10$, a jet power exceeding $10^{48}$ ergs s$^{-1}$ is implied (Fig.\\ 2).  The jet power could be reduced assuming a jet composed of e$^+$-e$^-$ plasma. The decay of $\\gamma$ rays in the two-component synchrotron model \\citep{ad03,ad04} will produce pairs, but at much higher energies than needed. If the pair plasma were produced in a compact inner jet, then the energy requirements would be hard to explain because of large adiabatic losses in the course of expansion of the blob from sub-parsec to kpc scales (see \\citet{cf93} for other arguments against a pair jet).  The SSC model for knot WK7.8 is ruled out. This model formally satisfies the radio and X-ray fluxes with $\\delta_{SSC} = 0.014$ and corresponding magnetic field $B = 0.07$ Gauss when $\\Sigma_{\\rm C} =10$, but requires comoving particle and field energies $\\gg 10^{61}$ ergs.  The synchrotron/EC model for knot WK7.8 allows a nonvariable X-ray spectrum that cannot be softer than the radio spectrum, and predicts a $\\gamma$-ray flux at the level of $\\approx 0.3\\times 10^{-8}$ ph$(> 100$ MeV) cm$^{-2}$ s$^{-1}$. This is detectable at strong significance with {\\it GLAST} in the scanning mode over one year of observation, but may be difficult to distinguish from the variable inner jet radiation. The two-component synchrotron model allows variability at X-ray energies, though at a low level because of the source size, with nonvarying $\\gamma$-ray flux below the {\\it GLAST} sensitivity (see Fig.\\ 2 in \\citet{ad04}). An interesting study for {\\it GLAST} is to separate a highly variable inner jet component from a stationary emission component to determine maximum fluxes of the extended jet.    \\vskip0.5in We thank Dan Schwartz and Andrew Wilson for questions and discussions about these issues, and the referee for helpful comments on a different, earlier approach to this problem.  The work of CD is supported by the Office of Naval Research.  Research of CD and visits of AA to the NRL High Energy Space Environment Branch are supported by {\\it GLAST} Science Investigation No.\\ DPR-S-1563-Y."
    },
    "0404/astro-ph0404413_arXiv.txt": {
        "abstract": "This paper presents the clustering properties of hard (2-8\\,keV) X-ray selected sources detected in a wide field ($\\approx \\rm 2\\,deg^{2}$) shallow [$f_X(\\rm 2-8\\,keV)\\approx 10^{-14}\\rm \\, erg \\, cm^{-2} \\, s^{-1}$] and contiguous  XMM-{\\it Newton} survey. We perform an angular correlation function analysis using a total of  171 sources to the above flux limit.  We detect a $\\sim 4\\sigma$ correlation signal out to 300\\,arcsec with $w(\\theta < 300^{''})\\simeq  0.13 \\pm 0.03$. Modeling the two point correlation function as a  power law of the form $w(\\theta)=(\\theta_{\\circ}/\\theta)^{\\gamma-1}$ we find: $\\theta_{\\circ}=48.9^{+15.8}_{-24.5}$ arcsec and $\\gamma=2.2\\pm {0.30}$. Fixing the correlation function slope to $\\gamma=1.8$ we obtain  $\\theta_{\\circ}=22.2^{+9.4}_{-8.6}$\\,arcsec. Using Limber's intergral equation and a variety of possible luminosity functions of the hard X-ray population, we find a relatively large correlation length, ranging from $r_{\\circ}\\sim 9$ to 19 $h^{-1}$ Mpc (for $\\gamma=1.8$ and the {\\em concordance} cosmological model), with this range reflecting also different evolutionary models for the source luminosities and clustering characteristics. The relatively large correlation length is comparable to that of extremely red objects and luminous radio sources. ",
        "introduction": "It is well known that the study of the distribution of matter on large scales, using different extragalactic objects provides important constraints on models of cosmic structure formation. Since Active Galactic Nuclei (AGN) can be detected up to very high redshifts they provide information on the underlying mass distribution as well as on the evolution of large scale structure (cf. Hartwick \\& Schade 1989; Basilakos 2001 and references therein).  The traditional indicator of clustering, the angular two-point correlation function, is a fundamental and simple statistical test for the study of any extragalactic mass tracer and is relatively straightforward to measure from observational data. The overall knowledge of the AGN clustering using X-ray data comes mostly from the soft ($\\le 3$keV) X-ray band (Boyle \\& Mo 1993; Vikhlinin \\& Forman 1995; Carrera et al. 1998; Akylas, Georgantopoulos, Plionis, 2000; Mullis 2002), which is however biased against absorbed AGNs. Hard X-ray surveys ($\\ge 2$keV) play a key role in our understanding of how the whole AGN population, including obscured (type II) AGNs, trace the underlying mass distribution. Furthermore, understanding the spatial distribution of type II AGNs is important since they are among the main contributors of the cosmic X-ray background (Mushotzky et al. 2000; Hasinger et al. 2001; Giacconi et al. 2002).  Recently, Yang et al. (2003) performing a counts-in cells analysis of a deep ($f_{2-8 keV} \\sim 3 \\times 10^{-15}$ erg s$^{-1}$ cm$^{-2}$) {\\em Chandra} survey in the Lockman Hole North-West region, found that the hard band sources are highly clustered with $\\sim$ 60$\\%$ of them being distributed in overdense regions. The XMM-{\\em Newton} with $\\sim 5$ times more effective area, especially at hard energies, and $\\sim 3$ times larger field of view (FOV) provides an ideal instrument for clustering studies of X-ray sources.  In this paper we estimate for the first time the angular correlation function of the XMM-{\\em Newton} hard X-ray sample. Using Limber's equation and different models of the luminosity funtion for these sources we derive the expected spatial correlation function which we compare with that of a variety of extragalactic populations. Hereafter, all $H_{\\circ}$-dependent quantities will be given in units of $h\\equiv H_{\\circ}/100$ km $s^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "In this paper we explore the  clustering properties of hard (2-8\\,keV) X-ray selected sources using a wide area ($\\rm \\approx 2\\,deg^2$) shallow [$f_X(\\rm 2-8\\,keV)\\approx 10^{-14}\\rm \\, erg \\, cm^{-2} \\, s^{-1}$] XMM-{\\it Newton} survey. Using an angular correlation function analysis we measure a clustering signal at the $\\sim 4\\sigma$ confidence level. Modeling the  angular correlation function by a power-law, $w(\\theta)=(\\theta_{\\circ}/\\theta)^{\\gamma-1}$, we estimate $\\theta_{\\circ}=48.9^{+15.8}_{-24.5}$\\,arcsec and  $\\gamma=2.2\\pm {0.30}$. Fixing the correlation function slope to $\\gamma=1.8$ we estimate $\\theta_{\\circ}=22.2^{+9.4}_{-8.6}$\\,arcsec. Using a variety of luminosity functions and evolutionary models the Limber's inversion provides correlation lengths which are in the range $r_{\\circ} \\sim 10 -19 \\; h^{-1}$ Mpc, typically larger than those of galaxies and optically selected QSO's but similar to those of strongly clustered populations, like EROs and luminous radio sources."
    },
    "0404/astro-ph0404249_arXiv.txt": {
        "abstract": "{We compiled and investigated the infrared/sub-mm/mm SED of the new outburst star IRAS~05436$-$0007 in quiescent phase.  The star is a flat-spectrum source, with an estimated total luminosity of $L_{\\rm bol}\\,{\\approx}\\,5.6\\,L_{\\sun}$, typical of low-mass T\\,Tauri stars. The derived circumstellar mass of $0.5$\\,$M_{\\sun}$ is rather high among low-mass YSOs.  The observed SED differs from the SEDs of typical T\\,Tauri stars and of 4 well-known EXors, and resembles more the SEDs of FU\\,Orionis objects indicating the presence of a circumstellar envelope. IRAS~05436$-$0007 seems to be a Class II source with an age of approximately 4${\\times}10^5$\\,yr. In this evolutionary stage an accretion disk is already fully developed, though a circumstellar envelope may also be present.  Observations of the present outburst will provide additional knowledge on the source. ",
        "introduction": "\\label{sect:Intro}  On 23 Jan 2004 the amateur astronomer J.W.~McNeil discovered a new nebula towards the Orion\\,B molecular cloud, close to the diffuse nebulosity Messier 78 (McNeil et al.~2004). The object was not visible in either of the two Palomar Surveys (1951, 1990), but a photograph taken in 1966 for the book ``The Messier Album''\\footnote{http://www.seds.org/messier/} (Mallas \\& Kreimer 1978) shows a bright nebulosity very similar to the one of today. Also in the very deep [SII] image of Eisl\\\"offel and Mundt (1997), taken in October 1995, parts of the nebula are clearly visible though fainter than in 1966. The alternation of active and quiescent periods, suggested by these earlier observations, indicates that the event, probably the eruption of a pre-main sequence star, may be similar to the well-known EXor-type outbursts.  At infrared and sub-millimetre wavelengths, however, the source was observable also during the quiescent periods (IRAS, 2MASS, Lis et al.~1999, Mitchell et al.~2001).  At these wavelengths the emission is due to thermal radiation of circumstellar dust. The infrared/sub-mm/mm data offer a possibility to study the circumstellar matter -- which is likely responsible for the explosion via a sudden rise of the accretion onto the star (Hartmann \\& Kenyon~1996) -- prior to an outburst.  In this paper we collect all infrared/sub-mm observations available in the literature and compile a spectral energy distribution (SED) representative of the quiescent phase. The SED will be analysed, and compared with SEDs of pre-main sequence stars, including several known FUORs and EXors. ",
        "conclusions": "We compiled and investigated the infrared/sub-mm/mm SED of the new outburst star IRAS~05436$-$0007 in quiescent phase.  The star is a flat-spectrum source, with an estimated total luminosity of $L_{\\rm bol}\\,{\\approx}\\,5.6\\,L_{\\sun}$, typical of low-mass T\\,Tauri stars. The derived circumstellar mass of $0.5\\,M_{\\sun}$ is rather high among low-mass YSOs.  The observed SED differs from the SEDs of typical T\\,Tauri stars and of 4 well-known EXors, and resembles more the SEDs of FU\\,Orionis objects indicating the presence of a circumstellar envelope. IRAS~05436$-$0007 seems to be a Class II source with an age of approximately 4${\\times}10^5$\\,yr. In this evolutionary stage an accretion disk is already fully developed, though a circumstellar envelope may also be present.  Observations of the present outburst will provide additional knowledge on the source."
    },
    "0404/astro-ph0404280_arXiv.txt": {
        "abstract": "\\noindent Deep optical observations of the spectra of 12 Galactic planetary nebulae (PNe) and 3 Magellanic Cloud PNe were presented in Paper~I by Tsamis et al. (2003b), who carried out an abundance analysis using the collisionally excited forbidden lines. Here, the relative intensities of faint optical recombination lines (ORLs) from ions of carbon, nitrogen and oxygen are analysed in order to derive the abundances of these ions relative to hydrogen. The relative intensities of four high-$l$ C~{\\sc ii} recombination lines with respect to the well-known 3d--4f $\\lambda$4267 line are found to be in excellent agreement with the predictions of recombination theory, removing uncertainties about whether the high C$^{2+}$ abundances derived from the $\\lambda$4267 line could be due to non-recombination enhancements of its intensity.  We define an abundance discrepancy factor (ADF) as the ratio of the abundance derived for a heavy element ion from its recombination lines to that derived for the same ion from its ultraviolet, optical or infrared collisionally excited lines (CELs). All of the PNe in our sample are found to have ADF's that exceed unity. Two of the PNe, NGC\\,2022 and LMC~N66, have O$^{2+}$ ADF's of 16 and 11, respectively, while the remaining 13 PNe have a mean O$^{2+}$ ADF of 2.6, with the smallest value being 1.8.   Garnett \\& Dinerstein (2001a) found that for a sample of about a dozen PNe the magnitude of the O$^{2+}$ ADF was inversely correlated with the nebular Balmer line surface brightness. We have investigated this for a larger sample of 20 PNe, finding weak correlations with decreasing surface brightness for the ADF's of O$^{2+}$ and C$^{2+}$. The C$^{2+}$ ADF's are well correlated with the absolute radii of the nebulae, though no correlation is present for the O$^{2+}$ ADF's. We also find both the C$^{2+}$ and O$^{2+}$ ADF's to be strongly correlated with the magnitude of the difference between the nebular [O~{\\sc iii}] and Balmer jump electron temperatures ($\\Delta T$), corroborating a result of Liu et al. (2001b) for the O$^{2+}$ ADF. $\\Delta T$ is found to be weakly correlated with decreasing nebular surface brightness and increasing absolute nebular radius.  There is no dependence of the magnitude of the ADF upon the excitation energy of the UV, optical or IR CEL transition used, indicating that classical nebular temperature fluctuations---i.e. in a chemically homogeneous medium---are not the cause of the observed abundance discrepancies. Instead, we conclude that the main cause of the discrepancy is enhanced ORL emission from cold ionized gas located in hydrogen-deficient clumps inside the main body of the nebulae, as first postulated by Liu et al. (2000) for the high-ADF PN NGC\\,6153. We have developed a new electron temperature diagnostic, based upon the relative intensities of the O~{\\sc ii} 4f--3d $\\lambda$4089 and 3p--3s $\\lambda$4649 recombination transitions. For six out of eight PNe for which both transitions are detected, we derive O$^{2+}$ ORL electron temperatures of $\\le$300\\,K, very much less than the O$^{2+}$ forbidden-line and H$^+$ Balmer jump temperatures derived for the same nebulae. These results provide direct observational evidence for the presence of cold plasma regions within the nebulae, consistent with gas cooled largely by infrared fine structure and recombination transitions; at such low temperatures recombination transition intensities will be significantly enhanced due to their inverse power-law temperature dependence, while UV and optical CELs will be significantly suppressed.  \\noindent {\\bf Key Words:} ISM: abundances -- planetary nebulae: general ",
        "introduction": "This is the second of two papers devoted to the study of elemental abundances in a sample of Galactic and Magellanic Cloud planetary nebulae (PNe). In a companion paper, Tsamis et al. (2003a) have presented a similar analysis of a number of Galactic and Magellanic Cloud {\\hii} regions. The main focus of these papers is on the problem of the optical recombination-line emission from heavy element ions (e.g. {\\cpp}, {\\npp}, {\\opp}) in photoionized nebulae.  The main manifestation of this problem is the observed discrepancy between nebular elemental abundances derived from weak, optical recombination lines (ORLs; such as {\\cii} $\\lambda$4267, {\\nii} $\\lambda$4041, {\\oii} $\\lambda\\lambda$4089, 4650) on the one hand and the much brighter collisionally-excited lines (CELs; often collectively referred to as forbidden lines) on the other (Kaler 1981; Peimbert, Storey \\& Torres-Peimbert 1993; Liu et al. 1995, 2000, 2001b; Garnett \\& Dinerstein 2001a; Tsamis 2002; Tsamis et al. 2003a), with ORLs typically being found to yield ionic abundances that are factors of two or more larger than those obtained from CELs emitted by the same ions. A closely linked problem involves the observed disparity between the nebular electron temperatures derived from the traditional {\\foiii} ($\\lambda$4959+$\\lambda$5007)/$\\lambda$4363 CEL ratio and the {\\hi} Balmer discontinuity diagnostic: the latter yields temperatures that are in most cases lower than those derived from the {\\foiii} ratio (Peimbert 1971; Liu \\& Danziger 1993b; Liu et al. 2001b; Tsamis 2002).  The ORL analysis of the current paper is based upon deep optical spectra of twelve galactic and three Magellanic Cloud PNe that were acquired by Tsamis et al. (2003b; hereafter Paper~I). Paper~I describes how the observations were obtained and reduced and presents tabulations of observed and dereddened relative intensities for the detected lines. Collisionally excited lines (CELs) in the spectra were used to derive nebular electron temperatures and densities from a variety of diagnostic ratios. They also derived CEL-based abundances for a range of heavy elements, using standard ionization correction factor (icf) techniques to correct for unobserved ion stages.  In the current paper we analyze the ORL data that were presented in Paper~I. In Section~2 we derive recombination-line ionic abundances for a number of carbon, nitrogen and oxygen ions. Section~3 presents a comparison between total C, N and O abundances derived from ORLs and from ultraviolet, optical and infrared CELs and derives abundance discrepancy factors (ADF's; the ratio of the abundances derived for the same ion from ORLs and from CELs) for a range of carbon, nitrogen and oxygen ions. In Section~4 we investigate how ORL/CEL ADF's correlate with other nebular parameters, such as the difference between [O~{\\sc iii}] forbidden line and H~{\\sc i} Balmer jump temperatures; the H$\\beta$ nebular surface brightness; and the nebular absolute radius. Section~5 looks at whether the observational evidence provides support for the presence of classical Peimbert-type temperature fluctuations within the nebulae, and whether the observational evidence points to strong density variations within the nebulae. In Section~6 we present evidence for the presence of cold plasma ({\\elt} $\\leq$ 2000~K) in a number of nebulae in our sample, making use of the fact that the strengths of several well observed O~{\\sc ii} and He~{\\sc i} recombination lines have sufficiently different temperature dependences for the relative intensities of two O~{\\sc ii} lines, or two He~{\\sc i} lines, to be used as diagnostics of the electron temperatures prevailing in their emitting regions. Section~7 summarizes our conclusions. ",
        "conclusions": "In Sections~3 and 5.2 it was demonstrated that there is no dependence of the magnitude of the nebular ORL/CEL abundance discrepancy factors upon the excitation energy of the UV, optical or IR CEL transition used (see Table~9), indicating that classical (i.e. in a chemically homogeneous medium) nebular temperature fluctuations are not the cause of the observed abundance discrepancies. This reinforces the same conclusion that was reached in Paper~I, based there upon the fact that [O~{\\sc iii}] electron temperatures derived from the ratio of the 52 and 88-$\\mu$m FS lines to the 4959 and 5007-\\AA\\ forbidden lines were greater than or comparable to those derived from the ratio of the higher excitation energy 4363-\\AA\\ transition to the 4959 and 5007-\\AA\\ lines---if temperature fluctuations in the ambient nebular material were the cause of the ORL/CEL abundance discrepancies then the IR line-based ratio should yield {\\em lower} temperatures. We conclude instead that the main cause of the abundance discrepancies is enhanced ORL emission from cold ionized gas located in hydrogen-deficient clumps inside the nebulae, as first postulated by Liu et al. (2000) for the high-ADF PN NGC~6153.  When nebular heavy element abundances exceed about five times solar, cooling by their collisionally excited infrared fine-structure (IR FS) lines is alone sufficient to balance the photoelectric heating from atomic species. Since cooling by the low-excitation IR FS lines saturates above a few thousand K, nebular electron temperatures in such high-metallicity regions will drop to values of this order, the exact value being determined by the H to heavy element ratios and the input ionizing spectrum. It is therefore physically plausible and self-consistent for the heavy element ORLs that yield enhanced abundances relative to those derived from CELs to also indicate very low electron temperatures for the regions from which they emit.  Liu et al.  (2000) in their empirical modelling of NGC\\,6153 had postulated the presence of H-deficient ionized clumps within the nebula that would be cool enough to suppress optical forbidden-line emission and even some infrared fine-structure line emission but which would emit strongly in heavy element recombination lines, due to the inverse power-law temperature dependence of ORL emission. P\\'{e}quignot et al. (2002) have constructed photoionization models of NGC\\,6153 and M~1-42 incorporating H-deficient inclusions and found equilibrium {\\elt}'s of $\\sim$~10$^3$\\,K in the H-deficient clumps and $\\sim$~10$^4$\\,K in the ambient gas, with the H-deficient ionized regions being within a factor of two of pressure equilibrium with the ambient nebular gas. In both cases, the H-deficient model components contained only $\\sim$~1\\% of the total ionized mass, so that the overall metallicity of the whole nebula was close to that of the `normal' high-temperature component.  Similar, or lower, mass fractions for the postulated H-deficient clumps in the typical PNe studied here, which have lower ADF's than the extreme cases discussed above, should ensure that their integrated IR FS line emission will not represent a significant perturbation to the integrated IR FS emission from the ambient high-temperature nebular material that forms the majority of the nebular mass, particularly if the H-deficient clump electron densities exceed the rather low critical densities of 500 -- 3000~cm$^{-3}$ that correspond to the [O~{\\sc iii}] and [N~{\\sc iii}] IR FS lines that have been investigated in this paper. The strong inverse power-law temperature dependence of ORL emission means that material at a temperature of 500\\,K will emit an ORL such as O~{\\sc ii} $\\lambda$4649 eighteen times more strongly than at 10$^4$~K. If 0.5\\% of the nebular mass was located in 500\\,K clumps having an electron density ten times that of the ambient nebular gas, then the integrated ORL emission from the clumps would exceed that from the ambient gas by a factor of nine.  Hydrogen is not expected to be entirely absent in the postulated H-deficient clumps. The steep inverse power-law temperature dependence of its recombination emission from the clumps could be responsible for the strong correlation, discussed in Section~4.1, between the C$^{2+}$ or O$^{2+}$ ADF's and the temperature difference, $\\Delta T$, between the [O~{\\sc iii}] optical forbidden line and H~{\\sc i} Balmer jump electron temperatures (see Fig.\\,~3). The correlation between ionic ADF's and decreasing nebular surface brightness (Fig.\\,~4), or increasing absolute nebular radius (Fig.\\,~5), may indicate that the density contrast between the H-deficient clumps and the ambient nebular material increases as the nebula evolves, i.e. that the clump density decreases less than the ambient nebular density does as the nebula expands. Such a situation might arise if the clump ionized gas originates from photoevaporation of dense neutral cores (cometary knots) of the kind found in the Helix Nebula (Meaburn et al. 1992; O'Dell et al. 2000) and the Eskimo Nebula (NGC\\,2392; O'Dell et al. 2002). Indeed, the location of NGC\\,2392 in the $\\Delta T$ vs. $S$(H$\\beta$) and $\\Delta T$ vs. Radius diagrams (Fig.\\,~6) is consistent with this picture and indicates that NGC\\,2392 is a candidate high-ADF nebula.  In confirmation of this, optical and ultraviolet large aperture measurements presented by Barker (1991) for six positions in NGC\\,2392 yielded $\\lambda$4267/$\\lambda$1908 C$^{2+}$ ADF's ranging from 7 to more than 24.  We thus show NGC\\,2392 in the ADF(C$^{2+}$) vs. $\\Delta T$ diagram (Fig.\\,~3) and the ADF(C$^{2+}$) versus Radius diagram (Fig.\\,~5) using a mean C$^{2+}$ ADF of 18 derived for the positions observed by Barker. It would be of interest to obtain deep optical spectra of this nebula in order to examine its heavy element ORL spectrum in detail.  Available spatial analyses of long-slit PN spectra show that ORL/CEL ADF's peak towards the center of nebulae in the cases of NGC\\,6153 (Liu et al. 2000) and NGC\\,6720 (Ring Nebula; Garnett \\& Dinerstein 2001b). An examination of the results presented by Barker (1991) for NGC\\,2392 shows that the same trend is present in those data too: the inferred ADF(C$^{2+}$) increases towards the centre of the PN (along the aligned positions 4, 2, and 1; cf. Fig.\\,~1 of that paper). In NGC\\,6720 especially, the location of peak O~{\\sc ii} ORL emission does not coincide with the positions of the {\\it HST}-resolved dusty cometary knots, which populate the main shell of the Ring, but is displaced inwards from that of the peak [O~{\\sc iii}] emission (Garnett \\& Dinerstein 2001b). We speculate that such an effect could be due to the advanced photo-processing of knots in the Ring Nebula that were overcome by the main ionization front in the past and whose relic material, rich in heavy elements, is now immersed in the He$^{2+}$ nebular zone, being subjected to the intense radiation field of the central star. The question of the possible relationship between the ORL/CEL abundance discrepancy problem and the cometary knot complexes observed in many PNe warrants further investigation.  Ruiz et al. (2003) and Peimbert et al. (2004) have argued against the presence of H-deficient knots in NGC\\,5307 and NGC\\,5315 on the grounds that their 34~km~s$^{-1}$ resolution echelle spectra did not reveal a difference between the radial velocities or line widths of the heavy element ORLs and those of the main nebular lines, of the type exhibited by the high-velocity H-deficient knots in the born-again PNe A30 and A58. We note however that the H-deficient knot model that has been invoked to explain the ORL/CEL ADF's of typical planetary nebulae makes no specific predictions as to whether the knots should exhibit a different kinematic pattern from the bulk of the nebula. Scenarios for the origin of such knots in `normal' PNe include i) the evaporation of primitive material (comets, planetesimals) left over from the formation of the progenitor star (e.g. Liu 2003), which would predict ORL C/O and N/O ratios typical of the unprocessed ISM material out of which the star formed; or ii) that they originated as incompletely mixed material brought to the surface by the 3rd dredge-up and ejected along with the rest of the AGB progenitor star's outer envelope during the PN formation phase. In the latter case, the knots' ORL C/O and N/O ratios should show the same nucleosynthetic signatures as the rest of the ejected envelope. Although the postulated H-deficient clumps may not exhibit different kinematics from the bulk of the nebular material, the low inferred electron temperatures of the clumps and the consequent very low thermal broadening means that heavy element ORL's in high-ADF nebulae should exhibit much narrower line widths than do the strong forbidden lines. For example, a C~{\\sc ii} $\\lambda$4267 line originating from 1000\\,K material should have a FWHM of 2~km~s$^{-1}$ if thermal broadening dominates, i.e.  if velocity broadening of the line is minimal, as when the emission comes from the edge of the nebula, where material is moving in the plane of the sky, or comes from well-separated approaching and receding velocity components, e.g. near the centre of the nebula. Thus observations of suitably chosen nebular sub-regions at a resolving power of 1.5$\\times10^5$ may be capable of confirming the presence of the postulated cold plasma clumps."
    },
    "0404/astro-ph0404555_arXiv.txt": {
        "abstract": "We study both analytically and numerically hydrodynamical effects of two colliding shells, the simplified models of the internal shock in various relativistic outflows such as gamma-ray bursts and blazars. We pay particular attention to three interesting cases: a pair of shells with the same rest mass density (``{\\it equal rest mass density}''), a pair of shells with the same rest mass (``{\\it equal mass}''), and a pair of shells with the same bulk kinetic energy (``{\\it equal energy}'') measured in the intersteller medium (ISM) frame. We find that the density profiles are significantly affected by the propagation of rarefaction waves. A split-feature appears at the contact discontinuity of two shells for the ``equal mass'' case, while no significant split appears for the ``equal energy'' and ``equal rest mass density'' cases. The shell spreading with a few ten percent of the speed of light is also shown as a notable aspect caused by rarefaction waves. The conversion efficiency of the bulk kinetic energy to internal one is numerically evaluated. The time evolutions of the efficiency show deviations from the widely-used inellastic two-point-mass-collision model. ",
        "introduction": "\\label{sec:intro}  The internal shock scenario proposed by Rees (1978) is one of the most promising models to explain the observational feature of relativistic outflows as in gamma-ray bursts, and blazars (e.g., Rees \\& Meszaros 1992; Spada et al. 2001). In this scenario, the bulk kinetic energy of the outflowing plasma is converted into thermal energy and non-thermal particle energy by the shock dissipation and particle acceleration, respectively, and explain the large power of these objects. Based on this scenario, a lot of authors have attempted to link the observed temporal profiles to multiple internal interactions (e.g., Kobayashi, Piran, \\& Sari 1997 (hereafter KPS97); Panaitescu, Spada \\& Meszaros 1999; Tanihata et al. 2002; Nakar \\& Piran 2002 (hereafter NP02)), looking for crucial hints on the central engine of these relativistic outflows.  Most of the previous works focus on the comparison with the observed light curves and model predictions employing a simple inelastic collision of two point masses (KPS97) and little attention has been paid to hydrodynamical processes in the shell collision. However, it is obvious that, in the case of relativistic shocks, the time scales in which shock and rarefaction waves cross the shells are comparable to the dynamical time scale $\\Delta^{'}/c$, where $\\Delta^{'}$ is the shell width measured in the comoving frame of the shell and $c$ is the speed of light. Since the time scales of observations of these relativistic outflows (e.g., Takahashi et al. 2000 for blazar jet; Fishman \\& Meegan 1995 for GRBs) are much longer than the dynamical time scales, the light curves should contain the footprints of these hydrodynamical wave propagations. Thus, it is very interesting to clarify the  difference between the simple two-point-mass-collision (hereafter two-mass-collision) model and the hydrodynamical treatment. The recent study by Kobayashi \\& Sari 2001 (hereafter KS01) reports that collided shells are reflected from each other by the thermal expansion. Since they perform a hydrodymamical simulation and show the reflection feature for a single case, the detail of propagations of rarefaction waves for various cases of collisions is not discussed. The aim of this paper is to clarify the hydrodynamical effects including the propagations of rarefaction wave. As the simplest case, we explore the hydrodynamics of two-shell-collisions in the  internal shock model. Since we are mainly interested in the hydrodynamical processes themselves, it is beyond the scope of this paper to make a detailed comparison of the observed phenomena with the model results.  We consider the time evolution of two colliding shells in relativistic hydrodynamics in \\S 2. In \\S 3, we discuss the application to GRBs and blazars. The summary and discussion are given in \\S 4. ",
        "conclusions": ""
    },
    "0404/astro-ph0404233_arXiv.txt": {
        "abstract": " ",
        "introduction": "The idea that the universe we see around us is a member of an ensemble of universes (the multiverse), the remainder of which are beyond our view, is an old one, But it is one under active investigation nowadays. Much of the contemporary motivation comes from rather grandiose theoretical conceptualizations. These include, but are not limited to, string theory ideology, eternal inflation, and even Everett's many-worlds interpretation of quantum mechanics~\\cite{ref:a}.  But another motivation comes from the great progress being made in the experimental investigation of the structure and early history of our own universe, as well as the successful establishment of the standard model of particle physics. While the accumulation of these data is no doubt the best stimulant, it nevertheless has led to a mixing of direct information from the experiments with highly speculative material, remote from experimental test. Indeed, there is plenty of skepticism extant as to whether the consideration of universes which are causally disconnected from our own is science at all.  The purpose of this note is to make the study of ensembles of universes as data-driven as possible. This will entail, first, a specific definition of a universe as the contents of a comoving spacetime box, large enough to contain essentially everything we can expect to be able to measure (based on the present consensus picture of cosmology and particle physics), but still small compared to the extent of the total spacetime domain envisaged by the majority of cosmologists and other theorists. This allows a credible way of constructing a local ensemble of universes very similar to our own.  One simply generalizes the construction of our universe to other spacetime domains causally disconnected from our own, yet near enough to allow reasonable extrapolation of properties of our universe to the remainder of the local ensemble.  We envisage this construction of our universe in analogy to the establishment of a homestead in Kansas by a pioneer family. They would see their spatial environment, as we do see our universe, as flatter than a pancake~\\cite{ref:c}, all the way to their property line. And it would be not unreasonable---and in fact even correct---for them to envisage other similar homesteads beyond their horizon. Were the family to contain a theorist, he or she would in fact probably erroneously conclude that Kansas, and hence the number of such homesteads, in fact was infinite in extent~\\cite{ref:d}. However, far beyond Kansas will be found other pieces of fertile terrain, {\\em e.g.} in the Central Valley of California, which also are flatter than a pancake and which bear other similarities to, as well as differences from, the Kansas homesteads.  The local, ``Kansas\" ensemble of universes we shall consider will be assumed to possess a spatially flat FRW metric,  with similar initial cosmological boundary conditions. This ensemble is already phenomenologically of some use, because it can be considered as the ensemble implicitly used by inflationary cosmologists in making quantum averages of the inflationary perturbation spectra. But more interesting, and the focus of this note, are the analogues of the Central Valley homesteads. We take these to be similar to our universe in the sense of again being described by a spatially flat FRW metric, with perhaps some also being fertile, {\\em i.e.} capable of supporting life as we know it. The features that we assume can vary from universe to universe will include the parameters of the standard model and of cosmology.  One can generalize much further than this, at the cost of increasing the level of speculation~\\cite{ref:e}. We will stop at this point, however. The reasons are, first, that obviously this level is already quite speculative, and second, that the data-driven (as well as theoretical!) motivations for considering ensembles of universes can already be expressed with specificity at this level of speculation. The main data-driven consideration is that the values of many standard model parameters seem to be fine-tuned in a way to allow our existence in the universe. The existence of an ensemble of universes with a variety of standard-model parameters allows an ``anthropic\" interpretation of this property of standard-model data~\\cite{ref:a,ref:f}. And the main theoretical consideration comes from string theory, where a variety of string vacuua, with different standard-model properties, may easily be envisaged to occur in various causally disconnected regions of spacetime~\\cite{ref:g}.  These issues can all begin to be addressed in the context of the subensemble of universes that we construct, namely those that differ in a rather minimal way from our own. It is already of great interest to try to delineate the existence and properties of these ``nearest-neighbor\" universes. And the possibility of erroneous speculation, while still huge, is clearly going to be less than dealing (in the absence of data) with a more broadly generalized ensemble of universes possessing properties vastly different from our own.  Another analogy to the ensemble of universes is the perhaps familiar one of the planets in our own universe. These are already known to come in many varieties, but an especially interesting subset of planets consists of those containing life as we know it. This ``nearest-neighbor\" ensemble might consist of only Planet Earth, or it may contain a large number~\\cite{ref:h}. The question of which alternative is correct is not only irresistible to ask, but is also undeniably a scientific one. In order to address it in a systematic way, one needs to parametrize properties of planets in such a way that one can look at the distribution function of those parameters and find the size and location of the habitable island in parameter space. The output of such a program could be an estimate or bound on the number of planets which can in principle support life as we know it. This is in the case of planets a very daunting task, one that will need to be data-driven to make progress~\\cite{ref:i}. If universes as we have defined them have a similarly complex parametrization, it will be extremely difficult to make progress, simply due to lack of data. It is a serious criticism by those who remain skeptical of the utility of multiverse ideology that this problem presents an insuperable barrier to scientific consideration of this line of thinking.  The use of ``anthropic\" reasoning when dealing with ensembles of universes only exacerbates this problem. If the distribution function of standard model parameters, taken over the multiverse ensemble, is broad in each of the parameters, and if there is little correlation between the probability of finding the value of one parameter and that of another, then the program of finding naturalistic explanations for the values of standard model parameters essentially grinds to a halt. The replacement is a statement, unsupported by data, that they are determined by historical accident. To this author, such an option seems not much more scientific in nature than interpreting the values of the standard model parameters via ``intelligent design\".  However, we cannot ourselves impose a decision on how things work at this level. If the above option is true, then it may be that the scientific method will never have the power to find out the answers to the Big Questions. But, it is not clear that the problem of characterizing the ensemble of universes is in fact as difficult as characterizing the ensemble of habitable planets. As long as there is an outside chance that the problem is simpler and more tractable, there is no reason not to pursue that possibility. This optimistic position will in fact be our working assumption. We shall assume that many of the standard model parameters are strongly correlated with each other, thereby simplifying the nature of the aforementioned distribution function. While anthropic reasoning will be utilized, the amount will be limited by the above simplistic assumption of the existence of correlations. And the test of whether the assumption makes sense will be whether it delivers insight into any of the unsolved problems facing particle physics and cosmology. We shall argue on the basis of the results that the answer to this is positive. In particular we shall obtain a partial understanding of the nature of the hierarchy problems of particle physics: why the ratios of many of the standard model parameters are such large numbers. We do not claim that the hierarchy problems are thereby solved, but rather that they are recast in a different form, one which in turn appears to be able to be attacked by scientific methods.  In the next section, we discuss in detail our definition of a universe, and in Section 3 provide a concise description of its major properties. In Section 4 we review the nature of our assumed correlations of standard model parameters , emphasizing their simplicity and cogency, as well as the consequence that the habitable range of sizes is limited to within a factor two of the size of our universe. In Section 5, we discuss the anthropic implications of this scenario, in particular that if the distribution in sizes is peaked at small sizes, there results some understanding of the hierarchy problems of particle physics.  The above material contains the main messages of this paper. The remaining sections are spinoff topics even more speculative. Section 6 explores a possible connection between Bekenstein-Hawking horizon entropy and the vacuum structure of QCD. Section 7 explores whether the inflationary phase of our universe might have its own set of standard model parameters, correlated with the very large value of the Hubble parameter (dark energy) during that epoch. Section 8 raises the question of whether the starting time we have used in our construction of the universe may contain physics, and represent the actual onset of the inflationary epoch. Section 9 is devoted to concluding comments. ",
        "conclusions": "One of the main purposes of this paper has been to cast multiverse ideology in as bottoms-up, phenomenologically driven a way as possible. This is expedited by the definition, by construction, of a universe which is big enough to contain essentially all of presently conceivable phenomenology, but small enough to exclude most of the commonplace and extravagant theoretical speculations enveloping the subject. The experimental accuracy of the cosmological principle, within the portion of the universe that we have observed, then allows a very credible extrapolation to relatively nearby regions of spacetime, where an ensemble of such universes can be similarly constructed.  Going only this far seems to us a quite conservative procedure, at least by contemporary standards of theoretical cosmology. More interesting is the next step, which posits that the ensemble can be generalized to include more distant members, of similar FRW spacetime structure. But these members are assumed to have different standard model parameters, in particular different cosmological constants. In the context of present activity in string theory, this step is also not very radical. However, acceptance of this step then invites introduction of anthropic reasoning, especially with respect to the role of the cosmological constant as an anthropically determined parameter.  At this point the level of controversy increases, because as soon as anthropic arguments are introduced, it is hard to determine when to stop. If the cosmological constant is determined only anthropically, what about all the other standard model parameters? Do they not also serve as labels for the gigantic number of vacuua in the string theory landscape? And if they are all anthropically determined, where does that leave, say, the future program of particle physics? The original dreams of a final theory, as visualized two decades ago~\\cite{ref:v}, lay at the opposite extreme, with specific and naturalistic explanations of those parameters expected to be provided by the future theory. Instead they may be only constrained by the fact that we exist in the universe to observe them~\\cite{ref:g}.  It is not for us to dictate the answer. Maybe all the parameters are anthropic, in which case the scientific method becomes relatively impotent. Maybe none of them are, despite the evidence of fine-tuning of parameters which has been uncovered and studied by the anthropic community over the last fifty years~\\cite{ref:n}. And maybe the answer lies in between the extremes.  It is this latter hypothesis which we have chosen in this paper. It is implemented by the presumption that, within the ensemble of universes we have constructed, at least the principal parameters of the standard model are strongly correlated with the value of the cosmological constant. We have also proposed a specific form of the correlation. This last step has been greeted  with much skepticism, most often in a completely dismissive way.  Perhaps this occurs because of the lack of any overarching theoretical ideology to motivate the scaling hypothesis. In defense, we argue from experience in searching for alternatives that the scaling hypothesis stands out in its simplicity and robustness~\\cite{ref:yy}. Dreamers of final theories would naturally expect most standard model parameters to be correlated with each other, because a good theory should have very few independent parameters. But it is {\\em a priori} unlikely that the correlations take a simple form, so simple that the right answer can be correctly guessed. We work from a naive sense of optimism at this point, but are rather convinced that if there {\\it is} a simple form of the correlation, the proposed correlation is likely to be, at the least, very similar to  the correct one.  Support for this point of view comes from the output, which provides some understanding of the hierarchy problems of particle physics. As we discussed, the hierarchy problem is traded in for the problem of understanding certain ``critical exponents\" such as 1/3 for QCD and 1/4 for the electroweak sector. An attempt at understanding the 1/3 of QCD was sketched in Section 6.  Here we only point out that that attempt, along with other spinoff ideas in the subsequent two sections, are not at all abstract. They deal with the understanding of very real issues existing within our own universe, such as the structure of the QCD vacuum, and the nature of the big-bang ignition (``reheating\") epoch  of our universe. It is evidence that consideration of multiple universes may conceivably help, at a quite phenomenological level, in uncovering the nature of physical processes in our own universe."
    },
    "0404/astro-ph0404005_arXiv.txt": {
        "abstract": "We discuss interstellar temperature determinations using the excitation equilibrium of the $^2P$ levels of \\SiII\\ and \\CII.  We show how observations of the \\upperfs\\ fine structure levels of \\SiII\\ and \\CII\\ (which have significantly different excitation energies, corresponding to $\\sim413$ and 92 K, respectively) can be used to limit gas kinetic temperatures.  We apply this method to the $z=4.224$ damped Lyman-$\\alpha$ system toward the quasar \\pss.  The lack of significant absorption out of the \\SiII\\ \\upperfs\\ level and the presence of very strong \\CII\\ \\upperfs\\ provides an upper limit to the temperature of the \\cstar -bearing gas in this system. Assuming a solar Si/C ratio, the observations imply a $2\\sigma$ limit $T<\\maxtemp$ K for this absorber; a super-solar Si/C ratio gives stricter limits, $T<\\alphamaxtemp$ K.  The observations suggest the presence of a cold neutral medium; such cold gas may serve as the fuel for star formation in this young galaxy. ",
        "introduction": "High-redshift damped \\lya\\ systems (\\dla s) are the highest column density class of QSO absorption lines.  Defined by $\\log N(\\mbox{\\HI}) \\ge 20.3$ (Wolfe et al. 1986), these systems are thought to trace the interstellar medium (ISM) of high-redshift galaxies.  Dedicated surveys over the past two decades have helped trace the global properties of high-redshift DLAs, including their contribution to the cosmological baryon density (Storrie-Lombardi \\& Wolfe 2000; Prochaska \\& Herbert-Fort 2004), their chemical enrichment (e.g., Prochaska et al. 2003), their dust content (e.g., Pettini et al.\\ 1994), and molecular fraction (Ledoux, Petitjean, \\& Srianand 2003).  These studies have demonstrated that the DLAs have a baryonic mass density comparable to the mass density of modern galaxy disks, that the metallicity of high-$z$ DLAs is slowly increasing, and that the majority of DLA sight lines have low dust-to-gas ratios and molecular fractions.  Understanding the detailed physics of the ISM in DLAs is an important step in understanding high-redshift galaxies in general.  Wolfe, Prochaska, \\& Gawiser (2003a) have constructed detailed models for the thermal equilibrium of the ISM in a set of DLAs, calculating the heating rate experienced by the gas due to the ultraviolet emission from young hot stars.  They used observations of absorption out of the \\upperfs\\ level of \\CII\\ (hereafter \\cstar) -- a direct indicator of the cooling rate through [\\CII] 158 $\\mu$m emission (Pottasch, Wesselius, van Duinen 1979) -- to infer the actual heating rate experienced by the gas (assuming thermal equilibrium).  Their comparison of the inferred and calculated heating rates suggests DLAs harbor significant star formation.  The models of Wolfe et al. (2003a,b) required that the observed \\cstar\\ in DLAs arise in a cold neutral medium (CNM), i.e., in gas with temperatures $T\\la1000$ K; their WNM models give SFRs that violate observations of the bolometric background.  The detection of \\htwo\\ absorption in some high-\\z\\ DLAs is further evidence that at least some of these systems contain a CNM (Ledoux et al. 2003; Hirashita \\& Ferrara 2005).  However, such temperatures are at odds with 21-cm absorption studies; all $z\\ga3$ DLAs searched for 21-cm absorption show $T_S\\ga1400$\\,K ($2\\sigma$) (Kanekar \\& Chengalur 2003).  Wolfe et al.  (2003b) argue that this discrepancy is likely due to the differing properties of the sight lines probed by the optical background sources and by the more extended radio sources.   In this Letter we present a method for determining the kinetic temperature of interstellar matter based solely on basic atomic physics.  Our method compares the excitation of the upper \\upperfs\\ fine-structure levels in \\SiII\\ and \\CII, which have excitation energies that differ by a factor of four.  We describe our approach in \\S \\ref{sec:method}.  We apply this technique to the $z_{abs}\\approx4.224$ DLA toward the quasar \\pss\\ in \\S \\ref{sec:psscnm}, demonstrating that this DLA contains a substantial reservoir of cold gas.  Lastly, we discuss the implications of this temperature determination in \\S \\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  We have presented a method for limiting gas temperatures in the ISM of galaxies through measurements and analysis of the \\upperfs\\ fine structure levels of \\SiII\\ and \\CII.  We have applied this method to limit the properties of the \\cstar -bearing gas in the $z=4.224$ DLA toward \\pss.  We rule out the hypothesis that the \\cstar\\ absorption arises in a WNM.  Our conservative temperature limit for this gas is $T \\la \\maxtemp$\\,K; we obtain stricter temperature limits if [Si/C]$\\, > 0$ or [C/Fe]$\\, \\approx 0$: $T\\la \\alphamaxtemp$ K. The detection of a CNM (Wolfire et al. 1995) in a high-redshift \\dla\\ is significant: while the gas seen in this DLA is likely not associated with the dense star-forming clouds, our result demonstrates that the physical conditions of the ISM in this system do not preclude the existence of cold material, including, in principle, the dense clouds from which stars could form.   The detection of a CNM in the z=4.224 absorber toward PSS 1443+27 is the first detection at such a large redshift.  We stress that our measurements allow a WNM as part of a multiphase medium toward \\pss, but the majority of the \\cstar\\ cannot come from warm material.\\footnote{We have tested the effects of multiphase absorbers on our technique, with as much as half of the ground-state ions arising in a WNM.  We find that there is very little impact on the derived maximum temperatures (Howk et al.  2005).}  The existence of CNM material may be a feature of many high-\\z\\ DLAs, as suggested by Wolfe et al. (2003a,b).  While $z \\la 2$ \\HI\\ 21-cm absorption-line measurements suggest the presence of cold \\HI\\ in DLAs, no 21-cm absorption has been found in DLAs with $z\\ga3$ (Kanekar \\& Chengalur).\\footnote{Only six $z>2.9$ DLAs have been observed; \\pss\\ has not been searched for 21-cm absorption.}  We note, however, that at least three $z>2.5$ DLAs show \\htwo\\ absorption (Ledoux et al. 2003), indicating the presence of cold gas (Hirashita \\& Ferrara 2005).  Searching for \\htwo\\ toward \\pss\\ might be difficult due to the strength of the \\lya\\ forest at $z\\sim4$.  In the future we will apply our method for determining gas kinetic temperatures to DLAs for which 21-cm and \\htwo\\ measurements have been attempted.  We note that the DLA toward \\pss\\ may not be typical.  It shows a high metallicity ([Si/H]$\\ga -1$), especially compared with other $z>4$ DLAs (Prochaska et al. 2003).  Furthermore, we chose this DLA for this experiment because it has very strong \\cstar.  It has the highest \\cstar /\\HI\\ ratio measured, some $+0.6$ dex above the next highest (the $z=1.92$ DLA toward Q2206-19; Wolfe et al.  2004).  The intensity of radiation in this DLA calculated following Wolfe et al. (2004) is higher than all other systems.  Although the fuel for star formation may be present in this DLA, no optical counterpart to this DLA has been identified in deep ground-based and {\\em Hubble Space Telescope} images of \\pss\\ (Prochaska et al.\\ 2002; L.  Storrie-Lombardi, private communication).  The method presented here will be discussed further, and temperature limits given for a larger number of DLAs, in Howk et al. (2005).  We note that the use of the relative populations of the \\upperfs\\ levels of \\SiII\\ and \\CII\\ may also be useful for constraining temperatures in Milky Way gas."
    },
    "0404/astro-ph0404469_arXiv.txt": {
        "abstract": "\\bigskip  There is evidence that the transiting planet HD~209458b has a large exosphere of neutral hydrogen, based on a 15\\% decrement in Lyman-$\\alpha$ flux that was observed by Vidal-Madjar et al.\\ during transits.  Here we report upper limits on H$\\alpha$ absorption by the exosphere.  The results are based on optical spectra of the parent star obtained with the Subaru High Dispersion Spectrograph. Comparison of the spectra taken inside and outside of transit reveals no exospheric H$\\alpha$ signal greater than 0.1\\% within a 5.1~\\AA\\ band (chosen to have the same $\\Delta\\lambda/\\lambda$ as the 15\\% Ly$\\alpha$ absorption).  The corresponding limit on the column density of $n=2$ neutral hydrogen is $N_2 \\lsim 10^9$~cm$^{-2}$. This limit constrains proposed models involving a hot ($\\sim$10$^4$~K) and hydrodynamically escaping exosphere. ",
        "introduction": "A milestone in extrasolar planet research was reached when Charbonneau et al.\\ (2000) and Henry et al.\\ (2000) observed the photometric signal of transits by the low-mass companion of HD~209458. The companion was originally discovered by radial velocity measurements (Mazeh et al.\\ 2000), which specified only the orbital period (3.5 days), orbital eccentricity ($e<0.03$), and minimum mass of the companion ($M\\sin i = 0.69$~$M_{\\rm Jup}$). The transit light curves allowed the measurement of the companion's mass (by breaking the $\\sin i$ degeneracy) and radius (from the depth of the transit), providing an unambiguous case of a planetary-mass companion to a main sequence star, and a demonstration that the companion's density was that of a gas giant planet.  Aside from this historic importance, the discovery made possible a number of unique and important follow-up studies, based on more subtle changes in the received starlight that should occur during transits. In this paper we are concerned with changes due to the passage of a small fraction of the starlight through the planetary atmosphere. Selective absorption by atmospheric constituents causes the transit depth to depend upon wavelength, an effect that forms the basis of ``transmission spectroscopy.'' Upper limits on various atmospheric absorption features have been given by Bundy and Marcy (2000), Moutou et al.\\ (2001), Brown, Libbrecht, and Charbonneau (2002), and Moutou et al.\\ (2003). The first successful detection using this technique was by Charbonneau et al.\\ (2002), who observed a $(0.023\\pm 0.006)$\\% increase in transit depth in the yellow light of the sodium resonance doublet. A strong sodium signal had been predicted by Seager and Sasselov (2000), and subsequently by Brown (2001) and Hubbard et al.\\ (2001).  More recently, Vidal-Madjar et al.\\ (2003) reported a remarkably large transit depth of $(15\\pm 4)$\\% in the ultraviolet light of the Lyman-$\\alpha$ (Ly$\\alpha$) transition of neutral hydrogen. To produce such a large absorption, a spherical cloud of neutral hydrogen would need a radius of $4.3R_{\\rm Jup}$ or greater (depending on the optical depth), as compared to the planetary radius of $1.3R_{\\rm Jup}$ inferred from broad-band observations, and the Roche lobe radius of $3.6R_{\\rm Jup}$. Such a signal could be produced by an extended envelope of hydrogen atoms that are evaporating from the planet's upper atmosphere.  Vidal-Madjar et al.\\ (2003) noted that the apparent blueshift of the absorbing atoms (up to $-130$~km~s$^{-1}$) provided additional evidence for atmospheric escape, and speculated that this evaporation process accounts for the rarity of extrasolar planets with orbital periods smaller than 3 days. Additional data obtained by Vidal-Madjar et al.\\ (2004) were consistent with the previous Ly$\\alpha$ result, and provided evidence (at the 2--3$\\sigma$ level) for exospheric absorption by oxygen (O~{\\sc i}) and carbon (C~{\\sc ii}).  Further confirmation of the existence of this ``exosphere,'' and characterization of its temperature, density, and composition are clearly important goals.  Unfortunately, theoretical predictions for the properties of the exosphere are not robust, making it difficult to design strategies for observing the exosphere.  One approach is to search for additional effects of hydrogen, which has already been implicated by the Ly$\\alpha$ measurement. If a significant fraction of the hydrogen exists in the first excited state ($n=2$), it will produce extra absorption in the Balmer lines of the stellar spectrum, and in particular H$\\alpha$ (6563~\\AA).  Thus we were motivated to search for H$\\alpha$ absorption from the exosphere of HD~209458b.  A positive detection would confirm the existence of the hydrogen exosphere and constrain its density and temperature, with implications for the long-term evolution of ``hot Jupiter'' extrasolar planets.  It would also represent a new and more practical means of searching for exospheres of other extrasolar planets, given the necessity of observing Ly$\\alpha$ absorption above the Earth's atmosphere, and the confusing effects of Ly$\\alpha$ absorption by interstellar hydrogen and emission from the geocorona. Previously, Bundy and Marcy (2000) placed upper limits of 3--4\\% on H$\\beta$ and H$\\gamma$ absorption in a 0.3~\\AA\\ band, although they did not examine H$\\alpha$ due to contamination of that region of the spectrum by I$_2$ lines. Moutou et al.\\ (2001) also searched for exospheric absorption in optical spectra. They did not remark on H$\\alpha$ in particular, but set upper limits of $\\approx$1\\% for any features of width 0.2~\\AA, and noted that the limits could be improved significantly with higher quality spectra taken in better atmospheric conditions.  We have been obtaining optical echelle spectra of HD~209458 with the Subaru Telescope, over a wide range of planetary orbital phases, in order to search for reflected light from the planet.  Data taken on the night of a planetary transit, described in section 2, are well-suited for the H$\\alpha$ search.  The procedure by which we compared the spectra taken at different times is given in section 3, and significantly improved upper limits on exospheric H$\\alpha$ absorption are derived from the data. In section 4 we discuss the physical interpretation and provide a brief summary. ",
        "conclusions": "Fundamentally, our null result corresponds to a maximum number of neutral hydrogen atoms in the $n=2$ state that are present in the exosphere of HD~209458b. Likewise, the detection of Ly$\\alpha$ absorption by Vidal-Madjar et al.\\ (2003) implies a minimum number of neutral hydrogen atoms in the $n=1$ state.  Unfortunately it is difficult to use these results to constrain interesting quantities such as the exospheric temperature and density, because the physical conditions in the exosphere are likely to be quite complex. The column density of neutral hydrogen may vary widely across the surface of the primary star due to evaporation, tidal forces, and radiation pressure (see, e.g., Lecavelier des Etangs et al.\\ 2004). Gas may be present at a variety of temperatures and densities, and may be far from local thermodynamic equilibrium (LTE). [Even in the denser atmosphere, non-LTE effects may be required to explain the unexpectedly weak detection of sodium by Charbonneau et al.\\ (2002), as proposed by Barman et al.\\ (2002) and Fortney et al.\\ (2003).]  Obviously the limited empirical information available is not sufficient by itself to determine a realistic physical exosphere models.  With these caveats, we can use a simple order-of-magnitude argument to provide an upper limit on the column density of $n=2$ neutral hydrogen that exosphere models will need to obey. In order to translate our upper limit on additional H$\\alpha$ absorption into a corresponding upper limit on the H$\\alpha$ equivalent width, we combined the in-transit spectra to form $S_{\\rm in}(\\lambda)$, and we combined the out-of-transit spectra to form $S_{\\rm out}(\\lambda)$.  Then the equivalent width $W_\\lambda$ of any transit absorption feature is \\begin{eqnarray} \\label{eq:eqwid} W_\\lambda \\approx \\displaystyle \\int d\\lambda \\hspace{0.05in} \\left[ \\frac{ S_{\\rm out}(\\lambda) - S_{\\rm in}(\\lambda) } { S_{\\rm out}(\\lambda) }  \\right]. \\end{eqnarray} Performing the integral over the 5~\\AA\\ band described in the previous section, we find $W_{{\\rm H}\\alpha} < 1.7$~m\\AA.  If we further assume that the fraction of the stellar surface covered by the exosphere is $\\Delta A / A \\approx 0.15$ (as estimated from the Ly$\\alpha$ result), then the null result implies that H$\\alpha$ absorption is weak, and the corresponding limit on the column density of $n=2$ atoms ($N_2$) is determined via \\begin{eqnarray} \\label{eq:halpha-width} \\left( \\frac{\\Delta A}{A} \\right) W_{{\\rm H}\\alpha} = \\frac{\\pi e^2}{m c^2} f_{23} \\lambda_{{\\rm H}\\alpha}^2 N_2, \\end{eqnarray} where $f_{23}\\approx 0.641$ is the oscillator strength of the $n=2\\rightarrow 3$ transition.  The resulting limit is $N_2 < 1.0\\times 10^9$~cm$^{-2}$.  In principle, $N_1$ could be estimated from the Ly$\\alpha$ decrement observed during transits, and the ratio $N_2/N_1$ could be related to the H~{\\sc i} excitation temperature. In reality, $N_1$ is highly uncertain because the absorption is probably saturated, and the line profile is unobservable due to interstellar absorption and geocoronal emission.  However, any exosphere model that specifies $N_1$ is constrained by our results to have a maximum excitation temperature $T_{\\rm ex}$, according to the relation \\begin{eqnarray} \\label{eq:n1n2ratio2} \\frac{N_2}{N_1} = \\frac{g_2}{g_1} \\exp \\left(-\\frac{E_2 - E_1}{kT_{\\rm ex}} \\right) = 4\\hspace{0.03in} \\exp\\left(-\\frac{10.2~{\\rm eV}}{kT_{\\rm ex}}\\right). \\end{eqnarray} This constraint appears to be important in the context of recently proposed models in which the exosphere is very hot and hydrodynamically escaping, as argued below.  The temperature of the lower atmosphere of HD~209458b is likely to be near the radiative effective temperature, $T_{\\rm eff} \\approx 10^3$~K, at which the $n=2$ population predicted from equation~(\\ref{eq:n1n2ratio2}) is utterly negligible. But, as argued by Moutou et al.\\ (2001), Lammer et al.\\ (2003), and Lecavelier des Etangs et al.\\ (2004), the exosphere is likely to be significantly hotter than $T_{\\rm eff}$, just as the exosphere of Jupiter is considerably hotter than its lower atmosphere ($\\sim$10$^3$~K vs.\\ 150~K; see, e.g., Atreya 1986). Lammer et al.\\ (2003) found that close-in gas giant planets could have exospheric temperatures up to $10^4$~K due to intense X-ray and extreme-ultraviolet (EUV) irradiation. Depending on the density of neutral hydrogen and other factors, such a hot exosphere could produce detectable H$\\alpha$ absorption in violation of our constraint.  For example, Lecavelier des Etangs et al.\\ (2004) presented 3 models for the thermosphere and exosphere of HD~209458b. In each case they specified the density profile of neutral hydrogen and computed the size and shape of the exosphere as a function of kinetic temperature $T_{\\rm k}$, taking tidal distortion into account. For the specific case of Model A, the density is $n_{\\rm H I}=2\\times 10^9$~cm$^{-3}$ at the base of the thermosphere, and decreases with radius according to the barometric law. We computed the average column density $N_1$ within the annulus reaching from the thermobase to the Roche radius, and then used equation~(\\ref{eq:n1n2ratio2}) to translate our upper limit on $N_2$ into an upper limit on the excitation temperature, finding $T_{\\rm ex} < 8000$~K. By comparison, Lecavelier des Etangs et al.\\ (2004) found that $T_{\\rm k}> 8000$~K in order to produce an escape rate of $10^{10}$~g~s$^{-1}$ [the lower limit inferred by Vidal-Madjar et al.\\ (2003)], and $T_{\\rm k}> 11100$~K through a simulation of EUV heating. Of course, one would not necessarily expect $T_{\\rm ex} \\approx T_{\\rm k}$ in the tenuous thermosphere and exosphere, because of departures from LTE; in general, $T_{\\rm ex}$ is neither an upper bound nor a lower bound on $T_{\\rm k}$. The $n=2$ state may be depopulated when collisions are infrequent, and likewise, whatever process accelerates the gas to $\\sim$100~km~s$^{-1}$ may overpopulate the $n=2$ state via recombination cascades and Ly$\\alpha$ resonant trapping. A full non-LTE exosphere model will need to accommodate the resulting upper limit on the excitation temperature.  In summary, we have used high resolution and high signal-to-noise ratio Subaru-HDS optical spectra to search for excess Balmer H$\\alpha$ absorption during a transit of the extrasolar planet HD~209458b. No excess absorption was detected. Our upper limit is two orders of magnitude below the Ly$\\alpha$ absorption reported by Vidal-Madjar et al.\\ (2003, 2004). It may be difficult to improve upon this limit with ground-based instruments, given the difficulty of correcting for telluric and instrumental spectral variations.  However, there should be no obstacle to improving upon our limit with space-based spectroscopy, and we are aware that such an effort is underway using data from the Hubble Space Telescope (Charbonneau, D., private communication). The current result, and any future refinements, will be useful in the further development of models for the escape of the hot exospheres of close-in gas giant planets.  \\bigskip  We are very grateful to Yutaka Abe, Dave Charbonneau, Bruce Draine, Bob Kurucz, Avi Loeb, Paul Martini, Dimitar Sasselov, and Motohide Tamura, for helpful consultations. This work is based on data from the Subaru Telescope, which is operated by the National Astronomical Observatory of Japan.  We wish to recognize and acknowledge the very significant cultural role and reverence that the summit of Mauna Kea has always had within the indigenous Hawaiian community.  We are most fortunate to have the opportunity to conduct observations from this mountain.  \\bigskip"
    },
    "0404/nucl-th0404089_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-Bethe-Goldstone many-body theory, with realistic two-body and three-body forces. For the description of quark matter, we employ both the MIT bag model with a density dependent bag constant, and the color dielectric model. We calculate the structure of NS interiors with the EOS comprising both phases, and we find that the NS maximum masses are never larger than 1.7 solar masses, no matter the model chosen for describing the pure quark phase. ",
        "introduction": "The appearence of quark matter in the interior of massive neutron stars (NS) is one of the main issues in the physics of these compact objects. Calculations of NS structure, based on a microscopic nucleonic equation of state (EOS), 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 the nucleon cores (dimension $\\approx 0.5\\;\\rm fm$) start to touch each other, and it is hard to imagine that only nucleonic degrees of freedom can play a role. On the contrary, it can be expected that even before reaching these density values, the nucleons start to loose their identity, and quark degrees of freedom are excited at a macroscopic level.  Unfortunately it is not straightforward to predict the relevance of quark degrees of freedom in the interior of NS for the different physical observables, like cooling evolution, glitch characteristics, neutrino emissivity, and so on. In fact, the other NS components can mask the effects coming directly from quark matter. In some cases the properties of quark and nucleonic matter are not very different, and a clear observational signal of the presence of the deconfined phase inside a NS is indeed hard to find.  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 NS. If the quark matter EOS is quite soft, the quark component is expected to appear in NS and to affect appreciably the maximum mass value. In this case the maximum mass is expected to be slightly larger than the observational limit (1.44 solar masses of the so-called Taylor pulsar \\cite{taylor}). The observation of a large NS mass (larger than 2 solar masses) would imply, on the contrary, that the EOS of NS matter is stiff enough to keep the maximum mass at this large values. Purely nucleonic EOS are able to accomodate masses comparable with these large values \\cite{gle,bbb,akma}. Since the presence of non-nucleonic degrees of freedom, like hyperons and quarks, tends usually to soften considerably the EOS with respect to purely nucleonic matter, thus lowering the mass value, their appearence would be incompatible with observations. The large value of the mass could then be explained only if both hyperonic and quark matter EOS are stiffer than expected. In particular, the quark EOS should be assumed to be stiff enough to render the deconfined phase energetically disfavoured.  In this paper we will discuss this issue in detail. Unfortunately, while the microscopic theory of the nucleonic EOS has reached a high degree of sophistication, 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 uncertainity. The best one can do is to compare the predictions of different models and to estimate the uncertainty of the results for the NS matter as well as for the NS structure and mass. In this paper we will use a definite nucleonic EOS, which has been developed on the basis of nuclear matter many-body theory, and two different models for the quark EOS, and compare the results. Confrontation with previous calculations shall also be discussed.  The paper is organized as follows. In section \\ref{s:bhf} we review the determination of the baryonic EOS comprising nucleons and hyperons in the Brueckner-Hartree-Fock approach. Section \\ref{s:qm} concerns the quark matter EOS according to the MIT bag model and the color dielectric model (CDM). In section \\ref{s:res} we present the results regarding neutron star structure combining the baryonic and quark matter EOS for beta-stable nuclear matter. Section \\ref{s:end} contains our conclusions. ",
        "conclusions": "\\label{s:end}  In this article we determine the structure of neutron stars, combining the most recent microscopic baryonic EOS in the BHF approach involving three-body forces and hyperons with different effective models describing the quark matter phase.  Without allowing for the presence of quark matter, the maximum neutron star mass remains below 1.3 solar masses, due to the strong softening effect of the hyperons on the EOS, compensating the repulsive character of nucleonic TBF at high density. The presence of quark matter inside the star is required in order to reach larger maximum masses.  We introduced a density dependent bag parameter $B(\\rho)$ in the MIT model in order to explore the maximum NS mass that can be reached in this approach. We compare with calculations using a fixed bag constant and using the color dielectric model. Joining the corresponding EOS with the baryonic one, all three quark models yield maximum masses in the range $(1.5,\\ldots,1.6)\\;M_\\odot$, while predicting slightly different radii.  Our results for the maximum masses are in line with other recent calculations of neutron star properties employing various phenomenological relativistic mean field nuclear EOS together with either effective mass bag model \\cite{bag} or Nambu-Jona-Lasinio model \\cite{njl} EOS for quark matter.  The value of the maximum mass of neutron stars obtained according to our analysis appears rather robust with respect to the uncertainties of the nuclear and the quark matter EOS. Therefore, the experimental observation of a very heavy ($M \\gtrsim 1.7 M_\\odot$) neutron star, as claimed recently by some groups \\cite{kaaret} ($M \\approx 2.2\\;M_\\odot$), if confirmed, would suggest that either serious problems are present for the current theoretical modelling of the high-density phase of nuclear matter, or that the assumptions about the phase transition between hadron and quark phase are substantially wrong. In both cases, one can expect a well defined hint on the high density nuclear matter EOS."
    },
    "0404/astro-ph0404432_arXiv.txt": {
        "abstract": "Neutrino-matter cross sections and interaction rates are central to the core-collapse supernova phenomenon and, very likely, to the viability of the explosion mechanism itself.  In this paper, we describe the major neutrino scattering, absorption, and production processes that together influence the outcome of core collapse and the cooling of protoneutron stars.  One focus is on energy redistribution and many-body physics, but our major goal is to provide a useful resource for those interested in supernova neutrino microphysics. ",
        "introduction": "Supernova explosions are one major means by which elements are injected into the interstellar medium and, hence, into subsequent generations of stars.  Therefore, supernovae are central to the chemical evolution and progressive enrichment of the universe. Most supernova explosions are the outcome of the dynamical collapse of the core of a massive star as it dies.  Collapse creates high temperatures ($> 1$ MeV) and densities ($10^7$ g cm$^{-3} < \\rho < 10^{15}$ g cm$^{-3}$) and produces (``after the dust settles\") either a neutron star or a black hole.  Under such extreme thermodynamic conditions, neutrinos are produced in abundance. The mechanism of core-collapse supernovae is thought to depend upon the transfer of energy from the inner core to the outer mantle of the iron core of the massive star.  Neutrinos seem to be the mediators of this energy transfer. Therefore, to fully understand core-collapse supernova explosions one must have a firm handle on the physics of neutrino production, absorption, and scattering.  In this paper, we summarize the neutrino-matter cross sections and the neutrino production rates in the core-collapse context. Some of this discussion can already be found in Burrows (2001). We do not attempt to explain the hydrodynamics of supernova explosions, but do try to present the relevant neutrino processes that play a role.  For the former, the reader is referred to \\cite{nature}, \\cite{bhf_1995}, \\cite{lieben2001}, \\cite{lieben20012}, \\cite{rampp20022}, \\cite{buras2003}, and Thompson, Burrows, \\& Pinto (2003).  In \\S\\ref{stimabs}, we present a physical derivation of stimulated absorption and then in \\S\\ref{cross6} we summarize the basic neutrino-matter cross sections.  In \\S\\ref{section:inelastic}, we discuss the neutrino-electron scattering kernel, along with a simple treatment of the collision integral.   In \\S\\ref{freegas}, we provide the relativisitic formalism for inelastic neutrino-electron and neutrino-nucleon scattering processes and energy redistribution for non-interacting nucleons.  This is followed in \\S\\ref{strongandeandm} with a discussion of the alternate, more powerful, formalism for determining differential interaction rates and redistribution in the many-body context, namely that of dynamical structure factors.  The role of strong and electromagnetic interactions between nucleons and leptons is explored, as well as collective excitations of the medium. Source terms for electron-positron annihilation (\\S\\ref{eplus}), neutrino-anti-neutrino annihilation (\\S\\ref{paira}), and nucleon-nucleon bremsstrahlung (\\S\\ref{bremsst}) cap off our review of the major processes of relevance in core-collapse simulations. ",
        "conclusions": "The processes that have been described in this paper are essential elements of the neutrino-driven supernova explosion mechanism.  Coupling these with radiation-hydrodynamics codes, an equation of state, beta-decay and electron capture microphysics, and nuclear rates, one explores the viability of various scenarios for the explosion of the cores of massive stars (Liebend\\\"orfer et al. 2001ab; Rampp \\& Janka 2000,2002).  Recently, Thompson, Burrows, \\& Pinto (2003) have incorporated this neutrino microphysics into simulations of 1D (spherical) core collapse and have investigated the effects on the dynamics, luminosities, and emergent spectra of weak magnetism/recoil, nucleon-nucleon bremsstrahlung, inelastic neutrino-electron scattering, and a host of the cross section corrections described above. The character of the spectra reflect the opacities and sources.  In particular, the energy hardness hierarchy from $\\nu_e$ (softer) to $\\nu_{\\mu}$ (harder) neutrinos is clearly manifest, as is the distinction between the $\\nu_e$ pre-breakout and post-breakout spectra.  To date, none of the detailed 1D simulations that have been performed explodes and it may be that multi-dimensional effects play a pivotal role in the explosion mechanism (Herant et al.~1994; Burrows, Hayes, \\& Fryxell 1995; Janka \\& M\\\"uller 1996; Fryer et al. 1999; Fryer \\& Warren 2002). Be that as it may, an understanding of neutrino-matter interactions remains central to unraveling one of the key mysteries of the nuclear universe in which we live."
    },
    "0404/astro-ph0404604_arXiv.txt": {
        "abstract": "Abell 1758 was classified as a single rich cluster of galaxies by Abell, but a ROSAT observation showed that this system consists of two distinct clusters (A1758N and A1758S) separated by approximately $8\\arcmin$ (a projected separation of 2~Mpc in the rest frame of the clusters). Only a few galaxy redshifts have been published for these two clusters, but the redshift of the Fe lines in the Chandra and XMM-Newton spectra shows that the recessional velocities of A1758N and A1758S are within 2,100~km~s$^{-1}$. Thus, these two clusters most likely form a gravitationally bound system, but our imaging and spectroscopic analyses of the X-ray data do not reveal any sign of interaction between the two clusters. The Chandra and XMM-Newton observations show that A1758N and A1758S are both undergoing major mergers.  A1758N is in the late stages of a large impact parameter merger between two 7~keV clusters.  The two remnant cores have a projected separation of 800~kpc.  Based on the measured pressure jumps preceding the two cores, they are receding from one another at less than 1,600~km~s$^{-1}$.  The two cores are surrounded by hotter gas ($\\mathrm{kT}=9$--12~keV) that was probably shock heated during the early stages of the merger.  The gas entropy in the two remnant cores is comparable with the central entropy observed in dynamically relaxed clusters, indicating that the merger-induced shocks stalled as they tried to penetrate the high pressure cores of the two merging systems. Each core also has a wake of low entropy gas indicating that this gas was ram pressure stripped without being strongly shocked.  A1758S is undergoing a more symmetric (lower impact parameter) merger between two 5 keV clusters.  The two remnant cores are nearly coincident as seen in projection on the sky.  The two cores are surrounded by hotter gas (9--11 keV) which was probably shock heated during the merger.  Based on the pressure jumps preceding the two cores, they must have a relative velocity of less than 1,400~km~s$^{-1}$. Unlike A1758N, there is no evidence for wakes of low entropy gas. ",
        "introduction": "Chandra observations over the past few years have shown that clusters of galaxies are dynamically complex systems.  These observations have shown that the hot gas in clusters is frequently perturbed by cluster mergers and nuclear outbursts and motion of the central dominant galaxy. Cluster mergers generate hydrodynamic shocks, cold fronts (contact discontinuities), and filaments (Markevitch et al.\\ 2000; Vikhlinin, Markevitch, \\& Murray 2001a; Markevitch et al.\\ 2002; Mazzotta et al.\\ 2002; Mazzotta, Fusco-Femiano, \\& Vikhlinin 2002; Kempner, Sarazin, \\& Ricker 2002). Residual motion or ``sloshing'' of the central dominant galaxy in clusters can also produce cold fronts and possibly wakes of stripped or cooled gas (Fabian et al.\\ 2001; Markevitch, Vikhlinin, \\& Mazzotta 2001). These discoveries have helped illuminate the roles of magnetohydrodynamics and transport processes during cluster mergers. The observed sharpness of cold fronts shows that thermal conduction is highly suppressed in these regions (Ettori \\& Fabian 2000; Vikhlinin, Markevitch, \\& Murray 2001a). The robustness of merging cores demonstrates that the growth of hydrodynamic instabilities is also suppressed, possibly by the compression of ambient magnetic fields along cold fronts and the subsequent increase in magnetic surface tension (Vikhlinin, Markevitch, \\& Murray 2001b). A full understanding of recent Chandra observations of clusters provides a significant challenge for future numerical simulations, which must resolve structure on scales of a few kpc and include the effects of magnetic fields.  While Abell 1758 was classified as a single cluster by Abell (1958), Rosat images show that there are two distinct clusters (A1758N and A1758S) separated by approximately $8\\arcmin$ (Rizza et al.\\ 1998).  Abell 1758N is a hot ($\\rm{kT} \\sim$~9--10~keV; Mushotzky \\& Scharf 1997) and X-ray luminous cluster ($L_{bol}=2.9 \\times 10^{45}$~ergs~s$^{-1}$; David, Jones, \\& Forman 1999). No temperature information is available for A1758S in the literature, but it is only slightly less luminous than A1758N ($L_{bol}=2.0 \\times 10^{45}$~ergs~s$^{-1}$; David, Jones, \\& Forman 1999). The Rosat HRI image shows that both A1758N and A1758S are highly disturbed systems (Rizza et al.\\ 1998).  A1758N has two X-ray peaks separated by approximately $3\\arcmin$, while the centroid of A1758S is offset from the centroid of the large scale emission by approximately $1\\arcmin$.  A1758N hosts one of the most powerful NAT radio sources known (O'Dea \\& Owen 1985, Rizza et al.\\ 1998). Head-tail and NAT radio galaxies are most common in clusters with perturbed X-ray morphologies and are probably produced as radio galaxies move through cluster atmospheres at high velocities (Burns et al.\\ 1994).   One of the X-ray peaks in A1758N is centered on a diffuse radio halo (Kempner \\& Sarazin 2001), which is further evidence for a recent major merger.  \\begin{figure*} \\plotone{f1.eps} \\caption{Adaptively smoothed, exposure corrected mosaic of the 3 XMM-Newton EPIC cameras.  The image spans $16\\arcmin$ on a side (4.1~Mpc at the redshift of the cluster). }\\label{fig:xmm_mosaic} \\end{figure*}  In this paper, we present Chandra and XMM-Newton observations of A1758. By utilizing Chandra's high spatial resolution and XMM-Newton's large throughput, we can search for signs of interaction between A1758N and A1758S and investigate the present status of the mergers in the two systems. This paper is organized in the following manner.  In $\\S 2$ we discuss the details of our Chandra and XMM-Newton data analysis.   Section 3 contains a discussion about the large scale X-ray properties of A1758 and a search for any interaction between A1758N and A1758S.  We discuss the dynamic states of A1758N and A1758S separately in sections 4 and 5.  Our proposed merger scenarios for the two clusters are summarized in $\\S 6$. We assume $\\mathrm{H_0}=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{M}=0.3$, and $\\Omega_{\\Lambda}=0.7$ throughout the paper. At a redshift of 0.279, the luminosity distance to A1758 is 1,430~Mpc, and $1\\arcsec = 4.24$~kpc. ",
        "conclusions": "The optically rich cluster Abell 1758 is located in a dynamically active region that contains two hot and X-ray luminous clusters (A1758N and A1758S) with recessional velocities within 2,100~km~s$^{-1}$ and within a projected separation of 2~Mpc. These two systems are most likely gravitationally bound, but there is no evidence for any interaction between the two gaseous atmospheres in the X-ray data. The X-ray morphologies of A1758N and A1758S are very complex indicating that both systems are undergoing major mergers. A1758N is in the later stages of a large impact parameter merger between two 7~keV clusters, while A1758S is in the earlier stages of a merger between two 5~keV clusters.  A1758N contains two low entropy cores that were not strongly shocked during the merger.  The robustness of the two cores in A1758N indicates that the core of the dark matter distribution in the two clusters is smaller than approximately 100 kpc. Each of these cores has a preceding cold front (separating the low entropy cores and shock heated gas) and a wake of low entropy gas.  The surface brightness profiles of the gas preceding the two cold fronts in A1758N are well represented by broken power-laws.  The slope of the outer power-law is consistent with the slope typically found in the outskirts of rich clusters. The gas temperature preceding the NW subcluster in A1758N increases with increasing distance from the cold front. A different trend would be observed if the break in the surface brightness was produced by a shock.  Based on the X-ray morphology of A1758N, we propose that both systems have undergone significant deflections during their merger.  Thus, these systems would have propagated through a broad range of ambient gas pressures producing a broad range of shock strengths.  The lower temperature gas located directly in front of the NW subcluster may have been shocked as the two cores passed close to one another while the gas pressure was high.  As the two systems receded from one another into lower pressure gas, the shock strength would increase producing higher post-shocked temperatures. The chaotic entropy distribution of the gas in A1758N will drive a significant redistribution of the energy deposited by the merger shock via convection.  A1758S is undergoing a major merger between two 5 keV clusters.  The X-ray morphology of A1758S is very similar to that seen in A2142 (Markevitch et al.\\ 2000). The cores of the two subclusters are nearly coincident on the plane of the sky. The X-ray morphology of A1758S is more symmetric than the morphology in A1758N, suggesting that the merger in A1758S is more nearly head-on.  Both subclusters in A1758S are preceded by cold fronts, which in turn are surrounding by hotter gas that was probably shock heated during the merger.  The upper limits on the pressure jump across the two cold fronts indicate that the relative velocity between the two subclusters must be less than 1,400~km~s$^{-1}$. The two cores in A1758S have radii of 200 kpc and 350 kpc, compared to 80 and 115 kpc in A1758N, indicating that these cores have not experienced as much ram pressure stripping.  This is consistent with A1758S being in the earlier stages of a merger, especially if the merger in A1758S is nearly head-on.  There is also no evidence in the XMM-Newton data for low entropy wakes in A1758S, which probably indicates that all of the gas that has been stripped from the two cores was shocked while it was being stripped. This is also reasonable since ram pressure stripping should be the most violent during the early stages of a merger while the outer, lower pressure gas is being stripped.  Abell 1758 is a remarkable system with 4 subclusters hotter than 5 keV within a region of 2~Mpc. Using the simulations of Randall, Sarazin, \\& Ricker (2002) as a guide, the final virialized temperature of A1758 should be about 12 keV. We are thus witnessing the formation of a cluster which will rank among the hottest and most massive clusters found in the local universe."
    },
    "0404/astro-ph0404097_arXiv.txt": {
        "abstract": "The binary system $\\gamma^2$ Velorum (WC8+O7.5) contains the nearest known Wolf-Rayet star to the Sun, at a distance of 258$_{-31}^{+41}$~pc. Its strong radio emission shows evidence for a partially absorbed nonthermal component, which has been interpreted as synchrotron emission from electrons accelerated in the colliding wind region. Inverse Compton cooling of these electrons in the intense UV radiation field from the O-type companion star could produce a significant hard X-ray and $\\gamma$-ray emission, whose flux depends on the ratio of the energy densities of magnetic to seed photon fields. The Vela region was observed with the {\\it INTEGRAL} satellite in 2003, as part of the Core Programme. No signals from $\\gamma^2$ Vel are detected in the images obtained with the IBIS/ISGRI coded aperture instrument in the energy ranges 20--40 and 40--80 keV. From the derived 3$\\sigma$ upper limits, we show that the average magnetic field near the region of stellar wind collision should be relatively high, $B$$\\gsim$1~G. The high-energy emission of $\\gamma^2$ Vel might be detected with the forthcoming {\\it GLAST} experiment. ",
        "introduction": "Radio continuum measurements of Wolf-Rayet (WR) stars have revealed at least ten sources with nonthermal (NT) emission, which presumably originates from an interaction between the WR stellar wind and the wind from a massive companion star (Chapman et al. 1999, and references therein). The observed NT radio spectra and luminosities are well explained as synchrotron emission from relativistic electrons accelerated at strong shocks in the colliding wind region (Eichler and Usov 1993).  Benaglia and Romero (2003) have recently studied the production of $\\gamma$-ray emission in the binary systems WR~140, WR~146 and WR~147. They showed that inverse Compton (IC) scattering of the accelerated electrons in the strong photon fields of the early-type stars could produce $\\gamma$-ray fluxes above the {\\it INTEGRAL}/IBIS continuum sensitivity. In the case of WR~140, they also showed that the expected high-energy emission can account for the unidentified EGRET source 3EG J2022+4317.  In the present paper, we consider the relatively close (i.e. with short orbital period $P_{orb}$=78.53~days) binary system $\\gamma^2$ Velorum (=WR~11, van der Hucht 2001). With a distance of 258$_{-31}^{+41}$~pc, as determined from {\\it Hipparcos} parallax measurements, this WC8+O7.5 binary contains the nearest WR star to the Sun. Radio and millimetre observations of $\\gamma^2$ Vel essentially revealed the strong thermal emission from the WR ionized wind (e.g. Leitherer et al. 1997). However, Chapman et al. (1999) found a significant steepening of the radio spectral index between 3 and 20~cm, which they interpreted as evidence for a highly attenuated NT component originating in the colliding wind region, well within the radio photosphere of the WR star (Chapman et al. 1999). This contrasts with the wider binary systems studied by Benaglia and Romero (2003), for which the free-free absorption of the NT radio emission by the WR wind is less important.  The existence of strong shocks in the colliding wind region of $\\gamma^2$ Vel is also supported by X-ray observations (Skinner et al. 2001, and references therein). In particular, {\\it ROSAT} and {\\it ASCA} data have revealed a hot plasma emission ($kT$$\\gsim$1~keV) with a strong phase-locked variability, which was interpreted as a colliding wind shock emission originating deep within the dense and opaque WR wind, and showing significantly less photoelectric absorption at orbital phases when the cavity around the O-type companion star crosses the line of sight (Rauw et al. 2000, and references therein).  The model for the high-energy emission from the stellar wind collisions in $\\gamma^2$ Vel is described in the next section. In Sect.~3, we present the {\\it INTEGRAL}/IBIS observations. The results are discussed in Sect.~4. ",
        "conclusions": "The high-energy IC radiation of $\\gamma^2$ Vel, which should accompany the synchrotron radio emission revealed by the ATCA data, is not detected with IBIS/ISGRI. A possible explanation is that the average magnetic field near the region of stellar wind collision is relatively high. Assuming that the intrinsic NT radio emission does not significantly vary over the orbital cycle, we obtain from the IBIS/ISGRI 3$\\sigma$ upper limits $B$$\\gsim$1~G (see Fig.~2). However, a significant variation of the NT radio emission is observed in the well-studied binary WR~140. From 8 years of monitoring the radio flux of WR~140 with the VLA, White and Becker (1995) have shown that the intrinsic NT component (before attenuation by free-free absorption) seems to become weaker near periastron, which is not expected in the model of spherically symmetric colliding winds. They proposed a new model in which the WR wind is strongly enhanced in the star equatorial disk. Such an effect could exist in the eccentric binary $\\gamma^2$ Vel as well, and radio and $\\gamma$-ray observations at the same orbital phases are required to specify the parameters of stellar wind collisions.  The forthcoming $\\gamma$-ray telescope {\\it GLAST} might detect the $\\gsim$20~MeV IC emission of $\\gamma^2$ Vel, provided that the proximity of the very bright Vela pulsar does not induce a contamination problem. For an estimated 3$\\sigma$ sensitivity $S_{ph}$$\\times$$E_{ph}^2$$\\sim$0.4~eV~cm$^{-2}$~s$^{-1}$ at 100~MeV (for one year of observation and $\\Delta E_{ph}$=$E_{ph}$), the IC counterpart of the NT radio emission could be observed if the average magnetic field near the colliding wind zone is $B$$\\lsim$10~G (see Fig.~2). In the case of a positive detection, it casts no doubt that $\\gamma^2$ Vel would become a key object for understanding the processes of particle acceleration in binaries of early-type stars."
    },
    "0404/astro-ph0404268_arXiv.txt": {
        "abstract": "We report the discovery of a new dwarf spheroidal satellite of M31, Andromeda IX,  based on resolved stellar photometry from the Sloan Digital Sky Survey (SDSS).  Using both SDSS and public archival data we have estimated its distance and other physical properties, and compared these to the properties of a previously known dwarf spheroidal companion, Andromeda V, also observed by SDSS. Andromeda IX is the lowest surface brightness galaxy found to date ($\\mu_{V,0} \\sim 26.8\\, {\\rm mag~arcsec}^{-2}$), and at the distance we estimate from the position of the tip of Andromeda IX's red giant branch, $(m - M)_0 \\sim 24.5$ (805 kpc), Andromeda IX would also be the faintest galaxy known ($M_V  \\sim -8.3$). ",
        "introduction": "\\label{txt:intro}  Hierarchical cold dark matter (CDM) models, while successful at large scales, predict too many low-mass dark subhalos to be consistent with the observed abundance of dwarf galaxies, by at least 1 -- 2 orders of magnitude \\citep[e.g.,][]{klyp99,moor99,bens02b}. This problem can be at least qualitatively addressed if star formation in low mass subsystems were inhibited, for example by photoionization in the early universe; this could lead to galaxy luminosity functions with shallow faint-end slopes at the present day \\citep[e.g.,][]{some02,bens02a}, with observed satellites embedded in much larger, more massive dark subhalos \\citep{stoe02}.   Observational efforts to constrain the form of the luminosity function for faint galaxies are hindered by the extremely low surface brightnesses expected of such galaxies \\citep[$\\mu_{V} \\gtrsim 26\\, {\\rm mag~arcsec}^{-2}$, e.g.,][]{cald99,bens02b}. Conducting a comprehensive ground-based  survey for such diffuse objects would be extremely difficult even in nearby galaxy groups. Fortunately, in the case of the Local Group (LG), galaxies can be resolved into stars, allowing one to reach much fainter limits \\citep[e.g.,][]{ferg02} and potentially place strong constraints on both the LG luminosity function and galactic formation models.   In this letter, we report the discovery, using resolved stellar data from the Sloan Digital Sky Survey (SDSS), of a new dwarf spheroidal companion to M31, one which is the lowest luminosity, lowest surface brightness galaxy found to date. For this work, we have assumed a distance to M31 of 760\\,kpc \\citep[$(m - M)_0 = 24.4$;][]{vand99}. ",
        "conclusions": "And IX is the lowest surface brightness, lowest luminosity galaxy found to date, with a metallicity comparable to the least chemically-evolved stellar systems in the LG; its distance from the center of M31, $\\sim 45$ kpc, places it well within, e.g., the virial radius (272 kpc) assumed by \\citet{bens02b} in their models, indicating that it is in all likelihood a satellite of that galaxy. At an absolute magnitude of $M_V \\sim -8.3$ ($\\sim 2 \\times 10^5$ L$_{\\odot}$), And IX is comparable in luminosity to many globular clusters, although two orders of magnitude larger ($r \\gtrsim 500 {\\rm pc}$) .   Given that current hierarchical CDM models for galaxy formation still generate a satellite galaxy luminosity function which rises at low luminosities (albeit with a shallower slope than some earlier models), the discovery of And IX raises some interesting questions. Is And IX a rarity, one of a small number of such extremely low luminosity galaxies in the LG, with the vast majority of the predicted large numbers of low mass subhalos surviving only as dark matter? Or could And IX be the tip of the iceberg, representative of a large population of low luminosity dwarf satellites that have remained undetected because of their extremely low surface brightnesses \\citep[e.g.,][]{bens02b}?   The discovery of And IX \\citep[like that of And NE,][]{zuck04a} also highlights the capabilities of large-scale imaging surveys with uniform photometry, like SDSS, for detecting incredibly diffuse stellar structures in the LG using their resolved stellar components \\citep[see][for a detailed discussion]{will02}. If And IX is but one of a large population of extremely faint M31 satellites, it is quite likely that further analysis of the SDSS data will yield more such objects, bringing the observed galaxy luminosity function into better agreement with model predictions."
    },
    "0404/astro-ph0404024_arXiv.txt": {
        "abstract": "We report the discovery of the scattered emission from a hidden broad-line region (BLR) in a Seyfert 2 galaxy, Mrk 573, based on our recent spectropolarimetric observation performed at the Subaru Telescope. This object has been regarded as a type 2 AGN without a hidden BLR by the previous observations. However, our high quality spectrum of the polarized flux of Mrk 573 shows prominent broad ($\\sim$3000 km s$^{-1}$) H$\\alpha$ emission, broad weak H$\\beta$ emission, and subtle Fe {\\sc ii} multiplet emission. Our new detection of these indications for the presence of the hidden BLR in the nucleus of Mrk 573 is thought to be owing to the high signal-to-noise ratio of our data, but the possibility of a time variation of the scattered BLR emission is also mentioned. Some diagnostic quantities such as the $IRAS$ color, the radio power, and the line ratio of the emission from the narrow-line region of Mrk 573 are consistent with the distributions of such quantities of type 2 AGNs with a hidden BLR. Mrk 573 is thought to be an object whose level of the AGN activity is the weakest among the type 2 AGNs with a hidden BLR. In terms of the systematic differences between the type 2 AGNs with and without a hidden BLR, we briefly comment on an interesting Seyfert 2 galaxy, Mrk 266SW, which may possess a hidden BLR but has been treated as a type 2 AGNs without a hidden BLR. ",
        "introduction": "\\begin{deluxetable}{lcccc} \\tablenum{1} \\tablecaption{Polarization Properties of Stars near the Line-of-Sight of Mrk 573\\tablenotemark{a}} \\tablewidth{0pt} \\tablehead{ \\colhead{Star Name} & \\colhead{Separation} & \\colhead{Distance} & \\colhead{Pol. Degree} & \\colhead{Pol. Angle} \\\\ \\colhead{} & \\colhead{(arcmin)} & \\colhead{(parsec)} & \\colhead{(\\%)} & \\colhead{(deg)} } \\startdata HD  9740 &  103  &  436  & 0.33 & 122.5 \\nl HD 10441 &  155  &  316  & 0.30 & 128.4 \\nl adopted  &\\nodata&\\nodata& 0.32 & 125 \\enddata \\tablenotetext{a}{Data are taken from Heiles (2000).} \\end{deluxetable}  After the discovery of scattered broad permitted lines in the polarized spectra of a Seyfert 2 galaxy (hereafter Sy2) NGC 1068 (Antonucci \\& Miller 1985), many attempts have been made to detect the polarized broad lines in type-2 active galactic nuclei (AGNs) up to now (e.g., Miller \\& Goodrich 1990; Tran, Miller, \\& Kay 1992; Young et al. 1993; Tran, Cohen, \\& Goodrich 1995; Tran 1995; Young et al. 1996; Kay \\& Moran 1998; Barth, Filippenko, \\& Moran 1999a, 1999b; Tran, Cohen, \\& Villar-Martin 2000; Kishimoto et al. 2001; Tran 2001; Lumsden et al. 2001). This is because the presence of scattered broad permitted lines is promising evidence for the AGN unified model, in which both the type-1 and type-2 AGNs possess a broad-line region (BLR) in their nucleus (see Antonucci 1993 for a review).  Although the scattered BLR emission has been found in many type-2 AGNs, it is now recognized that Sy2s do not always exhibit broad lines in the polarized spectra (see, e.g., Tran 2001). Does this suggest that not all Sy2s possess a BLR in their nucleus? This is very important issue because the presence of the two distinct populations of Sy2s, i.e., Sy2s with and without a BLR, is inconsistent to the simple AGN unified model. Heisler, Lumsden, \\& Bailey (1997) proposed that this dichotomy can be understood in the framework of the unified model if the scattering region resides very close to the nucleus and its visibility depends on the viewing angle (see also, e.g., Taniguchi \\& Anabuki 1999). On the other hand, Tran (2001) reported that the amount of the obscuration toward the central engine is indistinguishable between the Sy2s with and without polarized BLR emission (see also Alexander 2001). This suggests the presence of AGNs without BLRs (see, e.g., Heckman et al. 1995; Dultzin-Hacyan et al. 1999; Gu, Maiolino, \\& Dultzin-Hacyan 2001), which contradicts the current simple unified model. Investigating this issue further is crucially important not only to examine the AGN unified model but also to understand the nature of AGN phenomena themselves.  In order to discuss this issue, we should recognize correctly which Sy2 possesses a hidden BLR and which Sy2 does not. Here we report a clear detection of the hidden BLR in a Sy2, Mrk 573, which has been regarded as a Sy2 {\\it without} a hidden BLR (Tran 2001, 2003; see also Kay 1994). Its heliocentric radial velocity is 5156 $\\pm$ 90 km s$^{-1}$ (Whittle et al. 1988), giving a projected linear scale of 0.33 $h_{75}^{-1}$ kpc for 1 arcsec. ",
        "conclusions": "To examine the AGN unified model is one of the most interesting issues toward understanding the nature of the AGN phenomena. Thus various comparative studies between Sy2s with and without the polarized BLR emission have been performed up to now (e.g., Heisler et al. 1997; Awaki et al. 2000; Gu et al. 2001; Thean et al. 2001; Lumsden et al. 2001; Tran 2001, 2003; Gu \\& Huang 2002). In such studies, Mrk 573 has been regarded as a Sy2 without a hidden BLR. However it is now evident that Mrk 573 surely possesses a hidden BLR in its nucleus. If this kind of misclassification occurs frequently, the previous comparative studies would become rather senseless. Therefore the reason of the misclassification should be discussed here.  Kay (1994) presented the results of the spectropolarimetric observations for Mrk 573 which were performed in November 1987 -- December 1989. Although it was mentioned that the H$\\beta$ and the H$\\gamma$ emission in the polarized flux spectrum of Mrk 573 might be slightly broader than those in the total flux spectrum, it could not be concluded because the observation did not cover the H$\\alpha$ wavelength range. It is therefore crucially important for exploring the hidden BLRs to investigate the profiles of the polarized H$\\alpha$ emission, which is the easiest spectral feature to access the hidden BLR. It should be mentioned that, however, there may be a possibility that the scattered BLR emission is temporary variable. Although there are some reports that the scattered BLR emission of type-1 AGNs varies in a few years (e.g., Young et al. 1999; Smith et al. 2002; Nagao et al. 2004), there is no report of finding a significant temporal variation of the scattered broad BLR emission in Sy2s, so far. This is thought to be partly because the scattered photons cannot be observed due to the obscuration by a dusty tori if the scattering region resides at very close to the nucleus, as for Sy2s. Thus we can see only the polarized light scattered at far from the nucleus with a distance larger than the scale height of dusty tori, $\\sim$1 pc or more (e.g., Taniguchi \\& Murayama 1998). Therefore, we cannot detect the temporal variation of scattered BLR emission of Sy2s unless we monitor the spectropolarimetric data for several years at least. As for Mrk 573, the comparison of the spectropolarimetric data obtained at epochs separated $\\sim$15 years may enable us to access the temporal properties of the scattered BLR emission. If this is the case, most of the polarization of the hidden BLR in Mrk 573 is caused at the region within several pc from the nucleus where the obscuration by the dusty torus is not significant.  \\begin{figure*} \\epsscale{1.18} \\plotone{Nagao.fig05.eps} \\caption{ Diagnostic diagram of the ratio of the radio 20cm flux density to the $IRAS$ 60$\\mu$m flux vs. the $IRAS f_{25}/f_{60}$ color. The data are taken from Tran (2003). Filled and open circles denote the data of Sy2s with and without a hidden BLR, respectively. The data of Mrk 573 and Mrk 266SW are shown by a filled square and a thick circle, respectively. \\label{fig5}} \\end{figure*}  Then we discuss how our finding of the hidden BLR in Mrk 573 affects the previous studies. It is known that Sy2s with a hidden BLR tend to show hotter MIR colors (i.e., higher $f_{25}/f_{60}$ ratios where $f_{25}$ and $_{60}$ are $IRAS$ 25$\\mu$m and 60$\\mu$m fluxes, respectively) than Sy2s without a hidden BLR (e.g., Heisler et al. 1997; Gu et al. 2001; Tran 2001, 2003). Although Heisler et al. (1997) interpreted this tendency as a difference of viewing angles toward a dusty torus, Alexander (2001) reported that the $f_{25}/f_{60}$ ratio is not a reliable indicator for the viewing angle (see also Murayama, Mouri, \\& Taniguchi 2000); rather it denotes the relative strength of the AGN activity compared to the nuclear star-forming activity. In order to compare the relative strength of the AGN activity between Sy2s with and without a hidden BLR, a diagnostic diagram which consists of the $f_{25}/f_{60}$ ratio and the $S_{20}/f_{60}$ ratio (where $S_{20}$ is a radio 20 cm flux density) has been used since the $S_{20}/f_{60}$ ratio is also expected to be a good indicator for the relative AGN activity strength (e.g., Tran 2001, 2003). In Figure 5, we plot the data of Tran (2003) on the diagnostic diagram of $f_{25}/f_{60}$ and $S_{20}/f_{60}$. This is basically the same as Figure 1 of Tran (2001) (see also Figure 4 of Tran 2003), but the data of Mrk 573 (and Mrk 266SW, which is discussed below) is explicitly exhibited. We can see that the data of Mrk 573 is located the edge of the correlation between the two diagnostic flux ratios for the Sy2s with a hidden BLR. Actually the $f_{25}/f_{60}$ ratio of Mrk 573 is lower than all of the Sy2s with a hidden BLR in the sample of Tran (2003), i.e., $f_{25}/f_{60} = 0.23$. Despite the low ratios of $f_{25}/f_{60}$ and $S_{20}/f_{60}$, Mrk 573 is thought not to be a peculiar object as a Sy2 with a hidden BLR because its data is consistent with the correlation between the two diagnostic flux ratios for the Sy2s with a hidden BLR. The same conclusion can be seen in other diagnostic diagram; in Figure 6, we show the diagnostic diagram of $f_{25}/f_{60}$ and [O {\\sc iii}]$\\lambda$5007/H$\\beta$ in which the data of Tran (2003) are plotted. As reported by Tran (2001), there is a statistically significant difference in the distribution of the [O {\\sc iii}]$\\lambda$5007/H$\\beta$ flux ratio between Sy2s with and without a hidden BLR (see also Tran 2003). Again the data of Mrk 573 is consistent with the data of the Sy2s with a hidden BLR though it is located at the edge of the distribution of the data of Sy2s with a hidden BLR in this diagram. Note that it may be a part of the reason for the undetection of the scattered BLR emission of Mrk 573 by the previous observations that the strength of the AGN activity of Mrk 573 is very weak as inferred by the diagnostic diagrams in Figures 5 and 6.  Our study suggests that high-quality spectropolarimetric observations are crucial for dividing Sy2s between those with and without a hidden BLR correctly. In terms of this viewpoint, we briefly comment on an interesting Sy2, Mrk 266SW. This Sy2 has been also treated as a Sy2 without a hidden BLR (e.g., Tran 2001, 2003) but was suspected to possess the hidden BLR by Kay (1994), as well as Mrk 573. Being different from Mrk 573, however, the diagnostic quantities of Mrk 266SW are not similar to those of the Sy2s with a hidden BLR. As seen in Figures 5 and 6, Mrk 266SW shows lower [O {\\sc iii}]$\\lambda$5007/H$\\beta$ flux ratio than all of the Sy2s with a hidden BLR in the sample of Tran (2003), and shows cooler $IRAS$ color than Mrk 573. The data of Mrk 266SW does not follow the correlation for the Sy2s with a hidden BLR seen in Figure 5. Also in the diagnostic diagram of [O {\\sc iii}]$\\lambda$5007/H$\\beta$ vs. $f_{25}/f_{60}$, the data of Mrk 266SW is far from the data distribution of the Sy2s with a hidden BLR but is consistent to the Sy2s without a hidden BLR. Thus we can recognize that the general properties of Mrk 266SW is consistent with so-called ``pure Sy2s'', i.e., Sy2s without a BLR in its nucleus, not with Sy2s with a BLR. In order to understand the nature of the Sy2 populations, performing a high-quality spectropolarimetry for Mrk 266SW seems important to examine whether or not Mrk 266SW possesses a hidden BLR, because Mrk 266SW may be an atypical object if possesses a hidden BLR in its nucleus.  \\begin{figure*} \\epsscale{1.18} \\plotone{Nagao.fig06.eps} \\caption{ Same as Figure 5 but for other diagnostics; the [O {\\sc iii}]$\\lambda$5007/H$\\beta$ flux ratio vs. the $IRAS f_{25}/f_{60}$ color. \\label{fig6}} \\end{figure*}"
    },
    "0404/astro-ph0404212_arXiv.txt": {
        "abstract": "Recent results have shown that many of the known extrasolar planetary systems contain regions which are stable for both Earth-mass and Saturn-mass planets. Here we simulate the formation of terrestrial planets in four planetary systems -- 55 Cancri, HD 38529, HD 37124, and HD 74156 -- under the assumption that these systems of giant planets are complete and that their orbits are well-determined.  Assuming the giant planets formed and migrated quickly, then terrestrial planets may form from a second generation of planetesimals.  In each case, Moon- to Mars-sized planetary embryos are placed in between the giant planets and evolved for 100 Myr.  We find that planets form relatively easily in 55 Cnc, with masses up to 0.6 Earth masses and in some cases substantial water contents and orbits in the habitable zone.  HD 38529 is likely to support an asteroid belt but no terrestrial planets of significant mass.  No terrestrial planets form in HD 37124 and HD 74156, although in some cases 1-2 lone embryos survive for 100 Myr.  If migration occurred later, depleting the planetesimal disk, then massive terrestrial planets are unlikely to form in any of these systems. ",
        "introduction": "Most planets detected to date around main sequence stars are thought to be Jovian (gaseous) in nature.  This is known from their large masses, most of which are larger than 30 $\\mearth$ (although smaller planets have been detected -- e.g. Rivera \\etal, 2005), and from transit measurements of the size of HD209458b to be 1.27 Jupiter radii (Charbonneau \\etal 2000).  The radial velocity technique, which is sensitive to the reflex motion of a planet's parent star, is unlikely to ever be able to detect Earth-mass planets in the habitable zones of their parent stars.  The sensitivity of current surveys is 3-10 $m\\,s^{-1}$ (Butler \\etal 1996; Baranne et al. 1996; Marcy \\& Butler 1998), while the reflex velocity of the Sun due to the Earth is only about 9 $cm \\, s^{-1}$.  This signal is not likely to be detected by radial velocity surveys in the near future.  ESA's COROT and NASA's Kepler missions, to be launched in 2006 and 2007, respectively, hope to be the first to find Earth-like planets around other stars by looking for transits.  NASA's Terrestrial Planet Finder (TPF) and ESA's Darwin missions hope to spectroscopically characterize terrestrial planets around main sequence stars.  Recent results have shown that several of the known planetary systems contain regions in between the giant planets in which massless test particles remain on stable orbits for long periods of time.  Barnes \\& Raymond (2004; hereafter Paper I) mapped out these stable regions in semimajor axis $a$ and eccentricity $e$ space for HD 37124, HD 38529, HD 74156 and 55 Cnc.  These regions have been mapped in $a$ space (assuming circular orbits) for $\\upsilon$ And (Rivera \\& Lissauer 2000), GJ876 (Rivera \\& Lissauer 2001) and 55 Cnc (Rivera \\& Haghighipour 2003).  Menou \\& Tabachnik (2003) examined the possibility of Earth-sized planets residing in the habitable zones of known extrasolar planetary systems (including single planet systems), again using massless test particles.  They find that roughly one fourth of the known systems can support a planet in the habitable zone of its parent star, as defined by Kasting, Whitmire \\& Reynolds (1993).  Raymond \\& Barnes (2005; hereafter Paper II) tested the stability of Saturn-mass planets in the regions of four planetary systems in which test particles had been shown in Paper I to be stable: HD 37124, HD 38529, 55 Cnc and HD 74156.  They found that for Saturn-mass planets, the stable regions identified in Paper I shrank to a small fraction of the test particle stable region.  Barnes \\& Quinn (2004) tested the stability of seven known planetary systems, and found that several are on the edge of stability: a small change in orbital elements can lead to a catastrophic disruption of the system.  This idea led to the ``packed planetary systems'' (PPS) hypothesis first suggested by Laskar (1996), and presented in Paper I. The PPS model suggests that all planetary systems contain as many planets as they can support without becoming unstable, and implies that if a stable region exists within a planetary system, then it should contain an additional planet.  The systems studied in Paper II are not on the edge of stability, and therefore have enough ``dynamical space'' to harbor additional unseen planets.  Papers I and II dealt solely with the dynamic stability of hypothetical additional planets in planetary systems.  In this paper we examine the formation process of terrestrial planets in such a system.  Giant planets close to their parent stars (e.g., ``hot jupiters'') are thought to form farther out in the protoplanetary disk and migrate inward via torques with the gas disk (e.g., Lin \\etal 1996).  In order for a terrestrial planet to co-exist with a close-in giant planet, the terrestrial planet must either (i) form quickly and survive the inward gas giant migration, or (ii) form from material remaining after the giant planet has migrated through.  The probability of a planet surviving in the terrestrial region in scenario (i) is small.  Mandell \\& Sigurdsson (2003) showed that in some cases terrestrial planets can survive the migration of a Jupiter-like planet through the terrestrial zone.  The fraction of planets which survive such a migration is a function of the rate of migration (faster migration implies higher survival rate), and ranges between 15\\% and 40\\%.  However, only a small fraction (7-16\\%) of the surviving planets end up with orbits in the habitable zone, meaning that only 1-4\\% of terrestrial planets in the habitable zone are likely to remain on similar orbits after a migration event.  A much more likely outcome is that the planet is scattered onto a highly eccentric orbit with a large semimajor axis (Fig. 3 from Mandell \\& Sigurdsson, 2003).  Can terrestrial planets form from local material after a giant planet migrates through?  Armitage (2003) showed that in many cases the post-migration disk of planetesimals is depleted beyond repair. However, if giant planet migration occurs quickly and early enough in the disk's lifetime, then enough time remains for a second generation of planetesimals to form (Armitage, 2003).  Raymond, Quinn, \\& Lunine (2005a) argued that terrestrial accretion can therefore occur in the presence of one or more close-in giant planets via scenario (ii). Indeed, several recent results have shown that giant planets may form in less than 1 Myr via either the bottom-up, core-accretion scenario (Rice \\& Armitage 2003; Alibert, Mordasini \\& Benz, 2004; Hubickyj \\etal 2005) or the top-down, fragmentation scenario (Boss 1997; Mayer \\etal, 2002).  Migration begins immediately after (or even during) formation (Lufkin \\etal, 2004) and takes $\\sim 10^5$ years or less for planet larger than 0.1 Jupiter masses (D'Angelo, Kley \\& Henning, 2003, and references therein).  Raymond \\etal (2005a) show that potentially habitable, terrestrial planets can form in the presence of a close-in giant planet, assuming a substantial disk remains reforms after migration.  The orbit and composition of these planets are strongly affected by the position of the giant planet.  Hot/warm jupiters at larger orbital radii (up to 0.5 AU) cause the terrestrial planets to be iron-poor and in some cases drier than for closer-in giant planets, and may reduce the chances of habitable planet formation.  Here we simulate the final stages in terrestrial planet formation from disks of planetary embryos in four known systems: HD 37124 (Butler \\etal, 2003), HD 38539 (Fischer \\etal, 2003), 55 Cnc (Marcy \\etal, 2002), and HD 74156 (Naef \\etal, 2004), the same four systems we examined in Paper II.  We assume that embryos form via oligarchic growth (e.g., Kokubo \\& Ida 2000), and allow these bodies to accrete under their mutual gravity and the gravity of the known planets for 100-200 Myr.  We explore systems in which rapid, early migration has occurred, thus leaving behind a substantial planetesimal disk (Armitage 2003).  We also look at cases in which migration has occurred late in the disk lifetime, leaving behind a disk with little mass in planetesimals.  In addition, we make simple comparisons with previous simulations (e.g., Raymond \\etal 2004).  In $\\S$2 we describe our numerical method and initial conditions.  We present the results of our simulations in $\\S$3, and conclude in $\\S$4. ",
        "conclusions": "Our simulations show that certain systems of giant planets are conducive to forming terrestrial planets, others are likely to contain belts of debris or asteroids, and some are not likely to contain any rocky bodies.  Our systems of giant planets are drawn from observations, although we have made the important assumption that these systems are complete and have well-determined orbits.  We have used slightly outdated orbital parameters in order to remain consistent with previous work (specifically papers I and II).  If, for example, new values of $sin \\, i$ were determined for 55 Cnc and HD 38529, it would narrow the stable zones for additional planets and affect the region in which terrestrial planets could form.  Our ``early migration'' simulations follow the reasoning of Raymond \\etal (2005a), who argue that if giant planets form and migrate in less than roughly 1 Myr, then terrestrial planets may form via accretion in the standard way from a second generation of planetesimals (Armitage, 2003).  Our simulations therefore start with the gas giant planets already present, although their formation mechanism is unknown, be it core accretion (e.g. Rice \\& Armitage 2003) or gravitational collapse (e.g. Mayer \\etal 2002).  In either scenario, it is likely that the giant planets formed farther out in the disk and migrated inward via interactions with the gaseous disk (e.g. Lin \\etal 1996).  If giant planet migration occurred late in the evolution of protoplanetary disks, then the planetesimal disk would be severely depleted (Armitage 2003), and it is unlikely that habitable planets could form in any of the systems studied here.  The largest mass of planets that form in our ``late migration'' simulations is below the predicted lower limits of 0.2-0.3 $\\mearth$ for a 'tectonic' habitable planet (Williams \\etal 1997; Raymond, Scalo \\& Meadows 2006).  Late or slow migration could potentially deplete the disk by an order of magnitude more than we have simulated.  Since planet mass scales roughly linearly with surface density (Wetherill 1996; Kokubo \\etal 2006), these very low mass disks aren't capable of forming planets much more massive than the Moon.  Our simulations of 55 Cancri suggest that a potentially habitable planet could form {\\it in situ}.  Such a planet would have a small enough eccentricity to remain in the habitable zone throughout its orbit, substantial mass and water content.  However, as shown in Paper II, a Saturn-mass planet could exist on a stable orbit in the habitable zone of 55 Cnc.  Such a planet may preclude the existence of habitable planet, although there remains the possibility of a habitable satellite of the giant planet (Williams \\etal 1997).  The systems of terresrtial planets formed in 55 Cnc (and to a lesser extent in HD 38529) do not ressemble planets formed in preivous dynamical simulations (e.g. Agnor \\etal 1999; Chambers 2001; Raymond, Quinn \\& Lunine 2004, 2005a, 2005b, 2006).  Indeed, previous simulations tend to include systems of giant planets similar to Jupiter and Saturn, with relatively low eccentricities. The large masses and higher eccentricities of planets in 55 Cnc and HD 38529 increase the perturbations felt by embryos, causing a much higher rate of dynamical ejection than in previous simulations.  The zones in which accretion can occur correspond roughly to the stable zones from Papers I and II, and are much narrower than for systems with only an interior or exterior giant planets.  Thus, planets that form in 55 Cnc and HD 38529 are significantly smaller than in previous simulations for the same mass disk.  In addition, the distribution of planet masses tends to peak near the center of the stable regions from Papers I and II.  Strong perturbations mark the boundaries of the stable regions, so embryos that stray from the edges are quickly ejected.  The case of HD 38529 is an interesting one.  Several terrestrial bodies survive in our simulations of HD 38529, but do not accrete into large planets.  Despite strong giant planet perturbations, growing terrestrial planets do not reach high enough eccentricities to widen their feeding zones sufficiently to form Earth-sized planets.  This is likely because the stable region from Papers I and II extends only to eccentricities of 0.15 (0.3 in some areas).  Thus, planets which reach these high eccentricities are ejected rather than accreting into large terrestrial planets. We therefore speculate that a well-populated asteroid belt may exist in HD 38529, potentially including several Mars-sized planets but no Earth-sized planets.  Paper II found a wide zone ($0.3<a<0.8, e<0.15$) in HD 38529 stable for Saturn- or even Jupiter-mass planets.  Observations have ruled out the existence of such a massive planet, but not of a $\\sim$Neptune-mass planet.  However, it appears unlikely that accretion could form a large planet {\\it in situ}.  How do we reconcile this with the PPS hypothesis?  There certainly exists the possibility that a massive planet co-migrated with the inner giant planet, settling into the stable region (e.g., Thommes 2005).  Indeed, PPS predicts that such a large contiguous stable region {\\it must} contain a planet with significant mass.  This is the third paper of the Predicting Planets series.  We have conducted tests of dynamical stability and accretion in four systems of known extra-solar giant planets -- HD 37124, HD 38529, HD 74156, and 55 Cancri.  We have searched for regions in which the orbits of massless test particles are stable for millions of years (Paper I), which may correspond to stable Earth-mass planets.  We have self-consistently tested for regions which are stable for Saturn-mass test particles (Paper II).  In this paper, we have demonstrated that terrestrial planets may be able to form by accretion in one of these systems, 55 Cancri.  This series of paper has laid the foundation, along with Barnes \\& Quinn (2004), for the PPS hypothesis. These papers have clarified the conditions necessary for this hypothesis to be (dis)proven. We have also seen that some systems which initially appeared to contain stable regions, such as HD 168443, which was shown in Barnes \\& Quinn (2004) to lie far from instability, are in fact packed (Paper I; to be packed to the point that test particles cannot survive for even $10^6$ years). We have also shown that two formation avenues exist for a PPS system: simultaneous or episodic. Comparing Paper II to this work, we see that the only possibility for HD 37124 and HD 74156 to satisfy the PPS hypothesis is for an undetected giant planet to form coevally with the known planets.  However in 55 Cnc and HD 38529, additional planets could have formed and evolved together with the known planets, or possibly formed \\textit{after} a quick migratory epoch. In its current, nascent state, the PPS hypothesis cannot distinguish between these two formation events. If planets are found where we predict them to be in the 55 Cnc or HD 38529 system, then nothing can be said about their formation. Should they be found in HD 37124 or HD 74156, then only coeval formation is possible.  It remains to be seen whether the predicted planets exist.  We do not predict that stable regions will contain so much mass as to be borderline unstable. Rather, we suggest that any region in a planetary system which can support a massive planet will contain a planet.  Future observations of these well-studied planetary systems will test the credence of the PPS hypothesis."
    },
    "0404/astro-ph0404448_arXiv.txt": {
        "abstract": "Using data from the Southern Galactic Plane Survey we present a possible distant spiral arm in the fourth quadrant of the Milky Way.  The very distinct and cohesive feature can be traced for over 70\\arcdeg\\ as the most extreme positive velocity feature in the longitude-velocity diagram.  The feature is at a Galactic radius between 18 and 24 kpc and appears to be the last major structure before the end of the \\HI\\ disk.  We compare the feature with a Galactic spiral model and show that it is well reproduced by a spiral arm of pitch angle, $i \\sim 9\\arcdeg$.  The arm is quite well confined to the Galactic plane, dropping at most 1 kpc below the Galactic equator.  Over most of its length the arm is 1 - 2 kpc thick. ",
        "introduction": "\\label{sec:intro} Historically, neutral hydrogen (\\HI) in the disk of the fourth quadrant of the Milky Way has been understudied.  The last fully sampled survey of this region was made with the Parkes 18m Radiotelescope in the 1970's with an angular resolution of 48\\arcmin\\ \\citep{kerr86}.  Recently the International Galactic Plane Survey \\citep{taylor03,mcgriff01a} consortium has undertaken a large-scale project to produce a fully sampled survey of the Galactic plane with one arcminute resolution.  For the fourth quadrant, in particular, the Southern Galactic Plane Survey \\citep[SGPS;][]{mcgriff01a} constitutes more than an order of magnitude improvement in angular resolution over previous surveys.  In the Southern Milky Way little is known about the full extent of the \\HI\\ disk and the shape of the \\HI\\ density fall-off with Galactic radius. \\citet*{knapp78} made an ambitious attempt to measure the full extent of the \\HI\\ disk.  The idea behind their work was that for very large radii the radial velocity of the gas approaches the projection of the local standard of rest (LSR) velocity, $\\Theta_0$, onto the line of sight.  For an extended disk of \\HI\\ there should be a pile-up of gas at $v_{LSR} = \\Theta_0 |\\sin l|$.  Despite several very deep integrations, \\citet{knapp78} did not detect a pile-up.  Subsequent attempts, including a recent one with the Green Bank Telescope, have also failed to detect any significant pile-up of gas at extreme velocities \\citep*{lockman02}.  The outer disks of large spiral galaxies occasionally show clear \\HI\\ spiral arms extending to galactic radii in excess of 30 kpc, while the stellar arms stop at radii of $\\sim 10$ kpc.  A classic example is M83, where the \\HI\\ spiral arms extend three times as far as the stellar arms \\citep{tilanus93}.  Unfortunately, the positions of the \\HI\\ spiral arms in the outer Milky Way are not well known.  Early studies of Galactic spiral structure suggest possible arms out to distances of 20 kpc or more.  In particular, the \\citet{kerr69a} latitude-velocity diagrams show a ridge of \\HI\\ in the fourth quadrant at extreme positive velocities.  \\citet{kerr69b} maps this feature as a non-continuous spiral arm from $l=210\\arcdeg$ to $l=290\\arcdeg$. \\citet{davies72} also interpreted the low longitude ($l\\lesssim 260\\arcdeg$) end of this feature as an outer spiral arm, possibly associated with high-velocity clouds.  Despite these early studies, there has been very little interpretation of this ridge and no discussion of it beyond $l=290\\arcdeg$.  Here we present new data from the SGPS on this extreme positive velocity ridge from $l=260\\arcdeg$ to $l=330\\arcdeg$ and discuss its structure. We argue that it is a cohesive structure and interpret it as a distant \\HI\\ spiral arm at the edge of the Milky Way.  Using a simple Galactic spiral model, we show that the shape is well-fit by a spiral arm with a pitch angle of $i \\sim 9\\arcdeg$. ",
        "conclusions": "\\label{sec:conclusions} Using the SGPS we have mapped a large, coherent feature at extreme positive velocities in the fourth quadrant of the Milky Way.  The feature extends from $l=253\\arcdeg$, $v=+102$ \\kms\\ to at least $l=321\\arcdeg$, $v=+88$ \\kms, placing it at Galactic radii in excess of 18 kpc.  The feature appears to be approximately 1 to 2 kpc thick both in $z$ and along the line-of-sight.  We discuss several possible explanations for the origin of the feature, specifically that it is a pile-up in velocity space at $v = \\Theta_0 |\\sin l|$ due to an extended, smooth \\HI\\ disk, and that it results from either a non-circular LSR or an elliptical potential.  We exclude these and conclude that the feature is best described by a very distant spiral arm with a pitch angle of $i\\sim 9\\arcdeg$."
    },
    "0404/astro-ph0404481_arXiv.txt": {
        "abstract": "A main topic at this meeting is how galaxies are affected when they {\\it enter} for the first time the cluster environment from the outskirts. Most of the times we are forced to infer the environmental effects indirectly, relying on systematic variations of galaxy properties with environment, but there aren't many examples of direct observations able to unveil ongoing transformations taking place, and the corresponding mechanism producing it. We present a case in which it is possible to identify the cluster environment, and in particular the intracluster medium and the recent infall history of galaxies onto the cluster, as the cause for a recent, abrupt change in the evolutionary history of galaxies. ",
        "introduction": "This study is based on a photometric and spectroscopic survey of galaxies in the Coma cluster (Mobasher et al. 2001), which is distinctive from other surveys in three ways: for the large galaxy magnitude range covered (almost 7 mag, down to $M_B \\sim -14$), for the large area surveyed (two regions, towards the cluster center and to the South-West, $\\sim 1 \\times 1.5$ Mpc each), and for being a simply magnitude limited sample, with no morphological selection criteria adopted for spectroscopy. In this survey we have found that a significant fraction ($\\sim 10$ \\%) of the cluster dwarf galaxy population at $M_V>-18.5$ has post-starburst/post-starforming spectra (Poggianti et al. 2004, hereafter P04). This type of spectrum (``k+a'', or ``E+A'') indicates a galaxy with no current star formation activity which was forming stars at a vigorous rate in the recent past (last 1.5 Gyr). In the $B-R$ color-magnitude diagram, a group of blue and a group of red k+a's can be easily distinguished in Coma.  The average EW($\\rm H\\delta$) of the blue group is significantly stronger than that of the red group.  The blue, strong k+a's most likely correspond to ``young'' k+a's (observed soon after the termination of star formation, $<$ 300 Myr) and the red, weaker k+a's are ``old'' ones (observed at a later stage of the evolution, 0.5-1.5 Gyr).  \\begin{figure}% \\includegraphics{k+asmall.ps} \\caption{{\\it N.B. This plot is best viewed in color, see Poggianti et al. 2004.} Position of k+a galaxies with respect to X-ray substructure and X-ray temperature map. Only the central field of Coma is shown here, see P04 for a full map. Strong-lined k+a's with EW($\\rm H\\delta)>5$ \\AA $\\,$ are shown as light-coloured large dots, while weaker k+a's are plotted as darker large dots.  Tiny black dots are dwarf Coma members with velocities $> 7200 \\, \\rm km \\, s^{-1}$.  X-ray residuals from Neumann et al. (2003) are plotted as contours and clearly identify two substructures (Western and Eastern substructures), in addition to the NGC4839 peak in the South-West and the excess of emission towards the two central galaxies (NGC4874 and NGC4889).  The lowest contour and the step width between two contours are each 5 $\\sigma$.  The hardness ratio image (2-5 keV/0.5-2keV, Neumann et al. 2003) is also shown. Darker regions correspond to temperatures below 8 keV, intermediate-grey regions to $kT>8$ keV and light-coloured regions to $kT>10$ keV.  The rectangles show the limits of the two fields of our photometric and spectroscopic survey (Coma1 towards the cluster center and Coma3 in the South-West). The rectangle is about 1 by 1.5 Mpc.} \\end{figure}  A suggestive clue about the possible physical mechanism responsible for the k+a spectra comes from the recent X-ray mosaic observations of Coma obtained with {\\it XMM-Newton}. Coma has two central dominant galaxies, NGC 4874 (a cD galaxy) and NGC4889 (a very bright elliptical), and another cD galaxy, NGC4839, that dominates a substructure South-West of the center (Fig.~1).  Neumann et al. (2003) have identified and discussed X-ray substructure by fitting a smooth profile and subtracting it from the data. The residuals reveal several structures, that are shown as contours in Fig.~1: besides the well known NGC4839 South-West group, Neumann et al. find a large residual to the West of the cluster centre (``Western structure'' in Fig.~1) elongated along the North-South direction, and a filament-like structure South-East of the centre (``Eastern structure'' in Fig.~1), elongated along the East-West direction.  The temperature map shown in Fig.~1 sheds further light on the accretion history of Coma. Neumann et al. conclude that the region of high temperature observed between the Western structure and the Coma center is caused by the infall of this structure, either via compression or via shock waves.  These authors consider the two maxima in the western structure to be likely the result of the disruption of a galaxy group during its infall, instead of two galaxy groups falling at the same time.  In contrast, the South-Eastern structure is cooler than the mean cluster temperature and is associated with a low-mass galaxy group dominated by two large galaxies, NGC4911 and NGC4921.  Based on the filamentary form of this structure, the same authors conclude it is observed during the infall process while being affected by ram pressure stripping close to the cluster centre.  The coincidence of the position of the strongest k+a galaxies and the X-ray structures is striking. Four k+a's with EW($\\rm H\\delta)>5$ \\AA $\\,$ (light large dots in Fig.~1) trace the edge of the Western structure towards the Coma centre. Another three are associated with the Eastern structure, all at its western boundary. Thus, young post-starbursts are distributed close to the edge of infalling substructures. In the case of the Western substructure this edge is the infalling front, while for the Eastern substructure it is unclear whether the group is moving to the West, as suggested by the appearance of the X-ray residuals, or to the East, as suggested by the positions of NGC4911 and NGC4921 (Neumann et al. 2003). Overall, this strongly suggests that the k+a spectra, i.e. the truncation of the star formation activity in these galaxies and possibly the previous starburst, could be the result of an interaction with the hot intracluster medium (ICM).  In contrast, the location of the red k+a galaxies in Fig.~1 does not appear to be correlated with the X-ray residuals. The red k+a phase has a timescale that is comparable to the core crossing time in a cluster like Coma, and any signature of the link between the truncation of star formation and the location within a substructure is thus erased in the older k+a's, while it is still detectable in the youngest subsample of blue k+a's.  It is instructive to note that looking for a spatial segregation in the location of galaxies on the sky would not allow to establish a correlation between the star formation history of the k+a galaxies and the substructure: the link with the dynamical history of Coma appears evident only once a detailed X-ray map reveals the complicated structure in the hot intracluster gas.  The blue k+a's do show, however, a radial velocity distribution that is significantly different from that of the red k+a's and the global Coma dwarf population.  Their mean radial velocity is 8120$\\pm 709$ km $\\rm s^{-1}$, with all but one at $v>7200$ km $\\rm s^{-1}$.  In contrast, both the red k+a's and all faint galaxies with passive spectra have much lower mean velocities: 6992$\\pm 761$ and 6854$\\pm 244$ km $\\rm s^{-1}$, respectively. ",
        "conclusions": ""
    },
    "0404/astro-ph0404162_arXiv.txt": {
        "abstract": "We present optical light curves for the period 1996-2000, of two of the brightest known EGRET BL Lac objects:  PKS 2005-489 and % PKS 2155-304, the latter also one of the few known TeV sources. For both objects, quiescent epochs of low level of variability were followed by active periods, without any indication of periodicity. In PKS 2005-489, several variability events with duration of about 20 days were % observed. In PKS 2155-304 fast drops  and subsequent rises in luminosity occurred in time % scales of days. We proposed an explanation in which a region moving along the relativistic jet is eclipsed by broad line region clouds or star clusters in % the host galaxy. We compare our light curves with contemporaneous X-ray observations from All-Sky % Monitor/RXTE. Correlations between optical and X-ray activity were not found in any of the % sources at long time scales. However in PKS 2005-489 possible correlations were observed in 1997 and 1998 at % short time scales, with optical variability preceding X-rays by 30 days in 1997 and succeeding them by about 10 days in 1998. The analysis of the  SED, using the optical data presented here and BeppoSAX contemporaneous observations obtained from the literature, shows only % small shifts in the synchrotron peak as the X-ray flux density changes. ",
        "introduction": "Variability studies are a basic tool to understand the physical processes occurring in AGNs, especially those related to short time % scales and high amplitudes, as observed in BL Lacs. In this paper we made use of the five-year database of bright extragalactic sources obtained with the meridian circle of the Abrah\\~ao de Moraes Observatory (Valinhos, Brazil), between 1996 and 2000, to study light variations of two EGRET BL Lac sources: PKS 2005-489 and PKS 2155-304.   Both sources are high-frequency peaked BL Lacs (HBLs), with the synchrotron peak % located between UV and X-ray wavelengths and were also detected by EGRET in the 100 MeV % range (Lin et al. 1999). Moreover,  PKS 2155-304 is the only southern hemisphere BL Lac object detected  at TeV energies (Chadwick et al. 1999), while PKS 2005-489, also a strong candidate for detection at this energy range with the % available instrumentation (Stecker et al. 1996), has not yet been detected (Chadwick et al. % 2000).  The compact nature of PKS 2005-489 at radio frequencies was confirmed by VLBI % observations at 5 GHz  (Shen et al. 1998). Piner \\& Edwards (2004) were able to resolve PKS 2155-304 into a core and a jet % component at 15 GHz and obtained an upper limit of 4$c$ for its apparent velocity. Since even lower velocities were found for the other two TeV blazars observed in % that work, the authors concluded that stationary or subluminal velocities seem to be a characteristic of % this kind of sources, in contrast with the majority of blazars, for which parsec scale jet components are detected.  The interest in PKS 2005-489 increased after the detection of two strong X-ray % flares in 1998 by the Rossi X-Ray Timing Explorer (RXTE, Remillard, 1998, Perlman et al. 1998), however few optical data are available for this object. The only published information % about long term variations was presented by Wall et al. (1986), who found a change of % 0.5 magnitudes in the B band in two sets of observations spaced by one year. In time scales of a few days, % observations in 1994 by Heidt \\& Wagner (1996) showed 10\\% amplitude variations in the R band. Similar results were obtained by Rector \\& Perlman (2003) in observations taken some days after a third X-ray flare % in September 2000, while microvariability (variability of minutes to hours) was not detected neither at % that epoch nor in 1997 (Romero et al. 1999).  More observations at several wavelengths are available for PKS 2155-304. Using the optical and near infrared data from the literature after 1970, Fan \\& % Lin (2000) suggested the existence of 4 and 7 year periodicity in the light curves. Variability at optical wavelengths in time scales of less than 15 minutes during 1990 was reported by Heidt, Wagner \\& Wilhein-Erkes % (1997) and during 1995 by Paltani et al. (1997). Brightness variations in time scales from few days to % minutes were also observed in two large multiband monitoring campaigns, from radio to % X-rays, during the last years (e.g. Edelson et al. 1995, Pesce et al. 1997, Urry et al. 1997)  and the existence of lags between % variability at X-rays and ultraviolet wavelengths verified. However, no microvariability was found by Heidt et al. (1997) in 1994 and by % Romero et al. (1999) in 1997 and 1998, showing the existence of a duty cycle for this kind of variability.   In this paper we show that both objects present long  time scale variability, % alternating between high and low activity. In Section 2 we describe the main aspects of the instrumentation used in this work and the data reduction procedures. The % resulting differential light curves are shown in Section 3 and an analysis of the  % correlation between optical and All-Sky Monitor X-ray data is presented in Section 4. In % Section 5 we propose a scenario to explain  observed dips in the 1999 light curve of PKS % 2155-304. The implication of the data in the spectral energy distribution of the sources is % discussed in Section 6. Finally, in Section 7 we summarize our results. ",
        "conclusions": "This is the first time that a long term optical light curve for PKS 2005-489 is % presented. Moreover, the addition of more optical data to the well studied (and yet poorly understood) % BL Lac PKS 2155-304 allowed us to confirm that the two studied objects show high degree of variability at different time scales.    In PKS 2005-489 we could identify several variability events with duration of % about 20 days, although  light curves with better time sampling are necessary to verify if  in fact they are arising from similar % phenomena. PKS 2155-304 showed a significant activity in short time scales during 1999, with % amplitude variations of almost 1 magnitude in  a few days. Of special interest is the presence of at least two sharp dips in the light curve, with duration of about two days each. % We interpreted them as the result of eclipses, caused by BLR clouds or star clusters in the host galaxy.  To obtain the right time scales we had to assume that the the % eclipsed region is moving along the  jet with relativistic velocities. The detection of a superluminal % component formed at that epoch gives support to our assumption (Piner \\& Edwards 2004). Although our observations cover some months before and after the X and gamma ray activity of 1997, which resulted in % the detection of emission at  TeV energies (Chadwick et al. 1999), no unusual behavior in the % optical light curve was detected at those epochs.   We searched for correlation between the optical light curves and X-rays in the 1.5 - 12 keV range from the All-Sky Monitor project. We did not find significant correlation for any of the  sources in time scales of % years. However, in time scales of days, we found in PKS 2005-489 possible correlations % in the campaigns of 1997 and 1998. In the first year the optical variability preceded the X-rays by approximately 30 days and in 1998 X-rays preceded the optical % variability by about 10 days. Although we believe that the optical emission is always due to the synchrotron process, % the  X-ray activity can be explained  also as synchrotron radiation by several models (Geoganopoulos % \\& Marscher 1998, Blandford \\& K\\\" onigl 1979, Marscher 1987) only when it precedes the optical variability. Otherwise the same models can be applied, but assuming the the X-ray emission is % due to the inverse Compton process involving synchrotron radio or infrared photons. The 1999 optical and X-ray light curves of PKS 2155-304 seems to be correlated, but it was not possible to determine the presence of time % delays because of the width of the DCF function.   We also constructed the SED with our optical data and contemporaneous BeppoSAX % X-ray observations. First we fitted a simple parabola in the ${\\rm log} \\,F_{\\nu} \\times {\\rm log} \\, \\nu$ plane and then constructed the  SED, instead of using multi-parametric models, as usually done (eg. Padovani et al. % 2001). Although the observation are only contemporaneous, it seems that the  relative amplitude of the optical variability, if any, is much smaller than that of  % X-rays. The obtained SEDs confirm that the synchrotron peaks move to higher frequencies % during the occurrence of X-ray flares, although the shifts were much smaller than what was % observed for other blazars (Pian et al. 1998, Giommi et al. 1999)."
    },
    "0404/astro-ph0404354_arXiv.txt": {
        "abstract": "We present an {\\it XMM-Newton} spectrum of diffuse X-ray emission from within the solar system.  The spectrum is dominated by probable \\cvi\\ lines at 0.37~keV and 0.46~keV, an \\ovii\\ line at 0.56~keV, \\oviii\\ lines at 0.65~keV and $\\sim0.8$~keV, \\neix\\ lines at $\\sim0.92$~keV, and \\mgxi\\ lines at $\\sim1.35$~keV.  This spectrum is consistent with that expected from charge exchange emission between the highly ionized solar wind and either interstellar neutrals in the heliosphere or material from Earth's exosphere.  The emission is clearly seen as a low-energy ($E<1.5$~keV) spectral enhancement in one of a series of four observations of the {\\it Hubble} Deep Field North.  The X-ray enhancement is concurrent with an enhancement in the solar wind measured by {\\it ACE}, {\\it Wind}, and {\\it SoHO} spacecraft.  The solar wind enhancement reaches a flux level an order of magnitude more intense than typical fluxes at 1~AU, and has a significantly enhanced O$^{+7}$/O$^{+6}$ ratio. Besides being of interest in its own right for studies of the solar system, this emission can have significant consequences for observations of cosmological objects.  It can provide emission lines at zero redshift which are of particular interest in studies of diffuse thermal emission (e.g., \\ovii\\ and \\oviii), and which can therefore act as contamination in the spectra of objects which cover the entire detector field of view.  We propose the use of solar wind monitoring data as a diagnostic to screen for such possibilities. ",
        "introduction": "Diffuse X-ray emission from the solar system was clearly observed during the {\\it ROSAT} All-Sky Survey (RASS) as a background component with temporal variations on scales from a large fraction of a day to many days. These variations were dubbed Long-Term Enhancements (LTEs, \\citet{sea95}), and provided a significant background particularly at \\oqkev.  Although the origin of these LTEs at the time was unknown, the LTEs were associated with solar wind parameters \\citep{frey94}.  With the observation of X-rays from comets \\citep{lis96}, emission from solar wind charge exchange (SWCX) between solar wind ions and neutral material within the heliosphere was demonstrated \\citep{cra97}.  SWCX emission subsequently was suggested as the origin for LTEs \\citep{cox98,cra00}, and has recently been suggested as being responsible in quiescence for a significant fraction of diffuse X-ray emission at \\oqkev\\ \\citep{lal04,rc03b} previously attributed to the Local Hot Bubble \\citep{sea98}.  More recently, evidence for geocoronal emission (SWCX with exospheric material) has been detected in {\\it Chandra} observations of the dark moon \\citep{wea04}. This emission clearly originates in the near-Earth environment, which is consistent with the likely production of X-rays from the terrestrial magnetosheath \\citep{rc03a}.  The statistics of the dark moon data are somewhat limited but they clearly show excess \\ovii\\ and \\oviii\\ emission.  Analysis of cometary X-ray spectra suggest that solar wind charge state composition and speed affect the X-ray emission \\citep{lea01,kea03,sc00}.  Thus through a combination of techniques, including X-ray and neutral atom imaging among others, it may be possible to continuously monitor the solar wind from well inside Earth's magnetosphere.  However, the nature of solar system X-ray emission is strongly dependent on the point of view of the observer.  While the observation of SWCX X-ray emission samples aspects of the solar system of interest to astronomers, the emission can also provide a contaminating component which strongly impacts observations of extended cosmic X-ray sources.  The SWCX emission spectrum is dominated by highly ionized carbon, oxygen, neon, possibly iron, and magnesium lines, which are also of great astrophysical interest.  (The oxygen lines are particularly important as they are commonly used for temperature and density diagnostics of thermal emission from diffuse plasmas.) For objects at zero redshift (i.e., emission from the Milky Way and nearby galaxies) where the emission is expected to fill the entire field of view of the detector, SWCX emission is indistinguishable from that of the object, except to the extent that temporal variation can be detected.  While there is likely to be SWCX emission at some quiescent level \\citep{rc03b,lal04}, the strongest emission will be associated with flux enhancements of the solar wind.  These enhancements are long enough in duration that they can ``contaminate'' an entire observation, but variable enough that if there are multiple observations of a single target they can be identified, at least to some level.  This was the basis of the LTE ``cleaning'' of the RASS \\citep{sea95,sea97}, and of {\\it ROSAT} pointed observations which were typically distributed over the period of a few days if not weeks \\citep{sea94}.  In this paper we report the detection of SWCX emission during an {\\it XMM-Newton} observation of the {\\it Hubble} Deep Field North.   We correlate that detection with a concurrent enhancement in the solar wind density observed by Advanced Composition Explorer ({\\it ACE}), Solar and Heliospheric Observatory ({\\it SoHO}), and {\\it Wind}.  In \\S~\\ref{sec:data} we present the data, data reduction, and analysis, in \\S~\\ref{sec:discussion} we discuss the results, and in \\S~\\ref{sec:conclusions} we detail our conclusions both in regards to the emission itself and the implications for X-ray observations of more distant objects. ",
        "conclusions": "\\label{sec:conclusions}  The data from the {\\it XMM-Newton} observation of the {\\it Hubble} Deep Field North show clear evidence for a time-variable component of the X-ray background that is most reasonably attributed to charge exchange emission between the highly ionized solar wind and exospheric or interplanetary neutrals.  The SWCX emission is dominated by the lines from \\cvi, \\ovii, \\oviii, \\neix, and \\mgxi, and possibly lines from highly ionized iron (e.g., \\fexvii).  The emission is concurrent with the passage of a strong enhancement in the solar wind observed by the {\\it ACE} satellite, and which is also strongly enhanced in its O$^{+7}$/O$^{+6}$ ratio, and possibly other highly ionized species (although preliminary results for the O$^{+8}$/O$^{+7}$ ratio indicate otherwise). However, the light curve of the X-ray enhancement is relatively constant while the flux of the solar wind enhancement varies considerably.  The observation of SWCX emission allows the monitoring of interactions between the solar wind and solar system and/or interstellar neutrals without the need for {\\it in situ} measurements, albeit with the uncertainty of just where along the line of sight the emission arises.  However, in certain circumstances the distance ambiguity can be resolved allowing the detailed study of the related phenomena.  Two such situations are mentioned above: the case of SWCX emission from comets and exospheric material in the {\\it Chandra} dark moon observation.  The ability to remotely sense the solar wind and its interactions can also be used when the emission is expected to have a predictable variation, such as an X-ray scan across the magnetopause sub-solar point \\citep{rc03a}.  Also, for example, {\\it ROSAT} data from the All-Sky Survey suggest that the downstream helium focusing cone can be imaged in X-rays (e.g., \\citet{cea03}).  The charge exchange line emission can provide a significant contaminating background to observations of more distant objects (by any X-ray observatory with sensitivity at energies less than 1.5~keV) which use those lines for diagnostics.  However, the correlation of the X-ray enhancement with the solar wind density enhancement suggests a diagnostic.  For those observations at risk because of such contamination, e.g., observations of sources which cover the entire field of view and have thermal spectra, the data from solar wind monitoring observatories such as {\\it ACE}, {\\it Wind}, and {\\it SoHO} should be used for screening.  Should the SWCX emission be primarily exospheric in origin (e.g., from the magnetosheath), then the specific geometry of the observation should be examined as well.  An observation where the observatory is outside of the magnetosheath and looking away from Earth would clearly have significantly less exospheric SWCX emission than the case of the emission reported in this paper where the observation line of sight may be passing tangentially through the magnetopause near the sub-solar point.  While SWCX emission is likely responsible for the LTEs observed during the {\\it ROSAT} All-Sky Survey, the SWCX episode presented here was significantly brighter at \\tqkev\\ than what was typically observed (a factor of two brighter than the most intense \\tqkev\\ RASS LTE). However, one important issue to note is that none of the X-ray emission lines discussed here (or discussed in \\citet{wea04}) contribute significantly to the \\oqkev\\ band, where most of the LTE emission was observed.  Also of note is that the geometry of this observation relative to Earth's geomagnetic environment is considerably different from that of the RASS. In this case we may be looking tangentially through the magnetosheath near the sub-solar point while {\\it ROSAT} with its low circular orbit effectively looked radially outward through the flanks of the magnetosheath."
    },
    "0404/astro-ph0404389_arXiv.txt": {
        "abstract": "We investigate the expected gravitational wave emission from coalescing supermassive black hole (SMBH) binaries resulting from mergers of their host galaxies. When galaxies merge, the SMBHs in the host galaxies sink to the center of the new merged galaxy and form a binary system. We employ a semi-analytic model of galaxy and quasar formation  based on the hierarchical clustering scenario to estimate the amplitude of the expected stochastic gravitational wave background owing to inspiraling SMBH binaries and bursts owing to the SMBH binary coalescence events. We find that the characteristic strain amplitude of the background radiation is $h_c(f) \\sim 10^{-16} (f/1 \\mu {\\rm Hz})^{-2/3}$ for $f \\lesssim 1 \\mu {\\rm Hz}$ just below the detection limit from measurements of the pulsar timing provided that SMBHs coalesce simultaneously when host galaxies merge. The main contribution to the total strain amplitude of the background radiation comes from SMBH coalescence events at $0<z<1$. We also find that a future space-based gravitational wave interferometer such as the planned \\textit{Laser Interferometer Space Antenna} ({\\sl LISA}) might detect intense gravitational wave bursts associated with coalescence of SMBH binaries with total mass $M_{\\rm tot} < 10^7 M_{\\odot}$ at $z \\gtrsim 2$ at a rate $ \\sim 1.0 \\ {\\rm yr}^{-1}$. Our model predicts that burst signals with a larger amplitude $h_{\\rm burst} \\sim 10^{-15}$ correspond to coalescence events of massive SMBH binary with total mass $M_{\\rm tot} \\sim 10^8 M_{\\odot}$ at low redshift  $ z \\lesssim 1$ at a rate $ \\sim 0.1 \\ {\\rm yr}^{-1}$ whereas those with a smaller amplitude $h_{\\rm burst} \\sim 10^{-17}$ correspond to coalescence events of less massive SMBH binary with total mass $M_{\\rm tot} \\sim 10^6 M_{\\odot}$ at high redshift $ z \\gtrsim 3$. ",
        "introduction": "\\label{sec:intro} In recent years, there has been  increasing observational evidence that many nearby galaxies have central supermassive black holes (SMBHs) in the mass range of $10^6-10^9 M_{\\odot}$, and that their physical properties correlate with those of spheroids\\footnote{Throughout this paper, we refer to bulge or elliptical galaxy as {\\it spheroid}.} of their host galaxies. First, the estimated mass of a SMBH in a galactic center, $M_{\\rm BH}$, is roughly proportional to the mass of the spheroid, $M_{\\rm spheroid}$. The ratio of $M_{\\rm BH}/M_{\\rm spheroid}$ is $0.001-0.006$ in each galaxy (e.g. \\citealp{Kormendy95}; \\citealp{Magorrian98}; \\citealp{Merritt01b}). Second, the mass of a SMBH correlates with the velocity dispersion of stars in the spheroid, $\\sigma_{\\rm spheroid}$, as $M_{\\rm BH} \\propto \\sigma_{\\rm spheroid}^{n},n=3.7-5.3$ (e.g. \\citealp{Ferrarese00}; \\citealp{Gebhardt00}; \\citealp{Merritt01a}; \\citealp{Tremaine02}). These relations suggest that the formation of SMBHs physically links to the formation of spheroids that harbor the SMBHs.  In order to study the formation and evolution of SMBHs, it is necessary to construct a model that includes galaxy formation and merger processes. It has become widely accepted that quasars are fueled by accretion of gas onto SMBHs in the nuclei of host galaxies. Hydrodynamical simulations have shown that a merger of galaxies drives gas fall rapidly onto the center of the merged system and fuels a nuclear starburst leading to spheroid formation \\citep[e.g.][]{Mihos94,Mihos96,Barnes96}. Some observations of quasar hosts show that many quasars reside in spheroids of interacting systems or elliptical galaxies \\citep[e.g.][]{Bahcall97,McLure99,Dunlop03,Kauffmann03}. Thus the major merger of galaxies is a possible mechanism for quasar activity, SMBH evolution and spheroid formation.  In the standard hierarchical structure formation scenario in a cold dark matter (CDM) universe, dark-matter halos ({\\it dark halos}) cluster gravitationally and merge together. In each of merged dark halos, a galaxy is formed as a result of radiative gas cooling, star formation, and supernova feedback. Several galaxies in a common dark halo sometimes merge together and a more massive galaxy is assembled. When galaxies merge, SMBHs in the centers of the galaxies sink toward the center of the new merged galaxy and form a SMBH binary subsequently. If the binary loses enough energy and angular momentum, it will evolve to the gravitational wave emitting regime and begin inspiraling, eventually coalesces with a gravitational wave burst.  An ensemble of gravitational waves from a number of inspiraling SMBH binaries at different redshift can be observed as a stochastic background at frequencies $\\sim 1{\\rm n} - 1 \\mu {\\rm Hz}$, which can be detected by pulsar timing measurements \\citep{Detweiler79}. Gravitational wave bursts from coalescence events can be detected by a Doppler tracking test of interplanetary spacecraft \\citep{Thorne76}. Future space interferometers such as the Laser Interferometer Space Antenna \\textit{LISA} might detect quasi-monochromatic wave sources associated with inspiraling SMBH binaries and gravitational wave bursts associated with SMBH binary coalescence.   In order to estimate the feasibility of detecting such gravitational waves, we need to know the SMBH coalescing rate that depends on the complex physical processes before coalescence. A possible scenario of forming SMBHs consists of four stages: First, a seed massive black hole in the center of a galaxy grows to an intermediate mass black hole (IMBH) via run-away collapse in a dense star cluster \\citep{Zwart04}. Second, a galaxy that contains an IMBH in their center merges with another galaxy with an IMBH. Third, the two IMBHs sink to the center of gravitational potential well owing to dynamical friction until a ``hardening'' regime \\citep{Volonteri03}. Finally, physical processes such as three-body processes, gas dynamics, the ``Kozai'' mechanism, etc. drive the binary to the gravitational wave emission regime \\citep[e.g.][]{Begelman80,Makino97s,Gould00,Yu02,Armitage02,Blaes02,Volonteri03,Milosavljevic03}. However , the efficiency to bring two SMBHs together is still unknown.  To date, a number of attempts have been made to calculate the SMBH coalescence rate taking the physical processes in the ``second'' and ``third'' stages into account. Some authors use phenomenological models of galaxy mergers based on number counts of quasars and spheroids \\citep[e.g.][]{Thorne76, Fukushige92, Haehnelt94, Rajagopal95, Jaffe03}. Others use merger rates of dark halos (not galaxies) based on formalism of \\cite{PS74} and \\cite{Lacey93} \\citep[e.g.][]{Haehnelt94, Menou01}. \\cite{Volonteri03} calculated the SMBH coalescence rate taking conditions under which ``sub-halos'' with SMBHs in a dark halo sink to the dark halo center into account \\citep[see also][]{Wyithe03,Sesana04}. However, none of these models include baryonic gas evolution and galaxy formation processes. Because SMBH formation process is relevant to spheroids of host galaxies rather than to dark halos, we need to evaluate how the baryonic gas processes such as star formation and radiative cooling affect the SMBH coalescence rate.  In this paper, we estimate the SMBH coalescence rate using a new semi-analytic (SA) model of \\cite{Enoki03} (an extended model of \\citet{Nagashima01b}) in which the SMBH formation is incorporated into the galaxy formation. Then, we calculate the spectrum of gravitational wave background from inspiraling SMBH binaries, based on the formulation given by \\cite{Jaffe03} and we compare our result with that from a pulsar timing measurement. We also estimate the event rate of gravitational wave bursts from SMBH coalescence events that might be detected by future planned space laser interferometers, based on an argument in \\cite{Thorne76}.  In SA models, merging histories of dark  halos are realized  using a Monte-Carlo algorithm  and evolution of baryonic components within dark halos is calculated using simple analytic models for gas cooling, star formation, supernova feedback, galaxy merging and other processes. SA models have successfully reproduced a variety of observed features of galaxies, such as their luminosity functions, color distributions, and so on \\citep[e.g.][] {KWG93,Cole94,Cole00,SP99,Nagashima01b,Nagashima02}. \\cite{KH00} introduced a unified model that includes the formation of both galaxies and SMBHs within the framework of their SA model \\citep[see also][]{Cattaneo01,Menci03,Islam03,Granato04}. Our SA model reproduces not only these observational features but also the observed present black hole mass function and the quasar luminosity functions at different redshifts \\citep{Enoki03}.  The paper is organized as follows: in section \\ref{sec:merger} we briefly review our SA model for galaxy formation and SMBH growth; in section \\ref{sec:GWR} we calculate the spectrum of gravitational wave background and the event rate of gravitational wave bursts; in section \\ref{sec:summary} we provide summary and conclusions. ",
        "conclusions": "\\label{sec:summary}  In this paper, we have estimated the coalescence rate of supermassive black hole (SMBH) binaries in the centers of galaxies using a new semi-analytic model of galaxy and quasar formation (SA model) given by \\cite{Enoki03} based on the hierarchical structure formation scenario. Then, we calculated the spectrum of the gravitational wave background from inspiraling SMBH binaries based on the formulation given by \\cite{Jaffe03} and estimated the expected amplitudes and event rates of intense bursts of gravitational waves from coalescing SMBH binaries.  Our SA model includes dynamical friction and random collision as galaxy merging mechanisms, and assumes that a SMBH is fueled by accretion of cold gas during a  major merger of galaxies leading to a spheroid formation, and that SMBHs coalesce simultaneously when host galaxies merge. Many previous other studies have paid attention to only SMBH growth and did not take galaxy formation processes into account. For investigating the relations between SMBH growth and galaxy formation processes, SA methods of galaxy and SMBH formation are suitable tools \\citep[]{KH00,Cattaneo01,Enoki03, Menci03,Granato04}. Our SA model can reproduce not only observational properties of galaxies, but also the present SMBH mass function and the quasar luminosity functions at different redshifts \\citep{Nagashima01b,Enoki03}.  We have found that the gravitational wave background radiation spectrum for $f \\lesssim 1 \\mu {\\rm Hz}$ has a characteristic strain $h_c(f) \\sim 10^{-16} (f/1 \\mu {\\rm Hz})^{-2/3}$ just below the detection limit from the current measurements of the pulsar timing. The slope of the spectrum for $f \\gtrsim 1 \\mu {\\rm Hz}$ gets steep owing to the upper limit in frequency set by the radius of the innermost stable circular orbit. The stochastic background radiation mainly comes from inspiraling SMBH binaries at $0<z<1$. Therefore, the background radiation can probe inspiraling SMBH binaries at low redshift.   We have also found that {\\sl LISA} might detect intense bursts of gravitational waves owing to the SMBH coalescence events at a rate $0.1 \\sim 1.0 {\\rm ~yr}^{-1}$ and that the main contribution to the event rate comes from SMBH binary coalescence at high redshift $z \\gtrsim 2$. Our model predicts that burst signals with a large amplitude correspond to coalescence of large mass SMBH  binaries at low redshift while those with a small amplitude correspond to coalescence of small mass SMBH binaries at high redshift. This prediction can be tested by future measurements of the amplitude and the phase evolution in gravitational waves from inspiraling SMBH binaries \\citep{Hughes02}. Comparing these predictions with observations in future, we can put a stringent constraint on SMBH formation and evolution models."
    },
    "0404/astro-ph0404340_arXiv.txt": {
        "abstract": "We study the formation of molecular hydrogen on dust grain surfaces and apply our results to the high redshift universe. We find that a range of physical parameters, in particular dust temperature and gas temperature, but not so much dust surface composition, influence the formation rate of H$_2$.  The H$_2$ formation rate is found to be suppressed above gas kinetic temperatures of a few hundred K and for dust temperatures above 200-300 K and below 10 K. We highlight the differences between our treatment of the H$_2$ formation process and other descriptions in the literature.  We also study the relative importance of H$_2$ formation on dust grains with respect to molecular hydrogen formation in the gas phase, through the H$^-$ route. The ratio of formation rates of these two routes depends to a large part on the dust abundance, on the electron abundance, and also on the relative strength of the FUV (extra-)galactic radiation field.  We find that for a cosmological evolution of the star formation rate and dust density consistent with the Madau plot, a positive feedback effect on the abundance of H$_2$ due to the presence of dust grains can occur at redshifts $z\\ge 3$. This effect occurs for a dust-to-gas mass ratio as small as 10$^{-3}$ of the galactic value. ",
        "introduction": "\\hm\\ formation in the universe is a process which has been studied extensively in the past decades, but which is still not well understood. In the interstellar medium, \\hm\\ formation occurs on grain surfaces, permitting three-body reactions that are much more efficient than gas phase reactions (Gould \\& Salpeter 1963). Many studies have been performed over the years and focussed on explaining the mechanisms involved in the \\hm\\ formation process on grain surfaces. It is clear that this process is governed by the mobility of the H atoms on the grain, as well as the energies binding the H atoms to the surface. The mobility of an adatom (the atom that is bound to the surface) is the combination of two processes: thermal diffusion and quantum tunneling. Therefore, an adatom can move on the surface of the grain, from site to site, through these two processes according to the temperature of the grain and the characteristics of the surface (Barlow \\& Silk 1976; Leitch-Devlin \\& Williams 1984; Tielens \\& Alamandola 1987). On astrophysically relevant surfaces (olivine and amorphous carbon), an adatom can bind in two energetically different sites: physisorbed sites and chemisorbed sites. Physisorbed atoms are weakly bound to the surface and are mobile for low grain temperatures (around 10 K), whereas chemisorbed atoms are bound stronger to the surface and become mobile for a grain temperature of a few hundred K (Barlow \\& Silk 1976; Aronowitch \\& Chang 1980; Klose 1992; Fromherz et al.\\ 1993; Cazaux \\& Tielens 2002, 2004). Therefore, \\hm\\ formation strongly depends on the metallicity of the system (i.e., the dust abundance), the considered grains (olivine or amorphous carbon), the temperature of the grains and the atomic hydrogen abundance.  Studies of \\hm\\ formation in low metallicity systems -- e.g., the high redshift universe -- have already been performed (c.f.\\ Tegmark et al.\\ 1997; Glover 2003).  \\hm\\ at high redshift is of prime importance since it is considered to be the only coolant below $10^4$ K. Therefore, as discussed by many authors, \\hm\\ plays a crucial role in the formation of the first stars (Haiman, Rees, \\& Loeb 1996; Haiman, Thoul, \\& Loeb 1996; Tegmark et al.\\ 1997), and thus determines the end of the dark ages (Haiman, Abel, \\& Rees 2000). Other authors discussed the possibility of HD as a more efficient coolant because this molecule possesses a nonzero dipole moment, even though its abundance is much lower (Galli \\& Palla 1998; Nakamura \\& Umemura 2002). Norman \\& Spaans (1997) estimate a redshift of roughly unity, where \\hm\\ formation in the gas phase and \\hm\\ formation on grain surfaces are equally important. At high redshift, \\hm\\ formation proceeds through the H$^-$ route.  Below a redshift of unity, grain surface reactions become more important relative to the (inefficient) H$^-$ route. Note that at high densities H$_2$ is formed through three-body association.  These issues are of more than academic interest since the presence of H$_2$, either through the H$^-$ or dust grain routes, is instrumental in the ability of low metallicity gas to cool and contract to form stars. Because no dust grains are available prior to the formation of pop III objects, the H$^-$ route is crucial in starting up the star formation process, but the subsequent expulsion of metals in supernova explosions and ambient dust formation can strongly enhance the formation of molecular hydrogen (Hirashita, Hunt, \\& Ferrara 2002). \\\\The construction of comprehensive models for the effects of dust grains on early galaxy evolution and their observational consequences is underway (c.f.\\ Hirashita, Hunt, \\& Ferrara 2002; Spaans \\& Norman 1997; Kauffmann \\& Charlot 1998; Somerville, Primack \\& Faber 2001; Glover 2003), motivated by upcoming observatories like ALMA and JWST. In general, such studies adopt the H$_2$ formation rate of Hollenbach \\& McKee (1979). Although a good order of magnitude estimate, recent work (Cazaux \\& Tielens 2002, 2004) has allowed the formation of H$_2$ (and HD) on dust grains to be computed more accurately.  Because the H$_2$ chemistry plays such an important role in the high redshift universe it is crucial to put it on the firmest of physical bases.  The aim of this paper is then to provide detailed computations of the H$_2$ formation efficiency and formation rate as functions of ambient physical conditions and to extend earlier work on the effects of dust grains on the physical properties of primordial gas.  In this paper, we use the model developed by Cazaux \\& Tielens (2004) and extend it to accommodate the conditions in the high redshift universe.  This model, which describes the formation of \\hm\\ on grain surfaces, takes into account the presence of both physisorbed and chemisorbed sites on the grain surface, and allows quantum mechanical diffusion as well as thermal diffusion for adatoms to go from site to site. In Section two, we calculate first the \\hm\\ formation efficiency for typical H fluxes, different surface characteristics, and different dust grain abundances, and second the \\hm\\ formation efficiency through the gas phase route (i.e., the H$^-$ route). In Section three, we discuss which are the physical conditions (i.e., dust and gas temperature, dust abundance, electron fraction) at which the H$^-$ route and grain surface route make equal contributions to the \\hm\\ formation rate. Then, we adopt a cosmological model in order to assess at which redshift this equality occurs, consistent with the star formation rate history described by Madau ``the Madau plot'' (e.g., the universal star formation rate as a function of redshift; Lilly et al.\\ 1996; Madau et al.\\ 1996). Therefore, we estimate when, in the history of the universe, the dust route became the dominant process toward \\hm\\ formation. In Section four, we discuss the effects of this gas to dust route transition on the star formation rate and the ambient radiation field. ",
        "conclusions": "We have studied the formation of H$_2$ on dust grain surfaces at high redshift and we have related our results to the contribution made by gas phase reactions, i.e., through the H$^-$ route. We have found that the substrate (olivine or amorphous carbon) has a modest impact on the resulting H$_2$ formation rate. The formation efficiency depends strongly on the dust temperature below 10 K and above few hundred kelvins. The role of the gas temperature is more limited, but suppresses the H$_2$ formation rate above several hundred K due to a reduced sticking coefficient.  We wish to stress that the microphysical results on the formation efficiency and the H$_2$ formation rate on dust grains are robust and independent of any galactic context.  We adopted a cosmological model to determine at which cosmological parameters the dust and gas phase contributions to the H$_2$ formation rate are equal, and thus when, in the universe's history, the dust grain route becomes the dominant H$_2$ formation process. That is, when the presence of dust, a result of star formation, leads to an enhancement of the H$_2$ formation rate, which in turn can boost the H$_2$ abundance (depending on the internal FUV radiation field that can dissociate H$_2$) and hence the ambient cooling rate by driving a rapid ion-molecule chemistry that leads to the formation of species like CO (an excellent low-temperature coolant). Such a cycle constitutes a positive feedback loop (Hirashita, Hunt, \\& Ferrara 2002) and can enhance the star formation rate inside a galaxy.  Our results show that, within the uncertainties of our cosmological model for the evolution of disk galaxies, the conditions for this positive feedback can occur at a redshift above 3 which corresponds to a dust-to-gas mass ratio less than 10$^{-3}$ times the Galactic value. This redshift range is large and depends strongly on the dust grain and gas temperatures (microphysics) that we adopted as well as on the star formation rate (macrophysics) that we used.  Indeed, for high dust/gas grain temperatures, the atoms desorb/bounce from/off the grain surface which decreases the \\hm\\ formation rate. On the other hand, a high gas temperature favors gas phase formation of H$_2$. This clearly shows the importance of both dust and gas temperature to determine which of these two routes dominates. At very high redshift, $z>15-20$, the temperature of the dust is coupled to the temperature of the CMB, which increases like $1+z$.  The dust grains are then too warm to allow an important dust route contribution. So even if dust would be present at these redshifts, it would not boost the H$_2$ formation rate until the universe had cooled down considerably through cosmological expansion.   In fact, once the presence of dust boosts the H$_2$ formation rate, and hence the star formation rate through the enhanced cooling rate, for redshift higher than $z\\sim 3$, the production of stellar photons will raise the mean dust (and gas) temperature. This constitutes a minor effect when the shielding dust columns are large and the McKee criterium is satisfied (McKee 1989), but might be quite important in the first, metal-poor stages of star formation.  In any case, the magnitude of the positive feedback that the presence of dust has on the H$_2$ formation rate requires a careful treatment of the impact that (enhanced) star formation activity has on the dust (and gas) thermal equilibrium. We postpone these matters to a future paper.   Finally, recent observations of distant quasars (Bertoldi et al. 2003) at redshifts around z$\\sim$6 showed that these objects possess a metallicity close to solar. These quasars represent large over-densities in the Universe and our study concentrates on the evolution of a typical sub-L* galaxy, that we assume to be more representative of the average galaxy population. In any case, these distant quasars, rich in metals, possess both a high dust grain abundance and the physical conditions to allow efficient \\hm\\ formation on dust surfaces.   We would like to thank the anonymous referee for his careful reading of the manuscript and his comments that helped to improve this work."
    },
    "0404/astro-ph0404495_arXiv.txt": {
        "abstract": "We present {\\it Hubble Space Telescope} Wide Field Planetary Camera 2 images of a prominent externally-ionized molecular globule in the Carina Nebula (NGC~3372), supplemented with ground-based infrared images and visual-wavelength spectra.  This molecular globule has a shape resembling a human hand, with an extended finger that points toward its likely source of ionizing radiation.  Following an analysis of the spatially-resolved ionization structure and spectrum of the photoevaporative flow from the Finger, we conclude that the dominant ionizing source is either the WNL star WR25 (HD~93162), the adjacent O4~If-type star Tr16-244, or perhaps both.  We estimate a mass-loss rate of $\\sim$2$\\times$10$^{-5}$ M$_{\\odot}$ yr$^{-1}$ from the main evaporating surface of the globule, suggesting a remaining lifetime of 10$^{5.3}$ to 10$^6$ years. We find a total mass for the entire globule of more than 6 M$_{\\odot}$, in agreement with previous estimates. The hydrogen column density through the globule derived from extinction measurements is a few times 10$^{22}$ cm$^{-2}$, so the photodissociation region behind the ionization front should be limited to a thin layer compared to the size of the globule, in agreement with the morphology seen in H$_2$ images.  Although a few reddened stars are seen within the boundary of the globule in near-infrared continuum images, these may be background stars.  We do not detect a reddened star at the apex of the finger, for example, down to a limiting magnitude of $m_K\\simeq$17.  However, considering the physical properties of the globule and the advancing ionization front, it appears that future star formation is likely in the Finger globule, induced by radiation-driven implosion. ",
        "introduction": "Compact bright-rimmed molecular globules are common in evolved massive star forming regions, and are among the last vestiges of the molecular cloud that gave birth to the OB stars powering the H~{\\sc ii} region (e.g., Bok \\& Reilly 1947).  These globules are externally illuminated and are photoevaporated and photoionized by UV radiation.  They typically range in size from 0.1 to 1 pc, and appear striking in [S~{\\sc ii}] $\\lambda\\lambda$6717,6731 and H$\\alpha$ emission arising in the limb-brightened surfaces of their thin ionization fronts. Famous examples of this class of objects are Thackeray's globules in IC~2944 (Thackeray 1950; Reipurth et al.\\ 2003), as well as similar features in the Rosette Nebula (Herbig 1974), the Gum Nebula (Hawarden \\& Brand 1976; Reipurth 1983), and the Carina Nebula (Walborn 1975; Cox \\& Bronfman 1995).  \\begin{figure*}\\begin{center} \\epsfig{file=nsmith.fig1.eps,width=5.1in}\\end{center} \\caption{Large-scale image of the environment around the Finger, including the Keyhole and the background Carina nebula, $\\eta$ Car, WR25, and other stars.  The [S~{\\sc ii}] image shown here was obtained with the MOSAIC camera on the CTIO 4m telescope.  The inset shows an enlargement of the Finger in the H$\\alpha$ filter, along with the orientation of the slit apertures used for spectra in Figure 5.} \\end{figure*}  Material evaporated from these clumps partly fills the interior of an H~{\\sc ii} region and may help stall the advance of the main ionization front by absorbing incident Lyman continuum radiation and converting it to a recombination front, but perhaps the most pertinent role played by these globules is that they are potential sites of continued star formation.  Some globules show clear evidence for embedded star formation, while others do not (Reipurth 1983; Reipurth et al.\\ 2003).  If they are sites of star formation, one possibility is that they are simply dense cores that spontaneously collapsed to form stars and were then uncovered by the advancing ionization front. Another more intriguing idea is that pressure from the ionization-shock front at the surface propagates through a globule and helps to overcome the magnetic, turbulent, and thermal pressure that supports it against collapse, thereby inducing a star to form in the globule.  This process of radiation-driven implosion is often referred to as {\\it triggered star formation}.  However, direct and unambiguous observational evidence for actively triggered star formation remains elusive; the pillars in M16 may be an example of this phenomenon (White et al.\\ 1999; Hester et al.\\ 1996; Williams et al.\\ 2001; McCaughrean \\& Andersen 2002).  Numerous investigators have examined theoretical aspects of these evaporating globules, including details of the photoevaporative flows, their effect on the surrounding H~{\\sc ii} region, the globule's shaping, destruction, and acceleration (through the ``rocket effect''), and the possibility of triggered collapse and star formation due to the pressure of the impinging ionization front (Oort \\& Spitzer 1955, Kahn 1969; Dyson 1973; Dyson et al.\\ 1995; Elmegreen 1976; Bertoldi 1989; Bertoldi \\& McKee 1990; Bertoldi \\& Draine 1996; Lizano et al.\\ 1996; Gorti \\& Hollenbach 2002; Williams 1999; Williams et al.\\ 2001).  The evolution of these globules depends strongly on the incident UV radiation field, the initial size and density of the globule, and other properties, so observations of them in a wide variety of different environments are useful.  \\begin{table*}\\begin{minipage}{6in} \\caption{Observations of the Finger} \\begin{tabular}{@{}lllcl}\\hline\\hline Telescope &Instrument &Filter or     &Exp.\\ Time &Comment \\\\ &           &Emiss.\\ Lines &(sec)      &        \\\\ \\hline {\\it HST} &WFPC2   &F502N, [O~{\\sc iii}] $\\lambda$5007            &320  &image \\\\ {\\it HST} &WFPC2   &F656N, H$\\alpha$, [N~{\\sc ii}]                &400  &image \\\\ {\\it HST} &WFPC2   &F673N, [S~{\\sc ii}] $\\lambda\\lambda$6717,6731 &800  &image \\\\ CTIO 4m   &OSIRIS  &J                                      &120  &image \\\\ CTIO 4m   &OSIRIS  &H                                      &120  &image \\\\ CTIO 4m   &OSIRIS  &K                                      &120  &image \\\\ CTIO 4m   &OSIRIS  &He~{\\sc i} $\\lambda$10830              &720  &image \\\\ CTIO 4m   &OSIRIS  &Pa $\\beta$                             &720  &image \\\\ CTIO 4m   &OSIRIS  &H$_2$ 1-0 S(1) 2.122 $\\micron$         &1080 &image \\\\ CTIO 1.5m &RC Spec &blue; 3600-7100 \\AA &1200  &long-slit, P.A.=69$\\arcdeg$ \\\\ CTIO 1.5m &RC Spec &red;  6250-9700 \\AA &1200  &long-slit, P.A.=69$\\arcdeg$ \\\\ \\hline \\end{tabular}\\end{minipage} \\end{table*}  A rich population of these globules resides in the Carina Nebula (NGC~3372; d=2.3 kpc), with the most notable grouping clustered around the famous Keyhole Nebula.  Their emission from ionized, photodissociated, and molecular gas has been documented by several studies (Walborn 1975; Deharveng \\& Maucherat 1975; Cox \\& Bronfman 1995; Brooks et al.\\ 2000; Smith 2002; Rathborne et al.\\ 2002).  There are a dozen of these clumps around the Keyhole, with typical masses of order 10 M$_{\\sun}$ (Cox \\& Bronfman 1995).  A few of these are seen only in silhouette, but most have bright H$_2$ and polycyclic aromatic hydrocarbon (PAH) emission from their illuminated surfaces (Brooks et al.\\ 2000; Rathborne et al.\\ 2002).  In addition, several dozen smaller features with diameters of only $\\sim$5000 AU were recently discovered by Smith et al.\\ (2003a), and additional proplyd candidates are seen in our images (see below, and Barb\\'{a} et al.\\ in prep.). They may be either large photoevaporating circumstellar disks (proplyds) or very small cometary clouds that are remnants of larger globules in more advanced stages of photoevaporation (e.g., Bertoldi \\& McKee 1990).  It is not yet known if these smaller proplyd candidates or the larger globules in Carina contain embedded protostars or collapsing cores that may someday form stars.  However, one large globule in the southern part of the nebula has recently been shown to harbor a Class~{\\sc i} protostar that drives the HH~666 jet (Smith et al.\\ 2004).  In any case, these globules present a rare opportunity to study their associated phenomena in an environment powered by some of the hottest and most massive stars in the Galaxy (Walborn 1995; Walborn et al.\\ 2002), while still being located near to us and suffering little interstellar extinction.  In this paper we untertake a detailed analysis of one particularly striking bright-rimmed globule in the Carina Nebula (see Figure 1) that we call the ``Finger'' because of its gesticulatory morphology (it has previously been identified as clump 4; Cox \\& Bronfman 1995). The defiant finger seems to point toward its dominant UV source, and as such, allows for an interesting application of models of photoevaporating globules mentioned above.  In \\S 2 we discuss our observations, and in \\S 3 and \\S 4 we briefly discuss the results of our imaging and spectroscopy.  In \\S 5 we undertake a detailed analysis of the photoevaporative flow, and in \\S 6 we discuss implications for potential star formation in the globule.  \\begin{figure*}\\begin{center} \\epsfig{file=nsmith.fig2.eps,width=6in}\\end{center} \\caption{{\\it HST}/WFPC2 images of the Finger.  (a) F502N filter transmitting [O~{\\sc iii}] $\\lambda$5007; grayscale range from 2.4$\\times$10$^{-14}$ (white) to 1.9$\\times$10$^{-13}$ (black) ergs s$^{-1}$ cm$^{-2}$ arcsec$^{-2}$. (b) F656N filter transmitting H$\\alpha$; grayscale range from 5.4$\\times$10$^{-14}$ (white) to 3$\\times$10$^{-13}$ (black) ergs s$^{-1}$ cm$^{-2}$ arcsec$^{-2}$. (c) F673N filter transmitting [S~{\\sc ii}] $\\lambda\\lambda$6717,6731; grayscale range from 4.2$\\times$10$^{-15}$ (white) to 3.3$\\times$10$^{-14}$ (black) ergs s$^{-1}$ cm$^{-2}$ arcsec$^{-2}$. These observed flux levels were not corrected for reddening and extinction.  The axes display offset in arcsec from an arbitrary central position in the globule.} \\end{figure*}  \\begin{figure*}\\begin{center} \\epsfig{file=nsmith.fig3.eps,width=6in}\\end{center} \\caption{Near-IR images of the Finger.  (a) He~{\\sc i} $\\lambda$10830; grayscale range from 1.3$\\times$10$^{-14}$ (white) to 4.8$\\times$10$^{-14}$ (black) ergs s$^{-1}$ cm$^{-2}$ arcsec$^{-2}$. (b) Hydrogen Pa$\\beta$; grayscale range from 1.5$\\times$10$^{-14}$ (white) to 5.1$\\times$10$^{-14}$ (black) ergs s$^{-1}$ cm$^{-2}$ arcsec$^{-2}$. (c) H$_2$ $v=1-0$ S(1) $\\lambda$21218; grayscale range from 9.3$\\times$10$^{-16}$ (white) to 5.3$\\times$10$^{-15}$ (black) ergs s$^{-1}$ cm$^{-2}$ arcsec$^{-2}$. These observed flux levels were not corrected for reddening and extinction.} \\end{figure*}  \\begin{figure*}\\begin{center} \\epsfig{file=nsmith.fig4.eps,width=6in}\\end{center} \\caption{Selected flux-ratio images of the Finger.  (a) H$\\alpha\\div$[O~{\\sc iii}] (F656N$\\div$F502N) with the grayscale range from 1 (white) to 2.3 (black).  (b) [S~{\\sc ii}]$\\div$H$\\alpha$ (F673N$\\div$F656N) with the grayscale range from 0.033 to 0.13.  (c) The ground-based Pa$\\beta$ image divided by a smoothed {\\it HST}/WFPC2 H$\\alpha$ image, with the grayscale range from 0.07 to 0.12.  Note that the F656N H$\\alpha$ image is partially contaminated by adjacent [N~{\\sc ii}] lines, which affects the value of the flux ratios listed here.  These flux ratios have been corrected for reddening and extinction, using $E(B-V)$=0.37 and $R$=4.8 (see text).} \\end{figure*} ",
        "conclusions": "We have presented optical narrow-band {\\it HST}/WFPC2 images, ground-based optical spectra, and near-IR images of the ``Finger'' -- a photoevaporating molecular globule in the core of the Carina Nebula. The main conclusions of this work are the following:  1.  The Finger globule exhibits an interesting morphology, with a thin extended middle finger apparently pointing toward its source of ionizing photons.  The spatially-resolved structure of the stratified ionization fronts are consistent with the interpretation of the Finger as an optically-thick photoevaporating molecular globule, similar to structures often seen in {\\it HST} images of H~{\\sc ii} regions.  2.  Quantitatively, electron densities and the corresponding flux of ionizing photons incident upon the southward-facing ionization fronts of the Finger indicate that the dominant UV source is either WR25 (a late-type Wolf-Rayet star), Tr16-244 (O4 If), or perhaps both.  This is reassuring, since the finger points to within a few degrees of these stars.  3.  The mass-loss rate for the main evaporating surface of the globule is of order 2$\\times$10$^{-5}$ M$_{\\odot}$ yr$^{-1}$.  4.  From extinction measurements, we estimate an average hydrogen column density of a few times 10$^{22}$ cm$^{-2}$ through the globule, and a total mass (assuming a gas:dust mass ratio of 100:1) of at least 6 M$_{\\odot}$, in agreement with independent estimates from molecular studies.  This is an underestimate if the globule contains clumps that are optically-thick in the near-IR.  5.  The remaining lifetime of the globule before it is evaporated away is of order 10$^{5.3}$ to 10$^6$ years.  6.  Several reddened stars are seen projected within the boundaries of the Finger globule, but whether or not these sources are newly-formed stars that are physically associated with the globule is uncertain.  No reddened star is seen at the apex of the thin protruding finger or immediately behind the main ionization front, down to a limit of $m_K \\simeq 17$.  7.  Considering the properties of the advancing ionization front, it appears likely that stars are currently forming or will soon form in the globule, triggered by radiation-driven implosion.  At the main ionization front we find an external overpressure of a factor of $\\sim$5, and a smaller overpressure ($\\ga$1) at the end of the thin finger.    \\smallskip\\smallskip\\smallskip\\smallskip \\noindent {\\bf ACKNOWLEDGMENTS} \\smallskip \\scriptsize  We thank John Bally for supplying the ground-based [S~{\\sc ii}] image used in Figure 1, and we benefitted from helpful discussions with Paul Crowther regarding models of WR25.  Support was provided by NASA through grant HF-01166.01A from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS~5-26555.  NOAO funded N.S.'s travel to Chile and accommodations while at CTIO.  Some travel support was also provided by NASA grant NAG-12279 to the University of Colorado."
    },
    "0404/astro-ph0404206_arXiv.txt": {
        "abstract": "We apply the Gabor transform methodology proposed in (Hansen et al. 2002, 2003) to the \\emph{WMAP} data in order to test the statistical properties of the CMB fluctuation field and specifically to evaluate the fundamental assumption of cosmological isotropy. In particular, we apply the transform with several apodisation scales, thus allowing the determination of the positional dependence of the angular power spectrum with either high spatial localisation or high angular resolution (ie. narrow bins in multipole space). Practically, this implies that we estimate the angular power spectrum locally in discs of various sizes positioned in different directions: small discs allow the greatest sensitivity to positional dependence, whereas larger discs allow greater sensitivity to variations over different angular scales. In addition, we determine whether the spatial position of a few outliers in the angular power spectrum could suggest the presence of residual foregrounds or systematic effects. For multipoles close to the first peak, the most deviant local estimates from the best fit \\emph{WMAP} model are associated with a few particular areas close to the Galactic plane. Such deviations also include the ``dent'' in the spectrum just shortward of the first peak which was remarked upon by the \\emph{WMAP} team. Estimating the angular power spectrum excluding these areas gives a slightly higher first Doppler peak amplitude.  Finally, we probe the isotropy of the largest angular scales by estimating the power spectrum on hemispheres and reconfirm strong indications of a north-south asymmetry previously reported by other authors. Indeed, there is a remarkable lack of power in a region associated with the north ecliptic pole. With the greater fidelity in $\\ell$-space allowed by this larger sky coverage, we find tentative evidence for residual foregrounds in the range $\\ell=2-4$, which could be associated with the low measured quadrupole amplitudes and other anomalies on these angular scales (eg. planarity and alignment). However, over the range $\\ell=5-40$ the observed asymmetry is much harder to explain in terms of residual foregrounds and known systematic effects.  By reorienting the coordinate axes, we partition the sky into different hemispheres and search for the reference frame which maximises the asymmetric distribution of power. The north pole for this coordinate frame is found to intersect the sphere at $(80^\\circ,57^\\circ)$ in Galactic co-latitude and longitude over almost the entire multipole range $\\ell=5-40$. Furthermore, the strong negative outlier at $\\ell=21$ and the strong positive outlier at $\\ell=39$ as determined from the global power spectrum by the \\emph{WMAP} team, are found to be associated with the northern and southern hemispheres respectively (in this frame of maximum asymmetry). Thus, these two outliers follow the general tendency of the multipoles $\\ell=5-40$ to be of systematically lower amplitude in the north and higher in the south. Such asymmetric distributions of power on the sky provide a serious test for the cosmological principle of isotropy. ",
        "introduction": "One of the fundamental assumptions of cosmology is that the universe is isotropic. Such an  assumption necessarily implies therefore that the statistical properties of the cosmic microwave background (CMB) should be the same in all directions on the sky. The FIRAS experiment on the COBE satellite \\cite{cobe1,cobe2} demonstrated that the mean temperature of the CMB is isotropic to high precision, but it is only now with the superior sensitivity of the \\emph{WMAP}\\footnote{\\emph{Wilkinson Microwave Anisotropy Probe}} \\cite{WMAP} data that the isotropic nature of the angular fluctuations in the CMB can be tested.  In this paper we will consider the variation in the angular power spectrum of the CMB determined from patches sampling different locations on the sky. In an isotropic universe, there should be no preferred direction and these locally determined angular power spectra, after allowing for some sample variance, should demonstrate the same expectation value over all positions. Here, we adopt the Gabor transform method proposed in \\cite{hansen1,hansen2} and estimate the power spectrum in discs with 9.5 and 19 degree radius and then on hemispheres, thus allowing either high spatial localisation or high angular resolution, both of which are used to probe the directional dependence of the anisotropy. The discs on which we estimate the power spectrum are positioned such that they cover as much area as possible without overlapping so much as to introduce strong correlations between the locally estimated spectra. With the smaller discs, we can test the isotropy of the spectrum at a high spatial resolution, but small patches limit the resolution in multipole space. By increasing the size of the discs we test the isotropy of the spectrum at a refined resolution in multipole space but at the cost of a low spatial resolution. We apply different tests of consistency between the spectra in different multipole bins, comparing with a Monte Carlo ensemble generated with the best fit model spectrum obtained by the \\emph{WMAP} team and using the noise and beam properties of \\emph{WMAP}. We have chosen to focus on the co-added W+V map as these are the channels in which the foreground contamination is expected to be lowest. In all cases, only data outside of the \\emph{WMAP} Kp2 sky cut is included in the analysis.  With this approach we can also test for the presence of contaminating residual foregrounds or systematic effects. For example, incorrectly subtracted Galactic foregrounds would most likely manifest asymmetry in the Galactic frame of reference, with an increasingly likely level of residual contamination towards the plane. Similarly, systematic effects would plausibly align with the ecliptic frame of reference reflecting the scanning strategy of the satellite. In addition, contamination of the data due to solar system emission would be confined to the ecliptic plane. For these reasons, we compare the statistical distribution of the power spectrum estimates in both the Galactic and ecliptic northern and southern hemispheres, as well as the pole and equatorial regions for these reference frames.  We will in particular check the positional dependence of some features in the power spectrum observed by \\emph{WMAP} \\cite{hinshaw}: two at large scales corresponding to bins centered around $\\ell=21$ and $\\ell=39$, one close to the first Doppler peak and one close to the first trough (the latter was not seen in the analysis by \\emph{WMAP}). Moreover we investigate the recent claims of an asymmetry of the large scale structure between the northern and southern Galactic hemispheres \\cite{park,eriksen1,vielva,eriksen2,copi,curvat}. As the asymmetry is reported to be seen at the largest scales, we use the results of the power spectra estimated on hemispheres where the multipole resolution is large enough to resolve different features between $\\ell=2-63$.  The outline of the paper is as follows. In section \\ref{sect:gabtrans} we outline the method described in \\cite{hansen1,hansen2} to estimate the power spectrum on a patch on the sky using the so called Gabor transforms. The details of the data set used in the paper is outlined in section \\ref{sect:data}. In section \\ref{sect:small} we show the results of the local \\emph{WMAP} spectra estimated on discs of radius $9.5^\\circ$. We check for consistency between the locally estimated spectra in different multipole bins and different directions. To check the results at a higher multipole resolution, the analysis is repeated in section \\ref{sect:medium} with discs of radius $19^\\circ$. In section \\ref{sect:hemis} we check for an asymmetry in the large scale structure and report results of the \\emph{WMAP} spectra estimated on hemispheres. In particular, we test the robustness of the results using other frequency channels and a variety of sky cuts. Finally, section \\ref{sect:concl} summarise our results. ",
        "conclusions": "\\label{sect:concl}  In this paper, we have applied a method of local power spectrum estimation to small regions of the CMB sky as measured by the \\emph{WMAP} satellite in order to search for evidence of spatial dependence in the computed spectra beyond that due to sampling variance alone. In this way, we are able to confront one of the central tenets of modern cosmology, namely that of cosmological isotropy. As an important consequence of this analysis, we are also able to mitigate against foreground or systematic artefacts in the data. Specifically, we have estimated the angular power spectrum in different spatial directions within discs of radius $9.5^\\circ$ and $19^\\circ$ and on hemispheres. Since by increasing the disc size we increase the resolution of the power spectrum estimation in multipole space but lower the sensitivity to the spatial dependence of the estimate, the combination of these scales allows us to achieve high fidelity in both real and $\\ell$ space.  On small angular scales and specifically for two power spectrum bins, $\\ell=164-188$ and $\\ell=264-288$, close to the first Doppler peak, we found that the highest deviations, relative to the best fit model preferred by the \\emph{WMAP} team, was seen in the power spectrum estimates for those medium sized discs close to the Galactic plane. This was also the case for an interval in $\\ell$-space spanning the `dent' in the power spectrum as estimated by the global analysis of the \\emph{WMAP} team. By excluding those discs close to the Galactic plane with a local spectrum deviating by more than $2\\sigma$ from the best fit model, we find a resulting power spectrum with no dent and a slightly higher Doppler peak. We consider this to be a tentative indication of possible residual Galactic foreground contamination.  When studying the distribution of large scale power using the small and medium sized discs, we found that the northern Galactic and ecliptic hemispheres seemed to show a remarkable lack of large scale power. In order to further investigate this apparent lack of isotropy, we estimated the power spectrum for the first 60 multipoles on hemispheres centred on various directions on the sky. The multipole range $\\ell=2-4$ was found to be strongly asymmetric with the Galactic north-south axis as the axis of maximum asymmetry. Furthermore we found that the spatial dependence of the variations in amplitude for this multipole range was similar to the morphology seen in the integrated foreground template contribution, suggesting that the low order multipoles may remain compromised by residual Galactic emission (this was also noted by \\cite{oc,eriksen3,schwarz}). The other multipoles between $\\ell=5-40$ however, are asymmetric with the axis pointing in the direction $(80^\\circ,57^\\circ)$ in Galactic co-latitude and longitude, close to the ecliptic axis. In the northern hemisphere of the reference frame of maximum asymmetry, almost the full multipole range between $\\ell=5-40$ is below the average amplitude. The strong negative outlier at $\\ell=21$ found in the \\emph{WMAP} global estimate of the power spectrum, is associated specifically with this hemisphere. Conversely, in the southern hemisphere of maximum asymmetry, almost all multipoles in the range $\\ell=5-40$ are above the average amplitude, and a positive outlier at $\\ell=39$, as seen in the \\emph{WMAP} global spectrum, is found. The asymmetry does not appear to be associated with a particular anomalous region on the sky, but extends over a large area, as evidenced by the small and medium disc results.  We have checked the possibility that incorrectly removed foregrounds could cause the observed large scale asymmetry by testing the dependency on Galactic cut and frequency channel. The observed asymmetry is remarkably stable with respect to both frequency and sky coverage, thus arguing against this possibility. Moreover, the results seem unlikely to be compromised by systematics since we find supporting evidence (albeit at reduced statistical significance) for similar features in the \\emph{COBE}-DMR maps, which are susceptible to \\emph{different} parasitic signals. Similar asymmetric structures have been determined by a number of other groups \\cite{park,eriksen1,vielva,eriksen2,copi,curvat} using several alternative analysis techniques. Besides finding asymmetry, non-Gaussian features have also been detected in the northern and/or southern hemispheres. Since some of the non-Gaussian estimators used are power spectrum dependent, the detection of non-Gaussianity on opposite hemispheres may be related to the uneven distribution of large scale power studied in this paper. Given the large number of detections with different methods on different sky cuts and frequency channels, it seems inescapable that the \\emph{WMAP} data does indeed contain unexpected properties on large scales. In the absence of compelling evidence for a Galactic or systematic origin for the asymmetry, the intriguing possibility is raised that the cosmological principle of isotropy is violated and that fundamentally new physics on large scales in the universe is required. Further clarification of this scenario awaits further observations from \\emph{WMAP}, and ultimately the forthcoming Planck satellite mission.      \\clearpage  \\begin{figure} \\begin{center} \\leavevmode \\psfig {file=figure1.ps,width=10cm,height=7cm} \\caption{The position and numbering of the $9.5^\\circ$ discs on which the local power spectra have been calculated.} \\label{fig:discnumbers} \\end{center} \\end{figure}  \\clearpage  \\begin{figure} \\begin{center} \\leavevmode \\psfig {file=figure2a.ps,width=8cm,height=8cm} \\psfig {file=figure2b.ps,width=8cm,height=8cm} \\psfig {file=figure2c.ps,width=8cm,height=8cm} \\caption{The disc-disc correlation matrix for the multipole bins 0, 3 and 6 centred at  $\\ell=33$, $\\ell=189$ and $\\ell=339$ respectively.} \\label{fig:corplot} \\end{center} \\end{figure}  \\clearpage   \\begin{figure} \\begin{center} \\leavevmode \\psfig {file=figure3.ps,width=16cm,height=16cm} \\caption{Slice of the disc correlation matrix for some selected discs at the lowest multipole bin $\\ell=33$ and a bin centred close to the first power spectrum peak $\\ell=189$. We have chosen the discs 0,6 and 60 in order to study the correlations for discs which have many (disc 0) and few (disc 6 and 60) overlapping neighbours. The other discs have been chosen as particular bin values of these discs have been found in the \\emph{WMAP} data. We have chosen particular discs close to the ecliptic poles as some particular features are found in this area (see the text and figure \\ref{fig:plotdisc4_b034})}. \\label{fig:corcut} \\end{center} \\end{figure}  \\clearpage   \\begin{figure} \\begin{center} \\leavevmode \\psfig {file=figure4.ps,width=16cm,height=16cm} \\caption{Slice of the bin to bin correlation matrix for disc 0 taken at bins 1,4,5 and 9.} \\label{fig:binbincor} \\end{center} \\end{figure}  \\clearpage   \\begin{figure} \\begin{center} \\leavevmode \\psfig {file=figure5.ps,width=12cm,height=12cm} \\caption{The best fit \\emph{WMAP} running index power spectrum (solid line). The histogram shows the spectrum binned using the same bins as for the disc estimates. The crosses show the binned full sky \\emph{WMAP} spectrum obtained by the \\emph{WMAP} team and the shaded areas show the spread of the \\emph{WMAP} spectrum over 130 discs of radius $9.5^\\circ$, found from the Gabor analysis. The two shaded areas indicate where 67 and 95 percent of the 130 \\emph{WMAP} disc estimates are contained. Note that the shaded areas are NOT error bars for the \\emph{WMAP} estimates, but the spread found by estimating the spectrum in different positions.} \\label{fig:fullspectrum} \\end{center} \\end{figure}  \\clearpage   \\begin{figure} \\begin{center} \\leavevmode \\psfig {file=figure6.ps,width=12cm,height=12cm} \\caption{The best fit \\emph{WMAP} running index power spectrum (solid line). The histogram shows the spectrum binned in the same way as the disc estimates. The two shaded areas indicate where 67 and 95 percent of the disc estimates in the simulations are contained. The errorbars at each bin show the same distribution for the \\emph{WMAP} data, the left being the $1\\sigma$ level and the right being the $2\\sigma$ level.} \\label{fig:fulldist} \\end{center} \\end{figure}   \\clearpage   \\begin{figure} \\begin{center} \\leavevmode \\psfig {file=figure7a.ps,width=8cm,height=8cm} \\psfig {file=figure7b.ps,width=8cm,height=8cm} \\psfig {file=figure7c.ps,width=8cm,height=8cm} \\psfig {file=figure7d.ps,width=8cm,height=8cm} \\caption{The best fit \\emph{WMAP} running index power spectrum (solid line). The histogram shows the spectrum binned in the same way as the estimates on $9.5^\\circ$ discs. The two shaded areas indicate where 67 and 95 percent of the disc estimates in the \\emph{WMAP} map are contained. The errorbars at each bin show the same distribution (67 percent level) for a part of the \\emph{WMAP} data. {\\bf Upper left:} north Galactic hemisphere (left bar) / south Galactic hemisphere (right bar), {\\bf upper right:} Galactic polar region (left bar), Galactic equatorial region (right bar), {\\bf lower left:} north ecliptic hemisphere (left bar) / south ecliptic hemisphere (right bar), {\\bf lower right:} ecliptic polar region (left) / ecliptic equatorial region (right).} \\label{fig:partdist} \\end{center} \\end{figure}  \\clearpage  \\begin{figure} \\begin{center} \\leavevmode \\psfig {file=figure8.ps,width=12cm,height=12cm} \\caption{The disc estimates of the bin $\\ell=2-63$. The disc numbers correspond to the disc numbers in figure (\\ref{fig:discnumbers}). The shaded zones indicate the 1, 2 and 3 sigma spread of the estimated bin for the given disc as found from Monte Carlo simulations.} \\label{fig:plotdiscs2bin0} \\end{center} \\end{figure}  \\clearpage   \\begin{figure} \\begin{center} \\leavevmode \\psfig {file=figure9.ps,width=15cm,height=18cm} \\caption{The local power in the bin $\\ell=2-63$ (upper left plot), $\\ell=64-113$ (upper right plot), $\\ell=164-213$ (middle left plot), $\\ell=264-313$ (middle right plot) , $\\ell=364-413$ (lower left plot) and the bin $\\ell=414-463$ (lower right plot) estimated on $9.5^\\circ$ discs. The discs with power above the average set by simulations have yellow (below 1 $\\sigma$), red (between 1 and 2 $\\sigma$) and dark red (above 2 $\\sigma$) colour. The discs with power below the average have blue colour: light blue (above 1 $\\sigma$), blue (between 1 and 2 $\\sigma$) and dark blue (below 2 $\\sigma$). The two green discs show the positions of the ecliptic poles.} \\label{fig:plotdisc4_b034} \\end{center} \\end{figure}  \\clearpage     \\begin{figure} \\begin{center} \\leavevmode \\psfig {file=figure10.ps,width=12cm,height=12cm} \\caption{The joint likelihood estimation of 130 discs of radius $9.5^\\circ$. The histogram shows the mean and the shaded areas the 1 and 2 sigma levels from 1536 simulated maps with the same input spectrum. The crosses show the result of the same estimation procedure on the \\emph{WMAP} data. } \\label{fig:10degjoint} \\end{center} \\end{figure}  \\clearpage   \\begin{figure} \\begin{center} \\leavevmode \\psfig {file=figure11a.ps,width=8cm,height=8cm} \\psfig {file=figure11b.ps,width=8cm,height=8cm} \\psfig {file=figure11c.ps,width=8cm,height=8cm} \\psfig {file=figure11d.ps,width=8cm,height=8cm} \\caption{Results of the power spectrum analysis in different parts of the sky using a joint analysis of $9.5^\\circ$ discs in the Galactic reference frame (same as figure (\\ref{fig:10degjoint}) on a limited number of discs). The shaded zones indicate the 1 and 2 sigma errorbars calculated in each plot for one of the spectra shown. The errorbars of the other spectrum is very similar and is not shown. {\\bf Upper left plot:} black crosses:polar regions (discs 0-44, 66-110), gray crosses: equatorial regions (discs 29-65, 95-129). {\\bf Upper right plot:} black crosses: polar north (discs 0-44), gray crosses: polar south (discs 66-110). {\\bf Lower left plot:} black crosses: equatorial region north(discs 29-65), gray crosses: equatorial region south (discs 95-129). {\\bf Lower right plot:} black crosses: norther hemisphere (discs 0-65), gray crosses: southern hemisphere (discs 66-129).} \\label{fig:multi_galref} \\end{center} \\end{figure}  \\clearpage   \\begin{figure} \\begin{center} \\leavevmode \\psfig {file=figure12a.ps,width=8cm,height=8cm} \\psfig {file=figure12b.ps,width=8cm,height=8cm} \\psfig {file=figure12c.ps,width=8cm,height=8cm} \\psfig {file=figure12d.ps,width=8cm,height=8cm} \\caption{Results of the power spectrum analysis in different parts of the sky using a joint analysis of $9.5^\\circ$ discs in the ecliptic reference frame (same as figure (\\ref{fig:10degjoint}) on a limited number of discs). The shaded zones indicate the 1 and 2 sigma errorbars calculated in each plot for one of the spectra shown. The errorbars of the other spectrum is very similar and is not shown (this is not the case for the ecliptic polar regions versus ecliptic plane, but still the qualitative results are independent of the set of errorbars chosen). {\\bf Upper left plot:} black crosses:polar regions (distance to the poles less than $50^\\circ$), gray crosses: equatorial regions (distance from the poles larger than $50^\\circ$). {\\bf Upper right plot:} black crosses: polar north (distance to the north pole closer than $50^\\circ$), gray crosses: polar south. {\\bf Lower left plot:} black crosses: equatorial region north (distance from the poles larger than $50^\\circ$ but inside the northern hemisphere), gray crosses: equatorial region south. {\\bf Lower right plot:} black crosses: northern hemisphere, gray crosses: southern hemisphere.} \\label{fig:multi_eclref} \\end{center} \\end{figure}  \\clearpage   \\begin{figure} \\begin{center} \\leavevmode \\psfig {file=figure13.ps,width=10cm,height=7cm} \\caption{The position and numbering of the $19^\\circ$ discs on which the local power spectra have been calculated.} \\label{fig:discnumbers2} \\end{center} \\end{figure}    \\clearpage   \\begin{figure} \\begin{center} \\leavevmode \\psfig {file=figure14.ps,width=12cm,height=12cm} \\caption{The full sky \\emph{WMAP} power spectrum from a joint likelihood estimation of 34 discs of radius $19^\\circ$. The histogram shows the binned  \\emph{WMAP} best fit running index spectrum, the shaded areas the 1 and 2 sigma levels from 512 simulated maps with the same input spectrum. The crosses show the result of the joint disc estimation procedure on the \\emph{WMAP} data. } \\label{fig:20degjoint} \\end{center} \\end{figure}  \\clearpage   \\begin{figure} \\begin{center} \\leavevmode \\psfig {file=figure15.ps,width=15cm,height=18cm} \\caption{The local power in the bin $\\ell=2-63$ (upper left plot), $\\ell=64-113$ (upper right plot), $\\ell=164-188$ (middle left plot), $\\ell=214-238$ (middle right plot) , $\\ell=264-288$ (lower left plot) and the bin $\\ell=414-463$ (lower right plot) estimated on $19^\\circ$ discs. The discs with power above the average set by simulations are red and the discs with power below the average are blue.The intensity of the colour indicates whether the power is within $1\\sigma$ (light), between $1$ and $2\\sigma$ (medium) and outside of $2\\sigma$ (dark). The two green discs show the positions of the ecliptic poles.} \\label{fig:plotdisc4_20deg_b03} \\end{center} \\end{figure}   \\clearpage  \\begin{figure} \\begin{center} \\leavevmode \\psfig {file=figure16.ps,width=12cm,height=12cm} \\caption{The full sky \\emph{WMAP} power spectrum from a joint likelihood estimation of 19 discs of radius $19^\\circ$. The histogram shows the binned  \\emph{WMAP} best fit running index spectrum, the shaded areas the 1 and 2 sigma levels from 512 simulated maps with the same input spectrum. The crosses show the result of the joint disc estimation procedure on the \\emph{WMAP} data. } \\label{fig:20degjoint_exl} \\end{center} \\end{figure}  \\clearpage  \\begin{figure} \\begin{center} \\leavevmode \\psfig {file=figure17.ps,width=12cm,height=12cm} \\caption{The position and amplitude of the first peak. The solid line shows the best fit \\emph{WMAP} power spectrum. The shaded zones show the result of the full sky power spectrum estimation using the joint likelihood estimation procedure of 34 discs with radius $19^\\circ$. The zones indicate the $1$ and $2\\sigma$ spread of the peak position over 512 simulations, and the dark bold cross shows the result of the same procedure applied to the \\emph{WMAP} data. The white bold cross shows the peak position when some discs showing indications of Galactic foreground contamination are excluded. The crosses show the peak positions estimated on individual discs using the \\emph{WMAP} data, and the two circles show the corresponding $1$ and $2\\sigma$ spread over 34 discs in 512 simulations. Note that individual discs may have larger $1$ and $2\\sigma$ contours and are therefore not automatically outside of their individual contours even if they are outside of the circles in the plot.} \\label{fig:peak1} \\end{center} \\end{figure}  \\clearpage   \\begin{figure} \\begin{center} \\leavevmode \\psfig {file=figure18.ps,width=7.5cm,height=12cm} \\caption{The peak position (upper plot) and peak amplitude (lower plot) estimated on $19^\\circ$ discs. Red colour indicates that for the given disc, the estimate on \\emph{WMAP} was above the mean of simulations and blue colour shows that the \\emph{WMAP} estimate was below the mean. The significance is indicated by the intensity of the colour: light means within $1\\sigma$, medium means between $1$ and $2\\sigma$ and dark means outside of $2\\sigma$. The two green dots show the position of the ecliptic poles.} \\label{fig:discplot_peak} \\end{center} \\end{figure}  \\clearpage   \\begin{figure} \\begin{center} \\leavevmode \\psfig {file=figure19a.ps,width=7cm,height=7cm} \\psfig {file=figure19b.ps,width=7cm,height=7cm} \\psfig {file=figure19c.ps,width=7cm,height=7cm} \\psfig {file=figure19d.ps,width=7cm,height=7cm} \\psfig {file=figure19e.ps,width=7cm,height=7cm} \\psfig {file=figure19f.ps,width=7cm,height=7cm} \\caption{Results of power spectrum estimation on hemispheres: the left column shows the northern spectra, the right column shows the southern spectra. The first row is taken in the Galactic reference frame, the second row in the ecliptic reference frame and the last row in the reference frame of maximum asymmetry, where the north pole points at $(80^\\circ,57^\\circ)$. The solid line shows the best fit \\emph{WMAP} spectrum, the histogram is the same spectrum binned, the shaded zones show the one and two sigma error bars from simulations of the given hemisphere, grey crosses show the \\emph{WMAP} estimate on the full sky (with the Kp2 mask) and the black crosses show our estimate on the given hemisphere. Please note that the shaded zones show the errorbars for a hemisphere, the errorbars for the \\emph{WMAP} estimates (grey crosses) are smaller by a factor of roughly $\\sqrt{2}$.} \\label{fig:hemispec} \\end{center} \\end{figure}  \\clearpage    \\begin{figure} \\begin{center} \\leavevmode \\psfig {file=figure20.ps,width=7.5cm,height=18cm} \\caption{ The local power in the bin $\\ell=2-4$ (upper plot), $\\ell=20-22$ (middle plot) and $\\ell=38-40$ (lower plot)  estimated on hemispheres centred on the positions indicated by the discs. The hemispheres with power above the average set by simulations are indicated by red discs and the hemispheres with power below the average with blue discs. The significance is indicated by the intensity of the colour: light means within $1\\sigma$, medium means between $1$ and $2\\sigma$ and dark means outside of $2\\sigma$. The two green discs show the positions of the ecliptic poles.} \\label{fig:indbin} \\end{center} \\end{figure}  \\clearpage     \\begin{figure} \\begin{center} \\leavevmode \\psfig {file=figure21.ps,width=15cm,height=19cm} \\caption{The discs show the centres of the hemispheres where the power spectrum has been estimated. The colours show the ratio of the power spectrum bin $\\ell=2-4$ (upper left), $\\ell=5-16$ (upper right), $\\ell=17-28$ (middle left), $\\ell=29-40$ (middle left), $\\ell=41-63$ (lower left) and $\\ell=2-63$ (lower right) between the opposite hemispheres.} \\label{fig:allbins} \\end{center} \\end{figure}  \\clearpage  \\begin{figure} \\begin{center} \\leavevmode \\psfig {file=figure22.ps,width=12cm,height=12cm} \\caption{The disc estimates of the bin $\\ell=2-63$. The discs are ordered such that the values to the left are close to the north pole of the axis $(80^\\circ,57^\\circ)$ of maximum asymmetry and the values to the right are close to the corresponding south pole. The angular distance from this north pole is shown. The shaded zones indicate the 1, 2 and 3 sigma spread of the estimated bin for the given disc as found from Monte Carlo simulations.} \\label{fig:plotdiscs2max} \\end{center} \\end{figure}     \\clearpage  \\begin{figure} \\begin{center} \\leavevmode \\psfig {file=figure23.ps,width=12cm,height=12cm} \\caption{The percentage of maps with a higher asymmetry than in the \\emph{WMAP} data for a single multipole bin.} \\label{fig:binforbin} \\end{center} \\end{figure}  \\clearpage  \\begin{figure} \\begin{center} \\leavevmode \\psfig {file=figure24.ps,width=10cm,height=7cm} \\caption{The discs show the positions of the hemispheres with the 10 highest (black discs) and 10 lowest (white discs) bin values. The power spectrum bins considered were $\\ell=2-40$ (large discs), $\\ell=8-40$ (second largest discs), $\\ell=5-16$ (second smallest discs) and $\\ell=29-40$ (smallest discs).} \\label{fig:maxasspos} \\end{center} \\end{figure}  \\clearpage  \\begin{figure} \\begin{center} \\leavevmode \\psfig {file=figure25.ps,width=7.5cm,height=18cm} \\caption{The discs show the centres of hemisphere on which the power spectrum has been estimated. The colour in each disc represents the ratio between the power in the given hemisphere and the opposite hemisphere in the multipole range $\\ell=5-20$. The upper plot is for the \\emph{WMAP} W+V channel, the middle plot for the COBE data and the lower plot for the foreground template used for \\emph{WMAP}.} \\label{fig:wmap_vs_cobe} \\end{center} \\end{figure}  \\clearpage"
    },
    "0404/astro-ph0404030_arXiv.txt": {
        "abstract": "We present the international collaboration MINE (Multi-lambda Integral NEtwork) aimed at conducting multi-wavelength observations of X-ray binaries and microquasars simultaneously with the {\\it INTEGRAL} $\\gamma$-ray satellite. We will focus on the 2003 March--April campaign of observations of the peculiar microquasar GRS\\,1915+105 gathering radio, IR and X-ray data. The source was observed 3 times in the plateau state, before and after a major radio and X-ray flare. It showed strong steady optically thick radio emission corresponding to powerful compact jets resolved in the radio images, bright near-infrared emission, a strong QPO at 2.5\\,Hz in the X-rays and a power law dominated spectrum without cutoff in the 3--300\\,keV range. We compare the different observations, their multi-wavelength light curves, including JEM-X, ISGRI and SPI, and the parameters deduced from fitting the spectra obtained with these instruments on board {\\it INTEGRAL}. ",
        "introduction": "Microquasars are X-ray binaries that produce relativistic jets and thus appear as miniature replicas of distant quasars \\citep{mirabelrodriguez99}. Their emission spectra, variable with time, range from the radio to the $\\gamma$-ray wavelengths. We present here the first multi-wavelength campaign on GRS\\,1915+105 involving the {\\it INTEGRAL} satellite (3\\,keV--10\\,MeV). This campaign was conducted by the MINE (\\mbox{Multi-$\\lambda$} {\\it INTEGRAL} NEtwork, see {\\sf http://elbereth.obspm.fr/$\\sim$fuchs/mine.html}) international collaboration aimed at performing multi-wavelength observations of galactic X-ray binaries simultaneously with {\\it INTEGRAL}. ",
        "conclusions": "Here for the first time, we observed simultaneously all the properties of the \\emph{plateau} state of GRS\\,1915+105\\,: a powerful compact radio jet, responsible for the strong steady radio emission and probably for a significant part of the bright near-IR emission, as well as a QPO (2.5\\,Hz) in the X-rays and a power law dominated X-ray spectrum with a $\\Gamma$$\\sim$3 photon index up to at least 300\\,keV. Forthcoming works will study detailed fits of the X-ray spectra, to determine for example whether this power law is due to an inverse Compton scattering of soft disc photons on the base of the compact jet or not. In order to better understand the unusual behaviour of GRS\\,1915+105, we need to carry out similar simultaneous broad-band campaigns during the other states, in particular during the sudden changes that correspond to powerful relativistic ejection events."
    },
    "0404/astro-ph0404560_arXiv.txt": {
        "abstract": "{This paper reports on computational evidence for the formation of cloud-like dust structures around C-rich AGB stars. This spatio-temporal structure formation process is caused by a radiative/thermal instability of dust forming gases as identified by Woitke\\etal(2000)\\nocite{wsl2000}. Our 2D (axisymmetric) models combine a time-dependent description of the dust formation process according to Gail\\plus Sedlmayr (1988)\\nocite{gs88} with detailed, frequency-dependent continuum radiative transfer by means of a Monte Carlo method (Niccolini\\etal2003)\\nocite{nwl2003} in an otherwise static medium ($\\vec{v}\\!=\\!0$).  These models show that the formation of dust behind already condensed regions, which shield the stellar radiation field, is strongly favoured.  In the shadow of these clouds, the temperature decreases by several hundred Kelvin which triggers the subsequent formation of dust and ensures its thermal stability.  Considering an initially dust-free gas with small density inhomogeneities, we find that finger-like dust structures develop which are cooler than the surroundings and point towards the centre of the radiant emission, similar to the ``cometary knots'' observed in planetary nebulae and star formation regions. Compared to a spherical symmetric reference model, the clumpy dust distribution has little effect on the spectral energy distribution, but dominates the optical appearance in near IR monochromatic images. ",
        "introduction": "Dusty gases in space are often remarkably inhomogeneous.  Numerous observations of various dust forming objects like the circumstellar environments of AGB and post-AGB stars, R~Coronae~Borealis stars, planetary nebulae, the ejecta of novae and supernovae, and even the hot winds generated by Wolf-Rayet stars, have shown that the dust forming medium usually possesses a clumpy internal structure. Reviews of such observations are given by Lopez (1999) and Woitke (2001)\\nocite{woi2001,lop99}.  The best-studied object in that respect is probably the infrared carbon star \\object{IRC+10216}. Several infrared speckle observations of its innermost dust formation and wind acceleration zone show direct evidence for an irregular, possibly cloudy dust distribution around this late-type AGB star (Weigelt\\etal1998\\nocite{wbbfow98}, Haniff\\plus Buscher 1998\\nocite{hb98}, Tuthill\\etal2000\\nocite{tmdl2000}). Further evidence can be deduced from long-term $JHKL$ lightcurves. Concerning the carbon star \\object{II Lup}, Feast\\etal(2003)\\nocite{fwm2003} argue for a restricted epoch of dust formation in a limited region along the line of sight, in order to explain a large-amplitude long-term decrease in $J$ whereas no corresponding increase in $K$ and $L$ was detected. Direct IR imaging and multi-wavelength IR lightcurves of the oxygen-rich red giant \\object{L$_2$\\,Pup} point to similar wind asymmetries (Jura\\etal2002)\\nocite{jcp2002}.  Additional hints to inhomogeneities in the environments of AGB stars are given by the patchy SiO maser spots observed in oxygen-rich Mira variables (\\eg \\object{TX Cam}, Dia\\-mond\\plus Kemball 2003\\nocite{dk2003}). On a larger scale, CO rotational emission lines provide evidence for local density enhancements in the winds of late-type stars, \\eg concerning the carbon star \\object{TX Psc} (Heske\\etal1989)\\nocite{htm1989} or in high-resolution observations of the detached shell of \\object{TT Cyg} (Olofsson\\etal2000\\nocite{oblegb00}). Observations with higher spatial resolution, using instruments like the {\\sc Vlti, Ngst} or {\\sc Alma}, can be expected to reveal even more details in the near future.  Spherically symmetric dynamical models for AGB winds which include a time-dependent treatment of dust formation strongly suggest the formation of radial dust shells in more or less regular time intervals (\\eg Winters\\etal2000, Sandin\\plus H{\\\"o}fner 2003)\\nocite{wljhs2000,sh2003}, even if the pulsation of the star is neglected (e.g. Fleischer\\etal1995)\\nocite{fgs95}. The question arises whether these dust shells remain spherical symmetric (as suggested by the 1D models), or whether they might break apart into clouds due to instabilities. We focus in this paper on the second possibility and study the two-dimensional time-dependent behaviour of the dust/gas mixture just during the formation of a new dust shell.  Clumpiness may also play a vital role for the photo-chemistry during the AGB$\\,\\to\\,$PPN$\\,\\to\\,$PN transition phase. Opaque clumps can suppress the photodissociation of certain molecules like Benzene, such that their existence has been proposed to be an indicator for inhomogeneities with large density contrasts (Redman\\etal2003)\\nocite{rvcw2003}.  Regarding other classes of objects, numerous opaque structures have been discovered in proto-planetary nebulae (\\eg the Helix nebula \\object{NGC 7293} and the Eskimo nebula \\object{NGC 2392}) as well as in star formation regions (\\eg the \\object{Orion nebula}). These neutral, dense, probably dusty regions, designated as ``cometary knots'', ``globules'' or ``proplyds'', are located at the head of {\\sl radially aligned linear structures} in these nebulae and are particularly well visible in high resolution images of emission line ratios like [OIII]/H$\\alpha$ (O'Dell 2000\\nocite{ode2000}).  The physical cause of the observed clumpiness is still puzzling.  In Paper~I of this series (Woitke\\etal2000\\nocite{wsl2000}), we have formulated the hypothesis that a radiative/thermal instability in dust forming gases can provoke a self-organisation of the matter, which is possibly involved in the formation of the observed structures. This instability is characterised by a physical control loop between the radiative transfer, which determines the temperature structure of the medium, and the dust formation, which determines its opacity (see Fig.~1 in Paper~I).  This paper reports on computational evidence for this hypothesis. In Sect.~2, we outline the concept of our static model for the environment of a C-rich AGB star, which combines a time-dependent treatment of dust formation with two-dimensional radiative transfer.  Section~3 shows and discusses the results, including the calculated optical appearance of a clumpy dust distribution in the spectral energy distribution and in monochromatic images. In Sect.~4, our conclusions are drawn. ",
        "conclusions": "This paper has shown that the formation of dust shells in the circumstellar environments of late-type stars can be unstable. Spontaneous symmetry breaking may occur due to a radiative/thermal instability in the dust forming gas, which leads to the development of cloud-like dust structures close to the star.  These results have been obtained on the basis of time-dependent axisymmetric models, which combine a kinetic description of carbon dust formation/evaporation with detailed, frequency-dependent radiative transfer by means of a Monte Carlo method, in the static case. The simulations show that the dust preferentially forms behind already condensed regions, which shield the stellar radiation.  In the shadow of these clumps, the temperatures are lower by a few 100\\,K which triggers the subsequent formation and facilitates the survival of the dust close to the star.  As final result, numerous finger-like dust structures develop which may have an radial extension as large as $0.5\\,R_\\star$ and point towards the centre of the radiant emission, similar to the ``cometary knots'' observed in planetary nebulae and star formation regions.  The cloudy dust distribution has little effect on the calculated spectral energy distribution of the star, in comparison to a spherically symmetric model, but significantly influences the optical appearance of the circumstellar environment in near IR monochromatic images (\\eg at $\\lambda\\!=\\!2.2\\,\\mu$m).  In particular, an inhomogeneous dust distribution around the star leads to a likewise non-uniformly bright appearance of the stellar disk.  Comparable observations of late-type giants are commonly interpreted in terms of hot/cool spots on the stellar surface. However, our model suggests that a different physical explanation by dust is possible for these observations, even if the dust is barely visible in the image.  Due to computational time constraints, we have so far only been able to study the static case where velocity fields are ignored. In this case, the main feature of the model is the formation of a chemical wave which radially propagates outward, driven by dust formation on the outer edge and dust evaporation due to backwarming at the inner edge. The optical depth of this chemical wave reaches about $\\tau_{\\rm 1\\mu m}\\!\\la\\!1\\,...\\,3$ with a density-dependent propagation velocity of $\\approx\\!0.1\\rm\\,km/s\\,...\\,2\\rm\\,km/s$. The wave leaves behind a strongly inhomogeneous dust distribution close to the star, where the medium relaxes towards phase equilibrium. However, this process is unstable and results in the aforementioned simultaneous occurrence of cool dusty (optically thick) segments beside warmer, almost dust-free (optically thin) segments, through which the radiative flux finally escapes preferentially.  According to dynamical (but spherically symmetric) models of dust-forming AGB stars (\\eg Winters\\etal\\linebreak 2000, Schirrmacher\\etal2003, H{\\\"o}fner\\etal2003, Sandin \\& H{\\\"o}fner 2003)\\nocite{wljhs2000,sws2003,hgaj2003,sh2003}, the formation of dust mainly occurs in particular phases of the model triggered by the pulsation of the star, which results in radial dust shells. During such a dust shell formation phase, the effect of a re-evaporation from the inside is a typical feature, similar to the behaviour of our chemical wave. According to the present paper, this process should be unstable and might result in departures from spherical symmetry.   We believe that this instability can provide a basis for a better understanding of inhomogeneous dust distributions showing up in many observations. However, direct predictions for particular objects are difficult to make, because hydrodynamics is missing in the current model. Since the dust is blown away as soon as it forms, the system has only little time to relax towards phase equilibrium at the inner edge of the dust shell. On the one hand, this time may be too short to produce well-grown spatial dust structures as shown in this paper.  On the other hand, even small deviations from spherical symmetry may trigger important dynamical effects, \\eg \\begin{itemize} \\item The radiation pressure on dust grains, which is delivered to the gas via frictional forces, mainly depends on the radiative flux and the degree of condensation. If slightly more condensed regions exist, they might be accelerated outward, whereas less condensed regions stay behind or even fall back.\\vspace*{1.5mm} \\item Optically thick dust clouds will be confined by radiation pressure, because the bolometric radiative flux is larger at the inner edge of the cloud facing the star than at its self-shielded outer edge (this effect does not occur in spherically symmetric models where $r^2 F(r)=\\rm const$ in radiative equilibrium). Driven, pancake-like structures could evolve due to this effect. \\vspace*{1.5mm} \\item The hydrodynamical process of cloud acceleration due to radiation pressure is not well-studied, apart from the special case of spherical symmetry. This process may be dynamically unstable itself (\\eg Rayleigh-Taylor, Kelvin-Helmholtz). Velocity disturbances generated by these instabilities may have an important feedback on the dust forming medium. \\end{itemize} Thus, we propose a new hypothetical scenario for dust-driven AGB star winds: Excited by hydrodynamical, radiative or thermal instabilities, dust clouds are formed from time to time close to the star in temporarily shielded areas, which are accelerated outward by radiation pressure.  At the same time, thinner, dust-free matter falls back towards the star at different places. A highly dynamical and turbulent environment close to the star would be created in this way, which can be expected to bear again a strongly inhomogeneous dust distribution.  In order to verify this hypothetical scenario, much more elaborate model calculations would be required which may well exceed the present capabilities of parallel super-computers. Various processes must be traced in 3D, using dynamical models with detailed radiative transfer and time-dependent dust-chemistry."
    },
    "0404/astro-ph0404610_arXiv.txt": {
        "abstract": "{ We present \\XMM observations of $\\gamma^2$ Velorum (WR\\,11, WC8+O7.5III, $P$\\,=\\,78.53\\,d), a nearby Wolf-Ray binary system, at its \\xray high and low states. At high state, emission from a hot collisional plasma dominates from about 1 to 8 keV. At low state, photons between 1 and 4\\,keV are absorbed. The hot plasma is identified with the shock zone between the winds of the primary Wolf-Rayet star and the secondary O giant. The absorption at low state is interpreted as photoelectric absorption in the Wolf-Rayet wind. This absorption  allows us to measure the absorbing column density and to derive a mass loss rate $\\Mdot$\\,=\\,8$\\times$10$^{-6}$\\,M$_\\odot$yr$^{-1}$ for the WC8 star. This mass loss rate, in conjunction with a previous Wolf-Rayet wind model, provides evidence for a clumped WR wind. A clumping factor of 16 is required. The \\xray spectra below 1 keV (12~\\AA) show no absorption and are essentially similar in both states. There is a rather clear separation in that emission from a plasma hotter than 5\\,MK is heavily absorbed in low state while the cooler plasma is not. This cool plasma must come from a much more extended region than the hot material. The Neon abundance in the \\xray emitting material is 2.5 times the solar value. The unexpected detection of C~\\textsc{v} (25.3~\\AA) and C~\\textsc{vi} (31.6~\\AA) radiative recombination continua at both phases indicates the presence of a cool ($\\sim$\\,40,000~K) recombination region located far out in the binary system. ",
        "introduction": "}  The massive Wolf-Rayet binary system $\\gamma^2$~Velorum (WR\\,11, WC8+O7.5III, $P$\\,=\\,78.53\\,d) is an astrophysical laboratory in which many aspects of mass loss and wind-wind collision phenomena can be studied. The system is relatively nearby, its {\\sl Hipparcos} distance is $d\\,=\\,258\\pm35$\\,pc (van der Hucht et al. 1997; Schaerer et al. 1997). Both stars in the binary have been recently investigated with sophisticated model atmospheres and their stellar parameters are reasonably well known (De Marco \\& Schmutz 1999, De Marco et al. 2000).  The binary orbit has been re-determined by Schmutz et al. (1997) who combined recent with earlier observations (Niemela \\& Sahade 1980; Pike et al. 1983; Moffat et al. 1986; Stickland \\& Lloyd 1990). The orbit is mildly eccentric and has an inclination of 63$^\\circ$\\,$\\pm$\\,8$^\\circ$ (De Marco \\& Schmutz 1999). Because of the high orbital inclination, any emitting structures are seen through changing absorption columns as the stars revolve.   Since $\\gamma^2$~Vel is the nearest WR star, (see van der Hucht 2001), it is relatively bright and well observable at any wavelength, in particular in the X-ray domain.  It has been observed by all previous X-ray observatories, from the {\\sl Einstein} observatory (White \\& Long 1986; Pollock 1987) to {\\sl ASCA} (Stevens et al. 1996; Rauw et al. 2000), and, more recently, by {\\sl Chandra} (Skinner et al. 2001).  Its X-ray observational history has been reviewed by van der Hucht (2002) and Corcoran (2003).   \\begin{figure} \\hbox{\\hskip-0.2cm\\psfig{figure=35fig1.ps,height=7.7cm,width=8.5cm,clip=}} \\caption{Sketch of the $\\gamma^2$\\,Vel orbital configuration at the phases of our X-ray observations. The Wolf-Rayet star is at the center. Shape and orientation of the wind blown cavities around the O star are schematic only.} \\label{figure1} \\end{figure}   With a series of \\ROSAT observations covering the binary orbit of $\\gamma^2$\\,Vel, Willis et al. (1995) discovered that the X-ray emission is a factor of $\\sim$\\,4 enhanced during a brief time span when the O-type component is in front.  They also showed that the steep increase takes place only in \\ROSAT's hard \\xrays. They convincingly interpreted the variable \\xray emission to arise from colliding stellar winds.  The enhancement is explained by the viewing geometry, when the collision zone can be seen through a rarefied cavity that builds around and behind the O-type component (see Fig.\\,\\ref{figure1}). At other phases the dense WR wind absorbs the X-rays from the collision zone. The wind blown cavity is generally orientated away from the WC component but it is also somewhat warped because of the binary motion of the O star.  Here we present \\XMM observations of $\\gamma^2$\\,Vel, taken at two phases. The first phase is at the maximum X-ray flux, a few days after the O-type component passed in front of the WR star.  The second phase is intermediate between quadrature and superior conjunction. In this configuration the O star is seen through a large portion of the extended WR atmosphere (Fig.~\\ref{figure1}).  After describing the observations and the most interesting spectral features we analyze the data in two different ways. First, we simply take the \\xray emitting zone as a source of light with which the WR wind is irradiated.  The observed absorption changes between different orbital phases provide unique information about the structure of the WR wind.  Secondly, we interpret the X-ray emission at both phases by a spectral fitting procedure. This reveals new insights into the geometric and thermal structure and the elemental composition of the wind-wind collision zone. ",
        "conclusions": "High-resolution X-ray spectra obtained at different orbital phases provide a wealth of information about $\\gamma^2$\\,Vel.  Modeling the X-ray emission constrains the physical structure of the wind-wind collision zone, whereas the absorption observed at non-maximum phases gives indications about the geometric distribution of the emitting as well as the non-emitting material.  Both, emission and absorption are important and reveal different but linked aspects of the $\\gamma^2$\\,Vel system. It is indeed likely that a comprehensive tomographic analysis using \\xray spectra taken at many more orbital phases will allow a detailed mapping of the colliding wind region as well as of the ambient material. In particular, the hypothesis of a constant clumping factor around the orbit could be tested.   \\subsection{WR mass loss and wind clumping}  Phase dependent \\xray emission from  $\\gamma^2$\\,Vel can be used to analyze the Wolf-Rayet wind. In order to quantitatively interpret the absorption at low state, we apply a previously published WR model atmosphere with a smooth density distribution (De\\,Marco et al 2000). This model atmosphere is the result from a fit to the broad WR emission lines. The column density required by the observed \\xray absorption is a factor of 4 lower than what is predicted by this model. The mass loss rate that matches the \\xray absorption is correspondingly smaller. We conclude that the WC8 star in  $\\gamma^2$\\,Vel loses mass at a rate of only 8$\\times$10$^{-6}$\\Msolar/yr.  The discrepancy between our directly measured mass loss rate and the one required by the model atmosphere can be reconciled if the wind is clumped. In order to still fit the WR emission line spectrum with the reduced mass loss a clumping factor f = 16 is required.  \\begin{figure} \\hbox{\\hskip-0.2cm\\psfig{figure=35fig9.ps,height=7.7cm,width=8.5cm,clip=}} \\caption{Sketch of the $\\gamma^2$\\,Vel system (same as Fig.\\,1) with the probable locations of our components 1 to 4. The Wolf-Rayet star is at the center.} \\label{figure9} \\end{figure}   \\subsection{Size of line emitting regions}  The observed absorption behaviour also constrains the geometry of the \\xray line emitting region. It is very interesting that in our spectra the separation between absorption and no absorption is quite sharp. While the \\Neix lines remain unabsorbed the \\Nex lines are reduced by a factor of 5 (see Tab. 2). In terms of temperature this means that the plasma hotter than 5 MK is heavily absorbed at phase 0.37 while the cooler plasma is not.  This is also reflected in our emission model in which the components 1 and 2 with temperatures of 8 and 19 MK are strongly absorbed whereas component 3 with a temperature of 3 MK is only weakly absorbed. We conclude that components 1 and 2 are formed in the central part of the colliding winds which is deeply embedded in the WR wind. The cool (3\\,MK) component is clearly detached from this hot region (see Fig.\\,9).  Furthermore, the \\Neix lines that predominantly come from this region are not affected by the UV radiation of the O star. They either  are formed far away from the O star or they are shielded from that UV radiation by intervening material. In either case, the O star is not likely to contribute much to them and we conclude that firstly, this component is associated with the colliding winds and secondly that this region must be rather extended for it to still be well detectable at phase 0.37.  \\subsection{Neon abundance} A further interesting piece of evidence comes from the neon abundance we derive. Neon is considerably enriched through nuclear processing in WC stars and therefore differs significantly from solar abundance.  Neon is in fact the only element that emits copious line radiation from the collisionally ionized region, and that allows to discriminate between WC and solar composition. We find a clear  Neon enhancement compared to solar which indicates that Wolf-Rayet material is present at least in components 2 and 3.  It is noteworthy that much of what we learn about the collision zone actually comes from the \\Neix and \\Nex lines. Apart from discriminating between WC and solar abundance patterns they also provide a dividing line between absorption and no absorption at phase 0.37. These lines seem to hold the key for further progress and their behaviour at other phases should certainly be very interesting to follow.   \\subsection{\\xray emission variability} An interesting feature of our \\xray spectra is the high energy end. The section above 4.5\\,keV is little affected by intrinsic absorption and interestingly there is no observable difference between the two phases. In particular the highest temperatures in the wind-wind collision zone seem to be the same at both phases. This is remarkable because the binary separation has changed from 0.83 AU at phase 0.12 to 1.27 AU at phase 0.37. From a simple 1/D law one would expect a 50\\% flux reduction (see e.g., Stevens et al. 1992). This confirms the result of Rauw et al. (2000) that $\\gamma^2$\\,Vel  does not follow a 1/D distance relation. The \\xray  flux and temperature from the hottest plasma as detected by \\XMM is not affected by the orbital separation.  \\subsection{Recombining plasma} Apart from the shock excited components we also find a recombining plasma. The relation (if any)  between this fourth component to the shocked material is not clear. The recombining plasma is highly ionized and is not absorbed at phase 0.37. It therefore comes from far out in the binary system. We also know that it is of very low temperature of about 40\\,000\\,K and that it is of WC composition. Possibly this plasma is due to photoionization through the X-rays from the wind-wind collision region. This radiation propagates through the rarefied and warped cavity behind the O star and irradiates the higher regions of the WR wind. There it may re-ionize some of the material."
    },
    "0404/astro-ph0404426_arXiv.txt": {
        "abstract": "Despite many years of effort, observational studies have not found a strong correlation between the presence of any proposed fueling mechanism and low-luminosity AGN. After a discussion of the mass requirements for fueling, I summarize this observational work and provide a number of hypotheses for why the nature of AGN fueling has remained unresolved. In particular, I stress the potential importance of the increasing number of candidate fueling mechanisms with decreasing mass accretion rate, the relevant spatial scales for different fueling mechanisms, and the lifetime of an individual episode of nuclear accretion. The episodic AGN lifetime is a particularly relevant complication if it is comparable to or shorter than the time that the responsible fueling mechanisms are observationally detectable. I conclude with a number of relatively accessible areas for future investigation. ",
        "introduction": "One of the main, unsolved problems in AGN research is how the AGN fuel is transported to the central, supermassive black hole. For low-luminosity AGN, the likely source of this fuel is the host galaxy itself, and in particular the ISM. As most of this material is distributed in a rotating disk extending out to kiloparsec scales, the problem of AGN fueling is essentially a problem of angular momentum transport: All but approximately one part in $10^7$ of the angular momentum must be removed for material to flow from kiloparsec scales to the event horizon on AU scales.  The problem has remained unsolved in part because the most relevant spatial scales for fueling, the central parsec and inward, are only observable in the very nearest galaxies. An angular resolution of $0.1''$ only corresponds to 1~pc or better spatial resolution out to a distance of 2~Mpc. Even the central 10~pc are not readily resolved for large samples of nearby AGN. The present technical limitations on angular resolution have therefore made it impossible to study AGN fueling directly and instead driven investigations to indirect, statistical studies of active and inactive galaxy samples in order to identify the mechanism(s) responsible for fueling accretion.  \\subsection{Requirements for Fueling}  Before consideration of these investigations, it is valuable to first consider the mass accretion rates estimated to produce the observed population of low-luminosity AGN, which are defined here to be any AGN less luminous than a QSO. The luminosity of an AGN is related to the mass accretion rate by $L = \\epsilon \\dot{M} c^2$, where $\\epsilon$ is the radiative efficiency of the accretion. This accretion is commonly assumed to occur via a Shakura-Sunyaev thin disk (Shakura \\& Sunyaev 1973) with a constant radiative efficiency $\\epsilon = 0.1$. To maintain Eddington accretion onto a supermassive black hole requires a mass accretion rate of \\begin{equation} \\dot{M}_{Edd} = 0.2 \\epsilon^{-1}_{0.1} \\left( \\frac{M}{10^7 {\\rm M}_{\\odot}} \\right) {\\rm M}_{\\odot} {\\rm yr}^{-1} \\end{equation} where $\\epsilon_{0.1} = \\epsilon/0.1$. The mass accretion rate is commonly parametrized in terms of the Eddington rate as the dimensionless accretion rate $\\dot{m} \\equiv \\dot{M}/\\dot{M}_{Edd}$, similarly the bolometric luminosity can be expressed as $l \\equiv L_{bol}/L_{Edd}$. For a standard accretion disk $l = \\dot{m}$, while lower efficiency ADAF models, which likely become important below $\\dot{m} = \\dot{m}_{crit} \\sim 0.01$, predict $l \\propto \\dot{m}^2$ from a scaling of $\\epsilon_{0.1} = \\dot{m}/\\dot{m}_{crit}$ (Narayan, Mahadevan, \\& Quataert 1998). The required mass accretion rates from fueling may be even higher if ADIOS/CDAF models are important, as these models predict that only a small fraction of the supplied mass is actually accreted by the black hole (Blandford \\& Begelman 1999). The key importance of these relations is that the mass accretion rates required for low-luminosity AGN do not decline as rapidly as their luminosity.  The most luminous low-luminosity AGN, Seyfert 1s, appear to have central, supermassive black holes with $M_\\bullet = 10^7 {\\rm M}_{\\odot}$ (Ferrarese et al.\\ 2001) and $l = 0.1$, which corresponds to accretion at $\\dot{m} = 0.1$ or $\\dot{M} = 0.02\\,{\\rm M}_{\\odot}$/yr. More typical Seyferts, which constitute a total of approximately 10\\% of the luminous galaxy population (Ho, Filippenko, \\& Sargent 1997) likely accrete with $\\dot{m} = 0.01$, while the 30\\% of the luminous galaxy population that are LINERs likely accrete with $\\dot{m} = 10^{-2} - 10^{-4}$ based on estimates of $l$ (Ho 1999). If all luminous galaxies go through periodic episodes as AGN, then over a time period of $10^8$ years\\footnote{This time period is only adopted to illustrate the potential total fuel requirements of a low-luminosity AGN. The relative numbers of Seyferts and LINERs only constrain the duty cycle of these phases and not the lifetime}, a $10^7 {\\rm M}_{\\odot}$ black hole will appear as a Seyfert galaxy for a total of $10^7$ yr and accrete on order $10^{3} {\\rm M}_{\\odot}$, appear to be a LINER for $3 \\times 10^7$ yr and accrete $10^{3} {\\rm M}_{\\odot}$, and appear to be inactive for $6 \\times 10^7$ yr. Therefore the mass inflow rates and total mass reservoirs required to power low-luminosity AGN are relatively meager.  \\subsection{Proposed Fueling Mechanisms}  Figure~1 lists the many mechanisms proposed to drive angular momentum transport in the host galaxy and provide fuel to the central parsec. These mechanisms can be divided between gravitational and hydrodynamic mechanisms. Gravitational mechanisms, such as galaxy interactions (Toomre \\& Toomre 1972) and large-scale bars (Simkin, Su, \\& Schwarz 1980), remove angular momentum through torques, while hydrodynamic mechanisms, such as turbulence in the ISM (Elmegreen et al.\\ 1998), remove angular momentum through gas dynamical effects. Many of these mechanisms are discussed in the review by Shlosman, Begelman, \\& Frank (1990) and the more recent review by Wada (2004). The latter in particular provides an excellent overview of recent theoretical work on hydrodynamic fueling mechanisms and places particular emphasis on high resolution simulations of the multiphase ISM in the central kiloparsec.  \\begin{figure}% \\includegraphics{martini.f1.eps} \\caption{\\small Mechanisms proposed to fuel accretion onto black holes at the centers of galaxies. The fueling mechanisms are approximately ordered by their expected maximum accretion rate. Progressively larger numbers of progressively more common mechanisms may be responsible for supplying the lowest rates of accretion. This suggests that while the relevant question for luminous AGN may be ``Why are active galaxies active?'' the question should be ``Why are inactive galaxies inactive?'' at the lowest accretion rates. } \\end{figure}  While a large number of mechanisms have been proposed for AGN fueling, the vast majority are only likely to provide relatively low mass accretion rates. Mergers between galaxies, particularly major mergers, is the mechanism most commonly invoked to explain the high accretion rates required to power luminous QSOs, while at the somewhat lower mass accretion rates responsible for Seyfert-level luminosities, mechanisms such as bars and minor mergers have also been considered. A progressively larger number of candidate mechanisms could be important at yet lower mass accretion rates. The mechanisms listed in Figure~1 are approximately ordered by their relative {\\it maximum} mass accretion rates.  As the mass accretion rate to the nucleus is expected to gradually taper off, rather than stop abruptly, mechanisms invoked to explain high mass accretion rates may also produce low mass accretion rates as the fuel supply is gradually depleted.  \\subsection{Observational Searches}  As the spatial scales most important for fueling are currently inaccessible in most galaxies, searches have instead employed large samples of AGN and inactive, control galaxies to identify the fueling mechanism(s). Ideally, a candidate fueling mechanism should only be found in the AGN sample. These experiments are challenging in that they require both large samples of galaxies and very careful control of systematic effects in the selection of both the AGN and the control samples. Claims over the years of statistically significant excesses have all diminished after reevaluation of sample sizes and the details of the AGN and control sample selection (e.g.\\ different distance distributions, host galaxy types). Sample selection remains a very important consideration as progressively more AGN are found in progressively more sensitive surveys (e.g.\\ Ho et al.\\ 1997) and few studies of AGN fueling use samples selected with uniform and unbiased criteria, such as hard X-ray luminosity, nor do they have comparably sensitive observations of the inactive, control sample.  The two main and most readily observed `large-scale' mechanisms are bars and interactions. Neither of these features is seen in significant excess in AGN samples compared to carefully-matched control samples (Fuentes-Williams \\& Stocke 1988; Mulchaey \\& Regan 1997; Schmitt 2001; see also Schmitt, {\\it these proceedings}). One explanation for these results is that AGN fueling is predominantly mitigated by smaller-scale phenomenon than accessible in these ground-based surveys (e.g.\\ Martini \\& Pogge 1999), a point illustrated in part by Figure~1. While the fueling mechanisms are approximately ordered by accretion rate, the mechanisms proposed to produce lower accretion rates are also progressively smaller-scale phenomena and more difficult to identify observationally. The desire for finer spatial resolution motivated a careful study of the circumnuclear region (on 100 pc scales) with {\\it HST}, but still no significant differences between AGN and control samples were found (Martini et al.\\ 2003a). ",
        "conclusions": ""
    },
    "0404/nlin0404016_arXiv.txt": {
        "abstract": "We consider a solvable model of the decay of scalar variance in a single-scale random velocity field. We show that if there is a separation between the flow scale $\\ko^{-1}$ and the box size $\\kbox^{-1}$, the decay rate $\\lambda\\propto(\\kbox/\\ko)^2$ is determined by the turbulent diffusion of the box-scale mode. Exponential decay at the rate $\\lambda$ is preceded by a transient powerlike decay (the total scalar variance $\\sim t^{-5/2}$ if the Corrsin invariant is zero, $t^{-3/2}$ otherwise) that lasts a time $t\\sim1/\\lambda$. Spectra are sharply peaked at $k=\\kbox$. The box-scale peak acts as a slowly decaying source to a secondary peak at the flow scale. The variance spectrum at scales intermediate between the two peaks ($\\kbox\\ll k\\ll\\ko$) is $\\sim k + a k^2 +\\dots$ ($a>0$). The mixing of the flow-scale modes by the random flow produces, for the case of large P\\'eclet number, a $k^{-1+\\delta}$ spectrum at $k\\gg\\ko$, where $\\delta\\propto\\lambda$ is a small correction. Our solution thus elucidates the spectral make up of the ``strange mode,'' combining small-scale structure and a decay law set by the largest scales. ",
        "introduction": "The problem of the decay of passive-scalar variance has recently been reexamined in the literature following the realization that the decay rates, spectra, and higher-order statistics based on small-scale Lagrangian-stretching theories \\cite{Antonsen_etal,Son,Balkovsky_Fouxon,Falkovich_etal_review} are not consistent with either numerical \\cite{Pierrehumbert_strange_mode,Pierrehumbert_pdfs,Fereday_etal,Sukhatme_Pierrehumbert} or experimental \\cite{Rothstein_etal,Voth_etal} results in the long-time limit. Instead, the scalar decay is dominated by an eigenmodelike solution dubbed ``the strange mode'' \\cite{Pierrehumbert_strange_mode} because it combines intricate small-scale structure with globally determined decay rate and self-similar statistics (self-similarity is also seen in numerical simulations of the related problem of kinematic dynamo \\cite{SCMM_ssd}). There has been a growing understanding \\cite{Voth_etal,Thiffeault_Childress,Chertkov_Lebedev,Fereday_Haynes,Sukhatme_pdf} that the overall decay rate is set by the slowest-decaying system-scale modes. This brings to mind homogenization theory \\cite{Majda_Kramer}, which considers the turbulent diffusion of passive scalar at scales much larger than the flow scale and where it is the largest-scale mode that decays most slowly. In this paper, we use a simple solvable example to demonstrate that the strange-mode decay rate is the rate of turbulent diffusion of the box-scale mode and show how the spectra of scalar variance accommodate both this box-scale diffusion and small-scale structure.  Qualitatively, the key idea quantified by our theory is as follows. A scalar field whose variance is at the scale smaller than or equal to the scale of the ambient random flow is mixed at a rate determined by the Lyapunov exponent of the flow --- this is the Lagrangian-stretching approach. However, if the size of the box is larger than the scale of the flow, the scalar field can have variance at the scale of the box. The rate of transfer of this large-scale variance to the flow scale (turbulent diffusion) can be much smaller than the Lagrangian mixing rate, in which case this slow transfer sets the global decay rate.  Our model emphasizes scale separation between the box and the flow. Our results are complementary to \\cite{Haynes_Vanneste}, where the decay of a scalar field is studied with more generality (in two dimensions).  We consider the advection-diffusion equation \\bea \\label{ADEq} \\dd_t\\theta + \\vu\\cdot\\nabla\\theta = \\eta\\Delta\\theta, \\eea with a random Gaussian white-in-time velocity field $\\la u^i(t,\\vx)u^j(t',\\vx')\\ra = \\delta(t-t')\\kappa^{ij}(\\vx-\\vx')$ known as the Kraichnan model \\cite{Kraichnan1}. The mean scalar concentration has been subtracted --- i.e., $\\la\\theta\\ra=0$. For the Kraichnan velocity, the angle-integrated scalar-variance spectrum in $d$ dimensions $T(k)=\\int\\diff\\Omega_\\vk\\, k^{d-1} \\la|\\theta(\\vk)|^2\\ra$ satisfies an integro-differential equation valid at all~$k$~\\footnote{The derivation is analogous to the standard one in the dynamo theory: see, e.g., \\cite{SBK_review} and references therein. Note that the $x$ space version of \\eq{MCEq} is local, but we stay with the integral equation because we are interested in spectra. For $x$ space calculations, see \\cite{Chertkov_Lebedev,Haynes_Vanneste}.}: \\bea \\nonumber \\dd_t T(t,k) &+& (2\\eta + \\kappa_0) k^2 T(t,k) \\\\ \\label{MCEq} &=& k_i k_j\\int\\diff^d k'\\kappa^{ij}(\\vk-\\vk')\\,T(t,k'), \\eea where $\\kkappa^{ij}(\\vk)=\\kkappa(k)\\bl(\\delta^{ij}-k_i k_j/k^2\\br)$ is the Fourier transform of $\\kappa^{ij}(\\vx-\\vx')$ and $\\kappa_0=(1/d)\\kappa^{ii}(0)$ is the turbulent diffusivity ($d$ is the dimension of space).  In \\secref{sec_small_scales}, we review the theory of scalar decay at small scales, which leads to the standard Lagrangian-stretching results. In \\secref{sec_main}, the theory for a finite-scale flow is developed --- this is the main part of the paper. Concluding remarks are in \\secref{sec_disc}. ",
        "conclusions": "\\label{sec_disc}  We now have a physical picture of the ``strange mode'': the small-$k$ peak (singularity at the box scale) serves as a slowly decaying source to the flow-scale mode (singularity at $k=1$), which, in turn, is mixed by the random flow and thus excites the nonsingular modes at small [\\eq{FPSln}] and large [\\eq{T_series}] scales. The structure of the spectrum is illustrated in \\figref{fig_strange_mode}.  The low-wave-number behavior of the decaying scalar field was previously analyzed in a heuristic way by Kerstein and McMurtry~\\cite{Kerstein_McMurtry} (see also \\cite{Gonzalez} for a treatment based on one of the turbulence closure schemes, which gives mostly similar results). They considered advection by a narrow-band (i.e., single-scale) forced random flow in an unbounded domain --- i.e., in the regime that we call the transient (powerlike-decay) stage. They recognized the defining role of coupling between the large scales ($k\\lesssim\\kpeak$) and the flow scale ($k=\\ko=1$) and derived the $k^4$ scaling at $k\\ll\\kpeak$ with a exponential fall off at $k>\\kpeak$ [\\eq{HomSln_C2}] and the ensuing powerlike-decay laws for the case of $C_2=0$ (see the end of \\secref{sec_C2_zero}). For the intermediate rage $\\kpeak\\ll k\\ll1$, they predicted a $k^2$ spectrum (in 3D) --- in contrast to our $k^1$ result [\\eqsand{Tint_decay_C2}{T_series} and \\secref{sec_C2_zero}]. The reason for the discrepancy is as follows. The analysis of \\cite{Kerstein_McMurtry} is based on Taylor-expanding the flow around $k=1$ --- i.e., in terms of our theory --- setting $T(t,1+q)\\simeq T(t,1)$ in \\eq{HomEq}, which gives $S(t,k)\\simeq k^4 T(t,1)$. If we had used the resulting equation to solve for $T(t,k)$ at $\\kpeak\\ll k\\ll 1$, we would also have obtained $T(t,k)\\sim k^2$. However, as we have seen above, the width of the flow-scale singularity is $\\sim\\kpeak$ [\\eqsand{T1_decay_C2}{T1_decay_zeroC2}], so Taylor expansion cannot be used in \\eq{HomEq} for $k\\gg\\kpeak$. In this intermediate range, the integral in \\eq{HomEq} must instead be replaced by the integral over the entire flow-scale peak, resulting in our $k^1$ scaling. The $k^2$ term enters as a correction due to the interaction between nonsingular modes [\\eq{T_series}].  Finally, let us comment on our modeling assumptions. The white-noise approximation might appear drastic: the correlation time of any realistic flow is comparable to the flow time scale $\\sim(u\\ko)^{-1}\\sim\\kappa_2^{-1}$. However, since the scalar decay time is much longer than the flow time scale (provided $\\kbox\\ll\\ko$), the white-noise model appears reasonable. We believe it also correctly captures the small-scale structure: the key factor here is the statistics of fluid displacements, which are integrals of velocity and are finite-time correlated even for a white-in-time velocity.  Our model flow was single scale. Although such flows can be set up in the laboratory \\cite{Rothstein_etal,Voth_etal} \\footnote{It was pointed out to us by Kerstein \\cite{Kerstein_pc} that a single-scale random flow could also be set up by randomly stirring a granular material: studying mixing in such a flow would provide an interesting experimental test.}, the real-world mixing problems usually contain (at sufficiently small scales) a wide (inertial) scale range of three-dimensional turbulent motions. While the variance spectrum in the inertial range should follow the Obukhov-Corrsin $k^{-5/3}$ law \\cite{Obukhov_scalar,Corrsin_spectrum} and there will be another transient powerlike-decay stage \\cite{Corrsin_inv,Chasnov,Eyink_Xin,Chaves_etal,Chertkov_Lebedev}, the long-term decay (after the scalar variance reaches $k<\\ko$) should still be qualitatively described by our theory. Another modification that results from the relaxation of the single-scale assumption concerns the intermediate wave-number range $\\kpeak\\ll k\\ll\\ko$. As noted in \\cite{Kerstein_McMurtry} and confirmed in pipe-flow mixing experiments \\cite{Guilkey_etal_spectra}, the interaction between the $k=\\kpeak$ mode and the low-wave-number tail of the kinetic-energy spectrum ($\\sim k^4$) can change the scaling of the scalar-variance spectrum in this range.  In conclusion, we emphasize that, in any laboratory experiment aiming to test our results, the stirring must be done at scales substantially smaller than the system size to ensure that $\\kbox\\ll\\ko$. It was just such a set up (in 2D) that allowed Voth {\\em et al.} \\cite{Voth_etal} to show experimentally that the global mixing rate was much smaller than that predicted by the Lagrangian-stretching theories and consistent with the box-scale turbulent-diffusion rate --- precisely the point the theory presented above is meant to demonstrate."
    },
    "0404/nucl-th0404002_arXiv.txt": {
        "abstract": "Two new facilities have recently been proposed to measure low energy neutrino-nucleus cross sections, the $\\nu$-SNS (Spallation Neutron Source) and low energy beta beams. The former produces neutrinos by pion decay at rest, while the latter produces neutrinos from the beta decays of accelerated ions.  One of the uses of neutrino-nucleus cross section measurements is for supernova studies, where typical neutrino energies are 10s of MeV.  In this energy range there are many different components to the nuclear response and this makes the theoretical interpretation of the results of such an experiment complex. Although even one measurement on a heavy nucleus such as lead is much anticipated, more than one data set would be still better. We suggest that this can be done by breaking the electron spectrum down into the parts produced in coincidence with one or two neutrons, running a beta beam at more than one energy, comparing the spectra produced with pions and a beta beam or any combination of these. ",
        "introduction": "Neutrino-nucleus cross section measurements are desirable from the point of view of understanding nuclear structure, but they are perhaps even more desirable for astrophysical reasons.  The supernova is the best studied astrophysical environment where neutrino scattering reactions have significant impact. Proper inclusion of the reverse process, electron capture on nuclei, has recently been shown to factor significantly in the prospects for obtaining a supernova explosion \\cite{Hix:2003fg}. Furthermore, neutrino-nucleus interactions figure heavily in determining the nucleosynthesis that is produced during the course of a supernova explosion. Neutrino nucleosynthesis, which occurs when neutrinos spall neutrons and protons off of pre-existing nuclei, is driven entirely by neutrino-nucleus interactions \\cite{Woosley:1989bd,Heger:2003mm}.  Several papers have suggested that the r-process of nucleosynthesis may be impacted heavily by neutrino-nucleus interactions, e.g. \\cite{McLaughlin:1996eq,Haxton:1996ms}. In fact these reactions may have such a detrimental effect that they are an effective tool in constraining the environment \\cite{Meyer:1998sn}. Thirdly, neutrino-nucleus measurements are needed to calibrate supernova neutrino detectors.  For a description of supernova neutrino detection using lead, see \\cite{Boyd:em}.  A recent review of different techniques used to calculate neutrino-nucleus cross section measurements is given in \\cite{Kolbe:2003ys}.  Traditional neutrino beams are created using pions which produce both muon neutrinos and muon antineutrinos, and either electron neutrinos or electron antineutrinos. This is the case for the proposed $\\nu$-SNS which will produce neutrinos from pions decaying at rest by way of the Spallation Neutron Source at Oakridge National Laboratory \\cite{nuSNS}.  This facility will make improvements on existing measurements of nuclei such as carbon and iron \\cite{Krakauer:1991rf,Armbruster:vd,Berge:1989hr} and measure cross sections on new nuclei such as lead.  Newly proposed beta beam facilities \\cite{Zucchelli:sa} produce either electron neutrinos or antineutrinos from beta-plus or beta-minus decays of radioactive ions. Feasibility studies for beta beams are underway and a design is discussed in \\cite{Lindroos:2003kp,Bouchez:2003fy}. Beta beams were originally proposed as a way to make high energy neutrinos for use in long-baseline studies to determine the third, and as yet unknown mixing angle in the neutrino mixing matrix, $\\theta_{13}$ and to investigate CP-violation in the lepton sector \\cite{Mezzetto:2003ub,Bouchez:2003fy}.  However, lower energy beta beams have been proposed by Volpe \\cite{Volpe:2003fi} and an application for neutrino magnetic moment measurements has been discussed in \\cite{McLaughlin:2003yg}.  In this paper we consider neutrino-lead measurements using spectra that would be produced by the $\\nu$-SNS and from a beta beam facility. We consider a target mass of 10 tons, 20 meters away from the pion source and 10 meters away from the end of a straight section of a beta beam ring. We examine the spectra of the electrons that would be produced from charged current interactions.  We explore the signals which can be produced if the electrons can be identified as being created in coincidence with zero, one or two spalled neutrons. We also suggest the possibility that low energy beta beams be operated at more than one energy, therefore producing different neutrino energy spectra. ",
        "conclusions": "We have discussed the electron spectrum produced from a lead target for two different sources of neutrinos. Any future measurements of neutrino-nucleus cross sections are much anticipated. However, if possible it would be desirable to have more than one electron spectrum with which to calibrate theory.  We have illustrated the relative nuclear contribution to the electron energy spectra for various beta beam energies and compared this with the pion source.  We have also considered the contributions of the one neutron and two neutron spectrum. One could maximize the infomation with which to compare theory calculations by separating the spectra into parts associated with zero, one and two neutrons, by running the beta beam at more than one energy such as at $\\gamma = 5$, $\\gamma=10$ and $\\gamma=15$, or by using the electron spectra produced by both sources."
    },
    "0404/astro-ph0404536_arXiv.txt": {
        "abstract": "For an observer in the Hubble flow (comoving frame) the last scattering surface (LSS) is well approximated by a two-sphere. If a nontrivial topology of space is detectable, then this sphere intersects some of its topological images, giving rise to circles-in-the-sky, i.e., pairs of matching circles of equal radii, centered at different points on the LSS sphere, with the same pattern of temperature variations. Motivated by the fact that our entire galaxy is not exactly in the Hubble flow, we study the geometric effects of our galaxy's peculiar motion on the circles-in-the-sky. We show that the shape of these circles-in-the-sky remains circular, as detected by a local observer with arbitrary peculiar velocity. Explicit expressions for the radius and center position of such an observed circle-in-the-sky, as well as for the angular displacement of points on the circle, are derived. In general, a circle is detected as a circle of different radius, displaced relative to its original position, and centered at a point which does not correspond to its detected center in the comoving frame. Further, there is an angular displacement of points on the circles. These effects all arise from aberration of cosmic microwave background radiation, exhausting the purely geometric effects due to the peculiar motion of our galaxy, and are independent of both the large scale curvature of space and the expansion of the universe, since aberration is a purely local phenomenon. For a Lorentz-boosted observer with the speed of our entire galaxy, the maximum (detectable) changes in the angular radius of a circle, its maximum center displacement, as well as the maximum angular distortion are shown all to be of order $\\beta=(v/c)$ radians. In particular, two back-to-back matching circles in a finite universe will have an upper bound of $2|\\beta| $ in the variation of either their radii, the angular position of their centers, or  the angular distribution of points. ",
        "introduction": "\\label{sec:intro} Whether the universe is spatially finite and what its size and shape may be are among the fundamental open problems that high precision modern cosmology seeks to resolve. These questions of topological nature have become particularly topical, given the wealth of increasingly accurate cosmological observations, especially the recent results from the Wilkinson Microwave Anisotropy Probe (WMAP) experiment~\\cite{WMAP}, which have heightened the interest in the possibility of a finite universe. Indeed, reported non-Gaussianity in cosmic microwave background (CMB) maps~\\cite{CMB+NonGauss}, the small power of large-angle fluctuations~\\cite{WMAP-Spergel-et-al}, and some features in the power spectrum~\\cite{CMB+NonGauss,WMAP-Spergel-et-al} are large-scale anomalies which have been suggested as potential indication of a finite universe~\\cite{Poincare} (for reviews see~\\cite{Revs}).  Given the current high quality and resolution of such maps, the most promising search for cosmic topology is based on pattern repetitions of these CMB anisotropies on the last scattering surface (LSS). If a nontrivial topology of space is detectable% \\footnote{The extent to which a nontrivial topology may or may not be detected has been discussed in references~\\cite{TopDetec}.}, then the last scattering sphere intersects some of its topological images giving rise to the so-called circles-in-the-sky. Thus, the CMB temperature anisotropy maps will have matched circles: pairs of equal radii circles (centered at different points on the LSS sphere) that have the same pattern of temperature variations~\\cite{CSS1998}. These matching circles will exist in CMB anisotropy maps of universes with any detectable nontrivial topology, regardless of its geometry. Therefore, pairs of matched circles may be hidden in CMB maps if the universe is finite, and to observationally probe nontrivial topology on the largest scale available, one needs a statistical approach to scan all-sky CMB maps in order to draw the correlated circles out of them. The circles-in-the-sky method, devised by N. Cornish, D. Spergel and G. Starkman~\\cite{CSS1998} to search for a possible nontrivial topology of the universe, looks for such matching circles through a correlation statistic for sign detection (a function whose peaks indicate matched circles).  As originally conceived, the circles--in--the--sky method did not take into account the role of our galaxy's peculiar motion. In a recent paper~\\cite{Levin2004}, however, this point has been considered, and it has been argued, in a simplified context (flat spacetime), that, for any observer moving with respect to the CMB, two effects will take place, namely the circles will be deformed into ovals, and these ovals will be displaced with regard to the corresponding circles in the comoving frame. These effects were estimated to be, respectively, of order $\\beta^2$ and $|\\beta|$.  We show here that, regardless of any background curvature or expansion, the shape of these circles--in--the--sky, as locally detected by an observer in motion relative to the comoving one, \\emph{remains} circular. We derive explicit expressions for the radius and center position of such an observed circle-in-the-sky, as well as a formula for the angular displacement of points on the circle. In general, a circle is detected as a circle of different radius, displaced relative to its original position, and centered at a point which does not correspond to its detected center in the comoving frame. Further, there is an angular displacement of points on the circles. These effects arise all from aberration of CMBR, and exhaust the purely geometric effects due to our galaxy's peculiar motion on circles-in-the-sky. We also estimate the maximum values of these effects considering the peculiar motion of our galaxy. In particular, for two back-to-back matching circles in a finite universe the upper bounds in the variation of either their radii, the angular position of their center or in the angular distribution of points are all of order $2|\\beta|\\simeq 2.46\\times 10^{-3}$ radians. Although these effects are still below WMAP's angular resolution, we show that they are relevant for future CMB missions like the Planck satellite. ",
        "conclusions": "\\label{ConRem}  Motivated by the fact that our entire galaxy is not in the Hubble flow, we have studied all geometric effects due to our galaxy's peculiar motion on the circles-in-the-sky. As a consequence of Peronse's result~\\cite{Penrose1959} we have shown that circles-in-the-sky as detected by a comoving observer $O$ will also be detected as circles-in-the sky by any observer $O'$ locally coinciding with $O$ at the time they measure the CMBR. In general, a circle for $O$ will be detected by $O'$ as a circle of different radius, displaced relative to its original position, and centered at a point which does not correspond to its detected center in the comoving frame.  We have derived closed explicit expressions for the radius, $\\rho'$, the center position, $(\\theta_c', \\phi_c')$, and the angular displacement $\\alpha'$ of points on the circles as detected by the Lorentz-boosted observer $O'$. We have also plotted figures to illustrate several instances of these effects.  The maximum displacement in either radius or center or angular position of points on a circle is $\\pm \\beta$. From the value $\\beta=1.23~\\times~10^{-3}$  obtained from the dipole amplitude in the CMB spectrum, we have that the maximum displacement for each of these effects individually is $\\simeq \\pm 0.07^{\\circ}=4.2^{\\prime}$. Thus, two antipodal circles whose center vector positions  are parallel (antiparallel) to $\\vec{\\beta}$ will have a difference in radii of $0.14^{\\circ}$, as detected by $O$ and $O'$. This is below the current WMAP resolution (at best $0.25^{\\circ}= 15'$)~\\cite{WMAP}, but close to Planck's (at least $0.16^{\\circ}=10^{\\prime}$)~\\cite{Planck} and other more accurate forthcoming missions' resolution.  In Ref.~\\cite{Levin2004}, in a simplified flat spacetime, the spatial positions, relative to a ``moving'' frame, of the events defined by the intersection of the copies of the LSS spheres were determined; of course, these events are simultaneous in the CMB frame and, \\emph{ipso facto}, are non-simultaneous in the ``moving'' one. The specific procedure of collecting the spatial positions of these \\emph{non-simultaneous} events defines, in the ``moving'' frame, an oval figure, whose diameter along the boost direction is \\emph{Lorentz-dilated}, whereas its perpendicular diameter remains unchanged. Briefly, it was shown that a (spatial) circle in the CMB frame, as a geometric figure, is formally transformed into a (spatial) oval in the ''moving'' frame of the underlying flat spacetime. The approach we adopted to the effects of our galaxy's peculiar motion on the circles-in-the-sky is realistic and complete, since it does not rely on such a flat spacetime, and takes into account that a typical observation of the CMBR records simultaneously incoming light rays in an essentially infinitesimal detector used by a local observer, thus implying their projection onto the sky-sphere.  We note that the aberration of CMBR might be of a more general interest in the analysis of maps of CMB temperature anisotropies. Indeed, consider the celestial sphere pixelized by a comoving observer $O$ and consider another observer $O'$ coinciding with $O$ and with relative  velocity $\\vec{\\beta}$ along the positive $z$ direction. The coordinates $(\\theta_i, \\varphi_i)$ of the $i$--th pixel in the comoving frame are transformed to the moving frame according to~(\\ref{AberEsf}), which, at first order in $\\beta$, simply reads \\begin{equation} \\label{AberEsfNR} \\theta^\\prime_i = \\theta_i - \\beta \\sin \\theta_i \\qquad ; \\qquad \\varphi^\\prime_i = \\varphi_i \\; . \\end{equation} This displacement of the centers of the pixels gives rise to a distortion on the temperature variation pattern in CMB maps, which might be relevant for future missions.  Finally, we emphasize that all these effects arise from aberration of the CMBR, and exhaust the purely geometric effects due to the motion of our galaxy (or CMBR detectors). It is again worth noting that, since aberration of light is a purely local phenomenon, the effects considered in this work depend neither on the background curvature of space nor on the expansion of the universe.  \\bigskip \\noindent{\\bf Acknowledgments} \\medskip  We thank CNPq and FAPESP (contract 02/12328-6) for the grants under which this work was carried out."
    },
    "0404/astro-ph0404193_arXiv.txt": {
        "abstract": "We present    \\xmm and \\cha  observations of  \\oo , being the first to show evidence for a significant variation in the X-ray luminosity of  this Anomalous X--ray Pulsar (AXP). While during the first \\xmm (2000 December) and  \\cha  (2001 July) observations the source had a flux consistent with that measured on previous occasions ($\\sim$5$\\times$10$^{-12}$ erg cm$^{-2}$ s$^{-1}$), two more recent  observations found it at a considerably higher flux level of 2$\\times$10$^{-11}$ erg cm$^{-2}$ s$^{-1}$  (2002 August; \\cha) and 10$^{-11}$ erg cm$^{-2}$ s$^{-1}$ (2003 June; \\xmm). All the spectra are fit by the sum of a blackbody with kT$\\sim$0.6 keV and a power law with photon index $\\sim$3. No significant changes were seen in the spectral parameters, while the pulsed fraction in the 0.6-10 keV energy range decreased from $\\sim$90\\% in 2000 to $\\sim$53\\% in 2003. The spectral invariance does not support the presence of two physically distinct components in the AXP emission. The sparse coverage of the data does not permit us to unambiguously relate the observed variations to the two bursts seen from this source in the fall of 2001. ",
        "introduction": "The Anomalous X-ray Pulsars (AXPs) are a small group of  pulsars with a rotational period of a few seconds, a fairly stable spin-down and a very soft X--ray spectrum (Mereghetti \\& Stella 1995; van Paradijs, Taam \\& van den Heuvel 1995). If they are neutron stars, as their short spin period suggests, the loss  of rotational energy inferred from the observed period and period derivative values is too small to power their luminosity of 10$^{34}$-10$^{36}$ erg s$^{-1}$. AXPs are peculiar because of the strong evidence, accumulated over several years of intense observational effort, that they lack a main-sequence or giant mass-donor companion star (see Mereghetti et al. 2002 for a review). This has led to two classes of models involving isolated neutron stars. In the first class, the emission originates from accretion of material supplied by  a residual disk (e.g., Ghosh, Angelini \\& White 1997; Chatterjee, Hernquist \\& Narayan  2000; Alpar 2001), while  in the ``\\textit{Magnetar}'' model the  energy source is the decay of an extremely high (10$^{14}$--10$^{15}$ G) magnetic field (Duncan \\& Thompson 1992; Thompson \\& Duncan 1995, 1996). The latter model can explain the properties of the Soft Gamma-ray Repeaters (SGRs, see Hurley 2000 for a review), a class of bursting hard X--ray sources with quiescent soft X--ray emission quite similar to that of the AXPs. The magnetar interpretation has been supported by the recent discovery of short bursts in two AXPs, \\ee (Kaspi et al. 2003) and \\oo (Gavriil, Kaspi \\& Woods 2002).    \\oo was serendipitously diecovered with the \\textit{Einstein Observatory} as a 6.4 s pulsar near  the Carina Nebula (Seward, Charles \\& Smale 1986) and is one of the best studied AXPs. Early observations with EXOSAT and GINGA indicated an unusually soft (compared to accreting binary pulsars) power law spectrum and measured   a spin-down at a rate of $\\sim$1.5$\\times$10$^{-11}$ s s$^{-1}$ (Seward et al. 1986; Corbet \\& Day 1990). Higher quality spectra, not affected by contamination from bright nearby sources, could be obtained with the imaging instruments on board \\textit{BeppoSAX} (Oosterbroek et al. 1998), \\textit{ASCA} (Paul et al. 2000),  and \\xmm (Tiengo et al. 2002). These data showed a two component spectrum, composed of a blackbody with temperature kT$\\sim$0.64 keV plus a power law with photon index $\\alpha_{ph}\\sim$2-3, and a 2-10 keV luminosity of (1--2)$\\times10^{34}$ erg s$^{-1}$  (for an assumed distance d=5 kpc). The source lies at low galactic latitude, resulting in a significant absorption in soft X-rays (N$_H\\sim$10$^{22}$ cm$^{-2}$).  Deviations from a constant spin-down rate were first observed with ROSAT (Mereghetti 1995). Long term monitoring with \\xte carried out since 1996 (Kaspi et al. 2001) has shown that, compared to other AXPs for which phase coherent timing  could be obtained over extended time intervals, \\oo is characterized by a relatively high level of timing noise. The \\xte observations were also used to search for long term flux variability. Unfortunately, the presence of other X--ray sources in the field of view\\footnote{in particular the highly variable source $\\eta$ Carinae lies at $\\sim$40 arcmin} and the uncertainties in the subtraction of the time variable background, permitted us to use these data only to measure the pulsed component of the flux. Kaspi et al. (2001) concluded that the pulsed flux did not vary by more than 25\\% in 1996--2000.\\footnote{As noted by the same authors, this does not exclude flux variations, provided that the pulsed fraction was anticorrelated with the phase-averaged flux in order to cancel out variations in the pulsed flux.}  On the other hand, \\xte detected two bursts of short duration (about 51 and 2 s respectively)   from the direction of \\oo on 2001 October 29 and November 14 (Gavriil et al 2002). Although the possibility that they originated from a different source within the instrument field of view cannot be completely excluded, it is likely that they were due to \\oo, especially in light of the detection of short bursting activity also in another AXP (\\ee, Kaspi et al. 2003).   Here we present  \\xmm and \\cha observations showing  significant long-term variability in the X--ray flux from \\oo . ",
        "conclusions": "The data reported here provide solid evidence for long term variations in the  flux and pulsed fraction of the AXP \\oo . Independently of the model adopted to explain the X--ray emission, the spectral invariance argues against the presence of two physically distinct spectral components.  The sparse coverage of the data does not permit us to unambiguously relate the observed variations to the two bursts seen in the fall of 2001. Furthermore, the differences with respect to the behavior of \\ee after the 2002 June outburst suggest that the flux variations in \\oo are not necessarily related to SGR-like activity.    Long term flux variations have been reported for the AXP candidate AX J1845.0--0300 (Vasisht et al. 2000) and for the latest discovered member of the AXP class,  XTE J1810$-$197 (Ibrahim et al. 2003, Gotthelf et al. 2004). The latter source was first seen in a bright state (F$\\sim6\\times10^{-11}$erg cm$^{-2}$ s$^{-1}$, 2-10 keV) at the beginning of 2003. Its flux then decreased by a factor $\\sim$3 in the first half of 2003. Archival data from ROSAT and ASCA, in which XTE J1810$-$197 was two orders of magnitude fainter, indicate that this source behaves like a transient. Interestingly, XTE J1810$-$197 is quite similar to \\oo also because of its spectral parameters, nearly sinusoidal pulse profile with  a large pulsed fraction, spin-down at $\\sim$10$^{-11}$ s s$^{-1}$ with considerable timing noise, and IR counterpart (Israel et al. 2004).  These observations suggest that luminosity variations in AXPs are more common than previously thought and not necessarily associated with the emission of energetic flares as in the classical SGRs.   This work has been partially supported by the Italian Space Agency. Based on observations with \\xmm, an ESA science mission with instruments and contributions directly funded by ESA member states and USA."
    },
    "0404/astro-ph0404470_arXiv.txt": {
        "abstract": "We have detected statistically significant correlations between the cosmic microwave background and two tracers of large-scale structure, the HEAO1 A2 full sky hard X-ray map \\citep{Boldt} and the NVSS 1.4 GHz, nearly full sky radio galaxy survey \\citep{Condon}.  The level of correlations in these maps is consistent with that predicted for the integrated Sachs-Wolfe (ISW) effect in the context of a $\\Lambda CDM$ cosmological model and, therefore, provides independent evidence for a cosmological constant.  A maximum likelihood fit to the amplitude of the ISW effect relative to the predicted value is $1.13 \\pm 0.35$ (statistical error only). ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404316_arXiv.txt": {
        "abstract": "I will discuss a few select aspects of the most common and best understood galactic-scale outflow -- starburst-driven superwinds, focusing on winds from nuclear starburst galaxies. I will show that modern observations, in particular in the soft and hard X-ray bands, complement and reinforce the existing paradigm of superwinds as flows collectively driven by multiple SNe. The properties of the diffuse X-ray emission from dwarf starburst galaxies, $L_{\\rm BOL} \\sim L_{\\star}$ starbursts in spiral galaxies, and ULIRGS, are all consistent with superwind activity. Where appropriate, I  contrast the physics of starburst-driven winds with poorly collimated winds from AGN, and discuss what we know of the role of LLAGN and Seyfert nuclei in starburst superwind galaxies. ",
        "introduction": "It has long been appreciated that outflows from galaxies can have a major effect on galaxy formation and evolution, influencing such basic galactic properties such as the mean metal abundance \\citep{larson74,garnett02} or even the survival of low mass galaxies \\citep{dekel86}. More recent discoveries have reinvigorated interest in the role of galactic outflows, e.g.: The presence of metals in the true inter-galactic medium at low and high redshifts \\citep{songaila97,tripp00}; that a sizable fraction of all metals ever created now reside outside galaxies \\citep{pagel02}; and direct observational evidence for ubiquitous outflows from, and 100 kpc-scale cavities around, the Lyman break galaxies \\citep{adelberger03}.   By the far the best-studied, best-understood, and arguably most common form of galactic outflow capable of polluting the IGM are superwinds  \\citep{ham90,dahlem97}. These are loosely-collimated multi-phase outflows from actively-star forming galaxies, \\ie starburst galaxies. Within the local universe starbursts account for $\\sim 25$\\% of all massive star formation (and hence metal production), and starburst activity becomes progressively more important at higher redshifts \\citep{heckman98}. As all local starburst galaxies appear to have superwinds \\citep{lehnert96a}, it is clear that starburst-driven winds are of major importance.   Over the last 5 years new observations of superwinds, in particular satellite-based observations in the EUV and soft and hard X-ray bands, have substantially added to our understanding of the physics behind superwinds. I will discuss how these new observations support the long-standing conceptual picture of how superwinds work, specifically (a) how superwinds are driven by the collective mechanical power of multiple supernovae (SNe) occurring within the disks (in particular the nuclei) of starburst galaxies, and (b) X-ray emission from superwinds. Where appropriate I will contrast superwinds with AGN winds (\\ie loosely collimated outflows, and not large-scale jets). ",
        "conclusions": ""
    },
    "0404/astro-ph0404120_arXiv.txt": {
        "abstract": "{We present a detailed analysis of the evolution of a simulated isolated disc galaxy. The simulation includes stars, gas, star formation and simple chemical yields. Stellar particles are split in two populations: the old one is present at the beginning of the simulation and is calibrated according to various ages and metallicities; the new population borns in the course of the simulation and inherits the metallicity of the gas particles.  The results have been calibrated in four wavebands with the spectro-photometric evolutionary model GISSEL2000 (Bruzual \\&\\ Charlot~\\cite{BC93}). Dust extinction has also been taken into account.  A rest-frame morphological and bidimensional photometric analysis has been performed on simulated images, with the same tools as for observations.  The effects of the stellar bar formation and the linked star formation episode on the global properties of the galaxy (mass and luminosity distribution, colours, isophotal radii) have been analysed. In particular, we have disentangled the effects of stellar evolution from dynamic evolution to explain the cause of the isophotal radii variations.  We show that the dynamic properties (e.g. mass) of the area enclosed by any isophotal radius depends on the waveband and on the level of star formation activity.  It is also shown that the bar isophotes remain thinner than mass isodensities a long time ($> 0.7$~Gyr) after the maximum of star formation rate. We show that bar ellipticity is very wavelength dependent as suggested by real observations.  Effects of dust extinction on photometric and morphological measurements are systematically quantified. For instance, it is shown that, when the star formation rate is maximum, no more than 20\\%\\ of the B band luminosity can escape from the bar region whereas, without dust extinction, bar B band luminosity accounts for 80\\%\\ of the total B band luminosity. Moreover, the extinction is not uniformly distributed inside the bar. ",
        "introduction": "Early N-body models of isolated disc galaxies, using only collisionless particles, have clarified how several morphological features such as bars, spiral arms (e.g. Hohl \\cite{H78}) and boxy bulges (Combes \\& Sanders \\cite{CS81}), appear in stellar discs and evolve. Soon afterwards, it was realised that the role of the gas, the dissipative component, cannot be neglected. Collisionless N-body codes were thus coupled with hydrodynamic codes, either for cosmological purposes (Katz \\& Gunn \\cite{KG91}), or for detailed studies of isolated (Friedli \\& Benz \\cite{FB93}) or merging galaxies (Mihos \\& Hernquist \\cite{MH96}). For isolated disc galaxies, this kind of codes has brought clear evidence that the presence of gas can deeply change the morphology on a short timescale (less than 1~Gyr).  For instance, the gas can be responsible for the bar within bar phenomenon (Friedli \\& Martinet \\cite{FM93}) or the dissolution of the bar and the formation of a bulge by secular evolution (Pfenniger \\& Norman \\cite{PN90}, Norman et al. \\cite{NSH96}).  Studies of the fueling mechanisms at work in central regions (e.g. Friedli \\& Benz \\cite{FB93}, Shlosman \\& Noguchi \\cite{SN93}) have been also studied with such hybrid codes.  However, stars and gas are not only bound together by gravitation, but also by the process of star formation and the energy and chemical feedback of supernov\\ae\\ and stellar winds. Since star formation is not well understood even in our own Galaxy, several recipes have been implemented to mimic star formation in hydro+stellar dynamic codes. These recipes have allowed to study for the first time the influence of bars on galactic abundance gradient (Friedli et al. \\cite{FBK94}), the impact of the stellar ultraviolet radiation on star formation (Gerritsen \\& Icke \\cite{GI97}), the nuclear activity (Heller \\& Shlosman \\cite{HS94}), the starbursts induced by mergers (e.g. Mihos \\& Hernquist \\cite{MH96}), the renewal of bars by gas accretion (Bournaud \\& Combes \\cite{BC02}), etc. They also allow to perform self-consistent simulations of the formation and evolution of galaxies (e.g. Katz \\cite{K92}, Steinmetz \\& Muller \\cite{SM95}).  Whatever the kind of code used for the study of secular evolution, the morphological analysis of the simulations focuses on the properties of the mass distribution. To compare in details these models with observations it is implicitly assumed that all particles have the same mass-to-light ratio taken to unity for convenience. However, the real galaxies are built of composite stellar populations, then of various mass-to-light ratios and metallicities. Thus, in order to make more straightforward comparisons of simulations with multi-wavelength observations, we have used stellar population synthesis models to photometrically calibrate our self-consistent simulations which include stars, gas and star formation. We also take into account the effect of dust extinction.  This approach has been recently used to study the formation and the subsequent evolution of elliptical (e.g. Kawata \\cite{K01}) or disc galaxies in a $\\Lambda$ cold dark matter scenario (e.g. Westera et al. \\cite{WSBG02}, Abadi et al. \\cite{ANSE03}).  Our own approach slightly differs from previous ones since we do not simulate the formation of a disc galaxy in a primeval Universe. Indeed, we restrict ourselves to make a detailed study of the short-term evolution of an isolated disc galaxy already formed. All the particles are thus used to simulate the evolution of the galaxy.  We describe in Sect.~2 the numerical model and our technique of photometric calibration. The global evolution (i.e. integrated properties) of the model is presented in Sect.~3, and the morphological evolution in Sect.~4. We summarise our findings in Sect.~5. ",
        "conclusions": "We have carried out a photometric calibration of a self-consistent N-body simulation of an isolated disc galaxy, including stars, gas and star formation. Dust extinction has been taken into account. We have simultaneously study the dynamic and photometric evolution in various wavebands (B, R, H and K). We are able to reproduce observational properties such as \\ml\\ and B$-$H colour indices. A morphological and photometric analysis have also been performed on simulated images in various wavebands, using the same tools as for real observations.  Our main results can be summarised as follows: \\begin{enumerate}  \\item a dynamically self-consistent SFR peak produces a large increase in total luminosity in all wavebands whereas the stellar mass of the new population is much lower than the underlying old stellar mass. However, luminosity peaks can be almost completely obscured by dust extinction, since the young stellar population initially seats in gas-rich regions where extinction is maximum. This effect is particularly important on B$-$H colour maps since regions of highest SFR could become unobservable. It could lead to underestimate SFR from photometric or colour measurements.  \\item we show how isophotal radii evolve with respect to mass redistribution during the dynamic evolution but also with respect to stellar evolution. Hence, the dynamic properties of the area enclosed by any isophotal radius depends on the waveband and on the star formation activity. Extinction effects do not affect the determination of isophotal radii, even in B band. It is also noteworthy that the luminosity integrated inside the isophotal radius at 20.5~\\magunit\\ in the H band is not a good tracer of mass.  \\item the luminosity integrated into the bar length reaches $\\approx 80$\\%\\ of the total B band luminosity at the SFR maximum in absence of dust, and $\\ga 60$\\% when SFR is low, whereas the bar length remains smaller than or equal to half the isophotal radius.  \\item when SFR is high ($\\approx 30$~\\sfrunit) less than 20\\%\\ of the B band luminosity of the bar region succeed to escape in the presence of dust because most star formation occurs in gas-rich regions where extinction is the most efficient.  \\item as long as star formation is active, mass isodensities of the bar region are rounder than isophotes. Thus, the surface brightness distribution, even in H band, is a good tracer of mass only when SFR is below 1~\\sfrunit. Mass modelling that used surface brightness distribution might give biased results if star formation is active, even at a moderate rate.  \\end{enumerate}  The use of other morphological classification tools (e.g. concentration and asymmetry, Abraham et al. \\cite{Aetal96}) will be will be reported in a forthcoming paper."
    },
    "0404/astro-ph0404299_arXiv.txt": {
        "abstract": "We discuss observational constrains coming from supernovae Ia imposed on the behaviour of the Randall-Sundrum models. In the case of dust matter on the brane, the difference between the best-fit Perlmutter model with a $\\Lambda$-term and the best-fit brane models becomes detectable for redshifts $z > 1.2$. It is interesting that brane models predict brighter galaxies for such redshifts which is in agreement with the measurement of the $z = 1.7$ supernova. We also demonstrate that the fit to supernovae data can also be obtained, if we admit the \"super-negative\" dark energy (phantom matter) $p = - (4/3) \\varrho$ on the brane, where the dark energy in a way mimics the influence of the cosmological constant. It also appears that the dark energy enlarges the age of the universe which is demanded in cosmology. Finally, we propose to check for dark radiation and brane tension by the application of the angular diameter of galaxies minimum value test. We point out the existence of coincidence problem for the brane tension parameter. ",
        "introduction": "In recent several years a lot of effort has been done on the idea that our Universe is a boundary of a higher-dimensional spase time manifold \\cite{Arkani-Hamed98,Arkani-Hamed99}. Kaluza and Klein first discussed $5-dimensional$ space time to unify a gravity and electromagnetism. Among supersting theories which may unify all interactions M-theory is a strong candidate for the description of real world. In this theory, gravity is a truly higher-dimensional theory, becoming effectively 4-dimensional at lower energies. The standard model matter fields are confined to the 3-brane while gravity can, by its universal character, propagate in all extra dimensions. In the brane world models inspired by string/M theory (\\cite{rs1},\\cite{rs2}, \\cite{hw}) new two parameters which doesn't present in standard cosmology are introduced, namely brane tension $\\lambda$ and dark radiation $U$. One of new approaches was proposed by Randall and Sundrum \\cite{rs1}, \\cite{rs2} where our Minkowski brane is localized in 5dimensional anti-de Sitter space time with metric: \\begin{equation} \\label{eq:1} ds^2=\\exp(-2|y|l)(-dt^2+d\\vec{x}^2)+dy^2. \\end{equation}   For $y \\ne 0$, this metric satisfies the 5dimensional Einstein equation with the negative cosmological constant $\\tilde{\\check{\\Lambda_{(5)}}} \\propto -l^{-2}$. The brane is located at $y=0$ and the induced metric on brane is a Minkowski metric. The bulk is a 5-dimensional anti-deSitter metric with $y=0$ as a boundary.  We should mention that before the Randall and Sundrum work \\cite{rs2} where they proposed a mechanism to solve the hierarchy problem by a small extra dimension, large extra dimensions were proposed to solve this problem by Arkani-Hamed et.al. \\cite{Arkani-Hamed98,Arkani-Hamed99}. This gives an interesting feature because TeV gravity might be realistic and quantum gravity effects could be observed by a next generation particle collider. The Newtonian gravity potential on the brane is recovered at lowest order $V(r)=\\frac{GM}{r}(1+\\frac{2l^2}{3r^2})$. In this paper we demonstrate that if the brane world is the Randall-Sundrum version is realistic we may find some evidence of higher dimensions.   In  \\cite{Szydlo02} we gave the formalism to express dynamical equations in terms of dimensionless observational density parameters $\\Omega$. In this notation (see also \\cite{Dabrowski96,AJI+II,AJIII}) the Friedmann equation for brane universes takes the form \\begin{equation} \\label{eq:FriedCCC} \\frac{1}{a^2} \\left( \\frac{da}{dt} \\right)^2 = \\frac{C_{\\gamma}}{a^{3\\gamma}} + \\frac{C_{\\lambda}}{a^{6\\gamma}} - \\frac{k}{a^2} + \\frac{\\Lambda_{(4)}}{3} + \\frac{C_{\\cal U}}{a^4} , \\end{equation} where $a(t)$ is the scale factor, $k=0,\\pm1$ the curvature index, here we use natural system of units in which $8\\pi G=c=1$, $\\Lambda_{(4)}$ is the 4-dimensional cosmological contant, and $\\gamma$ the barotropic index ($p = (\\gamma - 1)\\varrho$, $p$ - the pressure, $\\varrho$ - the energy density), the constants $C_{\\lambda} = 1/6\\lambda \\cdot a^{6\\gamma} \\varrho^2$ and $C_{\\cal U} = 2/ \\lambda \\cdot a^4 {\\cal U}$, $C_{\\lambda}$ comes as a contribution from brane tension $\\lambda$, and $C_{\\cal U}$ as a contribution from dark radiation.  Because $\\rho^2$ term and dark radiation term do not appear in the standard cosmology, such terms could provide a smal window to see the extra dimensions.  In order to study observational tests we now define dimensionless observational density parameters \\bea \\label{Omegadef} \\Omega_{\\gamma}  &=&  \\frac{1}{3H^2} \\varrho , \\hspace{15pt} \\Omega_{\\lambda}  =  \\frac{1}{6H^2\\lambda} \\varrho^2 , \\hspace{15pt} \\Omega_{\\cal U}  =  \\frac{2}{H^2\\lambda} {\\cal U} ,\\nonumber \\\\ \\Omega_{k}  &=&  - \\frac{k}{H^2a^2} , \\hspace{15pt} \\Omega_{\\Lambda_{(4)}}  =  \\frac{\\Lambda_{(4)}}{3H^2} , \\eea where the Hubble parameter $H = \\dot{a}/a$, and the deceleration parameter $q  =  - \\ddot{a}a/\\dot{a}^2$ , so that the Friedmann equation (\\ref{eq:FriedCCC}) can be written down in the form \\begin{equation} \\label{Om=1} \\Omega_{\\gamma} + \\Omega_{\\lambda} + \\Omega_{k} + \\Omega_{\\Lambda_{(4)}} + \\Omega_{\\cal U} = 1  . \\end{equation} Note that $\\Omega_{\\cal U}$ in (\\ref{Omegadef}), despite standard radiation term, can either be positive or negative.  It is useful to rewrite (\\ref{eq:FriedCCC}) to the dimensionless form. Let us consider a standard Friedmann-Robertson-Walker universe (hereafter FRW) filled with mixture of matter with the equation of state $p_i = (\\gamma_i-1) \\rho_i$. Then we obtain the basic equation in the form:  \\begin{equation} \\label{eq:10} \\frac{\\dot{x}^{2}}{2} = \\frac{1}{2} \\Omega_{k,0} + \\frac{1}{2}\\sum_{i} \\Omega_{i,0} x^{2-3\\gamma_{i}} =-V(x)\\\\ \\end{equation} \\begin{equation} \\label{eq:10a} \\ddot{x} = - \\frac{1}{2} \\sum_{i} \\Omega_{i,0}(2-3\\gamma_{i}) x^{1-3\\gamma_{i}}=-\\frac{\\partial V(x)}{\\partial x} \\end{equation} where $i=(\\gamma,\\lambda,\\Lambda,U)$, and  \\begin{equation} \\label{eq:11} x \\equiv \\frac{a}{a_{0}}, \\qquad T \\equiv |H_{0}| t, \\qquad \\dot{}\\equiv \\frac{d}{dT}. \\end{equation} and $t$ is original cosmological time, $V$ is the potential function.  Therefore the dynamics of the considered model is equivalent to introducing fictitious fluids which mimic $\\rho^2$ contribution and dark energy term. For dark energy $\\gamma=4/3$ whereas for brane $\\gamma_{\\lambda}= 2\\gamma$. The presented formalism is useful in analysis of observational tests of brane models.   Above relations allow to write down an explicit redshift-magnitude relation  for the brane models to study their compatibility with astronomical data which is the subject of the present paper. Obviously, the luminosity of galaxies depends on the present densities of the different components of matter content $\\Omega_{i}$ given by (\\ref{Omegadef}) and their equations of state, reflected by the value of the barotropic index $\\gamma_i$. ",
        "conclusions": "We shown that there exists an effective method of constraining exotic physics coming from superstrings M theory from observation of distant supernovae. We obtain the estimated value the density parameters $\\Omega_{\\lambda,0}$ and $\\Omega_{\\Lambda,0}$ .  Finally, as a result we also obtain that at high redshifts the expected luminosity of supernovae Ia should be brighter then in the Perlmutter model. For the best fit value we obtain $\\Omega_{\\lambda,0} \\simeq 0.01$ which seems to be unrealistic. It is because if we consider pure Randall-Sundrum models, then there is a constraint on the parameter $\\Omega_{\\lambda,0}$ from the requirement of not violating the four-dimensional gravity on sufficiently large spatial scale. This constraint is that value of $\\lambda$ is to be no less than about $(100 \\, GeV)^4$, which means that during the late epoch the $\\rho^2$ term in the model should be small.   So, the obtained value of $\\Omega_{\\lambda,0} \\sim 0.01$ is the observational limit which is not based on theoretical model assumptions. The density $\\rho_{m,0}$ at the time relevant for supernovae measurements is about $(10^{-3}eV)^4$. Thus, the size of the parameter $\\Omega_{\\lambda,0}$ is on purely theoretical grounds, at most $10^{-56}$. The fits discussed in the paper innvolve $\\Omega_{\\lambda,0}$ of order $0.01$. Therefore similarly to the cosmological constant problem there is a coincidence problem with brane tension $\\lambda$ (it is treated as a constant) namely: why don't we see the large brane tension expected from the Randall-Sundrum theory which is about $10^{54}$ times larger than the value predicted by the Friedmann equation which fit SNIa data. A phenomenological solution to this problem seems to be dynamically decaying $\\lambda$."
    },
    "0404/astro-ph0404250_arXiv.txt": {
        "abstract": "{Deep, wide--field, continuum--subtracted, images in the light of the \\ha\\ + \\NII\\ and \\oiii\\ nebular emission lines have been obtained of the environment of the Luminous Blue Variable (LBV) star P~Cygni. A previously discovered, receding, nebulous filament along PA~50\\deg\\ has now been shown to extend up to 12\\arcmin\\ from this star. Furthermore, in the light of \\oiii, a southern counterpart is discovered as well as irregular filaments on the opposite side of P~Cygni.  Line profiles from this nebulous complex indicate that this extended nebulosity is similar to that associated with middle--aged supernova remnants. However, there are several indications that it has originated in P~Cygni and is not just a chance superposition along the same sight--line. This possibility is explored here and comparison is made with a new image of the LBV star R~143 in the LMC from which similar filaments appear to project.  The dynamical age of the P~Cygni giant lobe of $\\approx$ 5$\\times$10$^{4}$ yr is consistent with both the predicted and observed durations of the LBV phases of 50\\msun\\ stars after they have left the main sequence. Its irregular shape may have been determined by the cavity formed in the ambient gas by the energetic wind of the star, and shaped by a dense torus, when on the main sequence.  The proper motion and radial velocity of P~Cygni, with respect to its local environment, could explain the observed angular and kinematical shifts of the star compared with the giant lobe. ",
        "introduction": "The circumstellar environment of the proto-typical Luminous Blue Variable star (LBV - Conti 1984; Humphreys 1989; Davidson, Moffat \\& Lamers 1989) P~Cygni has been revealed at optical wavelengths in the work presented in six recent papers (Johnson et al. 1992; Barlow et al 1994; Meaburn et al 1996; O'Connor, Meaburn \\& Bryce 1998; Meaburn, L\\'{o}pez \\& O'Connor 1999; Meaburn et al 2000). Two nearly spherical, but complex, circumstellar shells were discovered. The bright \\NII\\ and \\NiII\\ emitting, 22\\arcsec\\ diameter inner shell (IS) was found (Barlow et al 1994) to be surrounded by a fainter \\NII\\ emitting, 1.6\\arcmin\\ diameter, outer shell (OS). The dynamical ages of the IS (Barlow et al 1994) and the OS (Meaburn et al 2000), for a distance to P~Cygni of 1.8~kpc (van Schewick 1968; Lamers, de Groot \\& Cassatella 1983) were derived from their expansion velocities as 880 and 2400 yr respectively.  This would place the creation of both of these shells as well before the great outburst of 1600 {\\sc ad} (de Groot 1969). Humphreys and Davidson (1994) emphasise that knowledge of P~Cygni's eruptive `geyser--like' behaviour prior to this date is unknown.  Potentially as interesting, is the presence of a filamentary, line emitting, giant lobe (GL), discovered by O'Connor, Meaburn \\& Bryce (1998) which could be a relic of the activity of P~Cygni close, or even prior, to its entry into its LBV phase. In the later work by Meaburn et al (2000) the northern ridge of GL had been traced for 7\\arcmin, along PA $\\approx$ 50\\deg, from P~Cygni and shown to connect morphologically and kinematically with the receding side of the  OS. However,  the previous observations in Meaburn et al (1999 \\& 2000) of the kinematical behaviour of this northern ridge of GL strengthen, but do not absolutely confirm, the suggestion that it is directly associated with P~Cygni (Meaburn et al 1999) and not a chance superposition along the same sight-line.  It is also significant that a morphologically similar feature (Meaburn 2001) to the P~Cygni GL has since been found (see Smith et al 1998 - their figure 2) to be apparently projecting from the LBV star R143 in the Large Magellanic Cloud.  The original observations by O'Connor et al (1998) of the P~Cygni GL, revealing its north eastern filamentary arc, were made in the light of the \\nii\\ nebular emission line. In the present paper these have  been supplemented by deep, wide--field, continuum--subtracted images in the light of \\ha\\ + \\NII\\ and more significantly \\oiii. Although many new, but connected, features are reported in the present paper the designation GL will be used to describe the whole of this extended nebular complex that has now been revealed around P~Cygni. Previously obtained line profiles will also be compared with these most recent images. Similarly, a new deep image of the candidate GL surrounding R~143 in the LMC will be presented and compared to the motions measured by Weis (2003). ",
        "conclusions": "The GL apparantly projecting from P~Cygni has now been shown to have a southern counterpart on the eastern side of P~Cygni. Also a more complex counterpart has been discovered to the west of the star.  The overall apparent extent of this GL is now found to be 9 pc.  These structures emit the \\oiii\\ line strongly enhancing their contrast against the confusing emission from the ambient ionized gas along the same sight--lines.  These newly discovered structures are now shown to have been detected kinematically in previous spectral observtions. Although receding radial velocities dominate, kinematical association with P~Cygni is strengthened, for the newly discovered southern ridge of the GL, on the eastern side of P~Cygni, has  radial velocities on either side of \\vsys\\ of this star.  It is proposed that the P~Cygni GL phenomenon was formed by continual activity between the age of the OS (2400 yr) and the dynamical age ($\\approx$ 5~x~10$^{4}$ yr) of the extreme extent of the GL. Within this interpretation sporadic LBV eruptions over this extended period appear to have ocurred.  The irregular shape of the GL could then be a consequence of the shape of the cavity formed by the wind of the 50~\\msun\\ star when on the main sequence immediately prior to its LBV phase.  Similar GL--like features, with an apparent extent of 5 pc, are shown to project  from the LMC, LBV star R~143. Their presence strenghtens the possibility that these and their P~Cygni equivelents are associated intimately with these stars and not just chance alignments of un-related supernova filaments some of whose characteristics they share."
    },
    "0404/astro-ph0404585_arXiv.txt": {
        "abstract": "We consider the consequences for the relic neutrino abundance if extra neutrino interactions are allowed, e.g., the coupling of neutrinos to a light (compared to $m_\\nu$) boson.  For a wide range of couplings not excluded by other considerations, the relic neutrinos would annihilate to bosons at late times, and thus make a negligible contribution to the matter density today. This mechanism evades the neutrino mass limits arising from large scale structure. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404066_arXiv.txt": {
        "abstract": "Using a sample of 91 galaxies distributed over 27 Compact Groups of Galaxies (CGs), we define an index that allows us to quantify their level of activity, be it AGN or star formation. By combining the mean activity index with the mean morphological type of the galaxies in a group we are able to quantify the evolutionary state of the groups. We find that they span a sequence in evolution, which is correlated with the spatial configuration of the galaxies making up a CG. We distinguish three main configuration Types, A, B and C. Type~A CGs show predominantly low velocity dispersions and are rich in late-type spirals that are active in terms of star formation or harbor an AGN. Type~B groups have intermediate velocity dispersions and contain a large fraction of interacting or merging galaxies. Type~C is formed by CGs with high velocity dispersions, which are dominated by elliptical galaxies that show no activity. We suggest that the level of evolution increases in the sense A$\\Rightarrow$B$\\Rightarrow$C. Mapping the groups with different evolution levels in a diagram of radius versus velocity dispersion does not reveal the pattern expected based on the conventional fast merger model for CGs, which predicts a direct relation between these two parameters. Instead, we observe a trend that goes contrary to expectation: the level of evolution of a group increases with velocity dispersion. This trend seems to be related to the masses of the structures in which CGs are embedded. In general, the level of evolution of a group increases with the mass of the structure. This suggests either that galaxies evolve more rapidly in massive structures or that the formation of CGs embedded in massive structures predated the formation of CGs associated with lower mass systems. Our observations are consistent with the formation of structures as predicted by the CDM model (or $\\Lambda$CDM), assuming the formation of galaxies is a biased process. ",
        "introduction": "Although it seems today an inescapable conclusion that the formation and evolution of galaxies is influenced by their environment, the details of how these processes occur in space and time are still largely unknown. One example is that of compact groups of galaxies (CGs). As a result of our studies of the activity in galaxies in CGs, we now have a better understanding of the evolution of galaxies in these systems (Coziol et al. 1998a,b; Coziol, Iovino \\& de Carvalho 2000). Our observations showed emission-line galaxies to be remarkably frequent, representing more than 50\\% of the galaxies in Compact Groups. Non-thermal activity, in the form of Seyfert~2s, LINERs and numerous Low Luminosity AGNs (LLAGNs; see Coziol et al. 1998a for a definition of this activity type in CGs), was found to constitute one important aspect of this activity, whereas nuclear star formation, although mildly enhanced in some groups, was noted to be generally declining. These observations were considered to be consistent with the effects of tidal forces exerted on disk galaxies when they fall into the potential well of a rich cluster or group of galaxies (Coziol, Iovino \\& de Carvalho 2000). Tidal stripping will remove gas from a galaxy, reducing star formation in the disk, whereas tidal triggering will start a short burst of star formation in the nuclear region and fuel an AGN (Merritt 1983, 1984; Byrd \\& Valtonen 1990; Henrikson \\& Byrd 1996; Fujita 1998).  The above processes may also produce the density-morphology-activity relation observed in CGs (Coziol et al. 1998a). By losing their gas and forming new stars near their nucleus, increasing their bulge, the morphology of spiral galaxies falling into groups is transformed to an earlier type. Assuming that groups are continuously replenished by spiral galaxies from the field (Governato, Tozzi \\& Cavaliere 1996; Coziol, Iovino \\& de Carvalho 2000), the cores of the groups are naturally expected to be populated by quiescent galaxies and AGNs (LLAGNs included), all having an early-type morphology, and their periphery to be richer in late-type star forming galaxies.  What is missing in the above description is a connection between galaxy evolution and the physical processes responsible for the formation and evolution of CGs. Our first interpretation of these systems, based on galaxy-galaxy interaction models, suggested they could not survive mergers over a long period of time (Barnes~ 1989), which seems in conflict with the high number of CGs observed today (for a different point of view, see Aceves \\& Vel\\'azquez 2002). It was then realized that this paradox may be explained, in part, by the fairly simplistic assumptions made about the nature of these systems. For example, if CGs are associated with larger and dynamically more complex structures, as many redshift surveys suggest (Rood \\& Strubble 1994; Ramella et al. 1994; Garcia 1995; Ribeiro et al. 1998; Barton, de Carvalho \\& Geller 1998), their formation and evolution must also be more complicated than previously thought.  In this study, we examine further the question of the formation and evolution of CGs, by extending our analysis of the activity to a larger sample of Hickson Compact Groups of galaxies (HCGs: Hickson 1982). ",
        "conclusions": "We have defined a new spectroscopic index that allows us to quantify the level of activity (AGN and star formation) in galaxies. By taking the mean of this index in a CG and combining it with the mean morphological type of the galaxies it proves possible to quantify the evolutionary state of these systems.  Applying our method to a sample of 27 CGs from Hickson's catalog, we have found an evolutionary sequence, which is correlated with the projected spatial configuration of the groups. Mapping the position of CGs with different levels of evolution and group morphology in a diagram of radius versus velocity dispersion, we did not observe the pattern predicted by the conventional fast merger model. Contrary to expectation, we found that the level of evolution of CGs increases with velocity dispersion. This trend was further shown to be possibly connected to the masses of the CGs or to the structures in which they are embedded. Assuming that the velocity dispersion increases with the mass, the trend we observe would thus imply that the most evolved groups are found within the most massive structures.  We propose two different hypotheses to explain our results. The first assumes that the evolution of galaxies accelerates with the number of interactions in massive structures. The other assumes that the formation of CGs follows the formation of the large scale structure, and that massive structures develop before less massive ones. Our observations are consistent with the formation of structures as predicted by the CDM model (or $\\Lambda$CDM), assuming the formation of galaxies is a biased process: galaxies developing first in high-density structures."
    },
    "0404/astro-ph0404558_arXiv.txt": {
        "abstract": "{Ionized gas and stellar kinematics have been measured along the major axes of seventeen nearby spiral galaxies of intermediate to late morphological type. We discuss the properties of each sample galaxy distinguishing between those characterized by regular or peculiar kinematics. In most of the observed galaxies ionized gas rotates more rapidly than stars and have a lower velocity dispersion, as  is to be expected if the gas is confined in the disc and supported by rotation while the stars are mostly supported by dynamical pressure. In a few objects, gas and stars show almost the same rotational velocity and low velocity dispersion, suggesting that their motion is dominated by rotation. Incorporating the spiral galaxies studied by Bertola et al. (1996), Corsini et al.\\ (1999, 2003) and Vega Beltr\\'an et al.\\ (2001) we have compiled a sample of 50 S0/a--Scd galaxies, for which the major-axis kinematics of the ionized gas and stars have been obtained with the same spatial ($\\approx1''$) and spectral  ($\\approx 50$ \\kms) resolution, and measured with the same analysis techniques. This allowed us to address the frequency of counterrotation in spiral galaxies. It turns out that less than $12\\%$ and less than $8\\%$ (at the $95\\%$ confidence level) of the sample galaxies host a counterrotating gaseous and stellar disc, respectively. The comparison with S0 galaxies suggests that the retrograde acquisition of small amounts of external gas gives rise to counterrotating gaseous discs only in gas-poor S0s, while in gas-rich spirals the newly acquired gas is swept away by the pre-existing gas. Counterrotating gaseous and stellar discs in spirals are formed only from the retrograde acquisition of large amounts of gas exceeding that of pre-existing gas, and subsequent star formation, respectively. ",
        "introduction": "\\label{sec:introduction}  Studying the interplay between ionized gas and stellar kinematics allows us to address different topics concerning the dynamical structure of disc galaxies and to constrain the processes leading to their formation and evolution. These topics include the study of the mass distribution of luminous and dark matter (see Sofue \\& Rubin 2001 for a review), the ubiquity of supermassive black holes and their relationship with the large-scale properties of the host galaxies (see Merritt \\& Ferrarese 2001 for a review), the discovery of kinematically decoupled components (see Bertola \\& Corsini 1999 for a review), the origin of disc heating (Merrifield, Gerssen \\& Kuijken 2001 and references therein), and the presence of pressure-supported ionized gas in bulges (Bertola et al.\\ 1995; Cinzano et al.\\ 1999).  All these issues will greatly benefit from a survey devoted to the comparative measurements of ionized gas and stellar kinematics in S0s and spiral galaxies. Since these  are available only for a limited number of objects, in the past years we began a scientific programme aimed at deriving detailed velocity curves and velocity dispersion radial profiles of ionized gas and stars along the major axes of disc galaxies (Bertola et al.\\ 1995, 1996; Corsini et al.\\ 1999, 2003; Vega Beltr\\'an et al.\\ 2001; Funes et al.\\ 2002) to be used for mass modelling (Corsini et al.\\ 1999; Cinzano et al.\\ 1999; Pignatelli et al.\\ 2001).  This paper is organized as follows. An overview of the properties of the sample galaxies as well as the spectroscopic observations and data analysis are presented in Section \\ref{sec:observations}. The resulting ionized gas and stellar kinematics are given and interpreted in Section \\ref{sec:kinematics}. In particular, we derived the central velocity dispersion of stars of all the sample galaxies with the aim of studying the relationship between that the disc circular velocity and bulge velocity dispersion (Ferrarese 2002; Baes et al.\\ 2003) in a forthcoming paper (but see also Pizzella et al.\\ 2003). The fraction of spiral galaxies hosting a counterrotating component is estimated in Section \\ref{sec:counterrotation} by analysing the major-axis kinematics of all the spiral galaxies we have observed in recent years. Our conclusions are discussed in Section \\ref{sec:conclusions}.   \\begin{table*}[t] \\caption[]{Parameters of the sample galaxies} \\begin{center} \\begin{footnotesize} \\begin{tabular}{lllrrcrrrrrcc} \\hline \\noalign{\\smallskip} \\multicolumn{1}{c}{Object} & \\multicolumn{2}{c}{Type} & \\multicolumn{1}{c}{$B_T$} & \\multicolumn{1}{c}{PA} & \\multicolumn{1}{c}{$i$} & \\multicolumn{1}{c}{$V_{\\odot}$} & \\multicolumn{1}{c}{$D$} & \\multicolumn{1}{c}{Scale} & \\multicolumn{1}{c}{$R_{25}$} & \\multicolumn{1}{c}{$M_{B_T}^0$} & \\multicolumn{2}{c}{$R_{\\rm last}/R_{25}$} \\\\  \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{[stars]} & \\multicolumn{1}{c}{[gas]} \\\\  \\multicolumn{1}{c}{[name]} & \\multicolumn{1}{c}{[RSA]} & \\multicolumn{1}{c}{[RC3]} & \\multicolumn{1}{c}{[mag]} & \\multicolumn{1}{c}{[\\degr]} & \\multicolumn{1}{c}{[\\degr]} & \\multicolumn{1}{c}{[\\kms]} & \\multicolumn{1}{c}{[Mpc]} & \\multicolumn{1}{c}{[pc$/''$]} & \\multicolumn{1}{c}{[$''$]} & \\multicolumn{1}{c}{[mag]} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{} \\\\  \\multicolumn{1}{c}{$^{(1)}$} & \\multicolumn{1}{c}{$^{(2)}$} & \\multicolumn{1}{c}{$^{(3)}$} & \\multicolumn{1}{c}{$^{(4)}$} & \\multicolumn{1}{c}{$^{(5)}$} & \\multicolumn{1}{c}{$^{(6)}$} & \\multicolumn{1}{c}{$^{(7)}$} & \\multicolumn{1}{c}{$^{(8)}$} & \\multicolumn{1}{c}{$^{(9)}$} & \\multicolumn{1}{c}{$^{(10)}$} & \\multicolumn{1}{c}{$^{(11)}$} & \\multicolumn{1}{c}{$^{(12)}$} & \\multicolumn{1}{c}{$^{(13)}$} \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\object{NGC 210}  & Sb(rs)  & SABb(s)     & 11.60 & 165 & 49 & 1650 & 23.2 & 113 & 150 & $-20.5$ & 0.3&0.8 \\\\ \\object{NGC 615}  & Sb(r)   & SAb(rs)     & 12.47 & 155 & 67 & 1849 & 25.7 & 125 & 108 & $-20.3$ & 0.7&0.9 \\\\ \\object{NGC 1620} & ---     & SABbc(rs)   & 13.08 &  25 & 70 & 3509 & 46.1 & 224 &  86 & $-21.1$ & 0.7&1.2 \\\\ \\object{NGC 2590} & ---     & SAbc(s):    & 13.94 &  77 & 72 & 4960 & 63.4 & 308 &  67 & $-21.0$ & 0.5&1.0 \\\\ \\object{NGC 2708} & ---     & SABb(s)pec? & 12.80 &  25 & 60 & 1984 & 23.5 & 114 &  78 & $-19.4$ & 0.6&0.8 \\\\ \\object{NGC 2815} & Sb(s)   & SBb(r):     & 12.81 &  10 & 72 & 2535 & 30.0 & 145 & 104 & $-21.0$ & 0.4&1.0 \\\\ \\object{NGC 3054} & SBbc(s) & SABb(r)     & 12.35 & 118 & 52 & 2430 & 28.5 & 138 & 114 & $-20.3$ & 0.5&0.9 \\\\ \\object{NGC 3200} & Sb(r)   & SABc(rs):   & 12.83 & 169 & 73 & 3526 & 43.4 & 211 & 125 & $-21.5$ & 0.5&0.9 \\\\ \\object{NGC 3717} & Sb(s)   & SAb:sp      & 12.24 &  33 & 81 & 1748 & 19.6 &  95 & 180 & $-20.5$ & 0.4&0.5 \\\\ \\object{NGC 4682} & Sc(s)   & SABcd(s)    & 13.14 &  83 & 61 & 2344 & 28.8 & 140 &  77 & $-20.0$ & 0.3&0.8 \\\\ \\object{NGC 5530} & Sc(s)   & SAbc(rs)    & 11.79 & 127 & 62 & 1092 & 11.7 &  57 & 125 & $-19.4$ & 0.6&0.9 \\\\ \\object{NGC 6118} & Sc(s)   & SAcd(s)     & 12.42 &  58 & 65 & 1580 & 21.4 & 104 & 140 & $-20.4$ & 0.6&0.8 \\\\ \\object{NGC 6878} & Sc(r)   & SAb(s)      & 13.45 & 125 & 40 & 5821 & 77.3 & 375 &  48 & $-21.3$ & 0.6&1.0 \\\\ \\object{NGC 6925} & Sbc(r)  & SAbc(s)     & 12.07 &   5 & 75 & 2783 & 37.7 & 183 & 134 & $-21.9$ & 0.6&0.9 \\\\ \\object{NGC 7083} & Sb(s)   & SAbc(s)     & 11.87 &   5 & 53 & 3093 & 39.7 & 192 & 116 & $-21.5$ & 0.5&1.0 \\\\ \\object{NGC 7412} & Sc(rs)  & SBc(s)      & 11.88 &  65 & 42 & 1712 & 22.7 & 110 & 116 & $-20.1$ & 0.4&0.8 \\\\ \\object{NGC 7531} & Sbc(r)  & SABbc(r)    & 12.04 &  15 & 67 & 1591 & 20.9 & 102 & 134 & $-20.2$ & 0.4&0.6 \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\noalign{\\smallskip} \\noalign{\\smallskip} \\end{tabular} \\begin{minipage}{18cm} NOTES -- $^{(2)}$Morphological classification from Sandage \\& Tammann (1981, RSA hereafter). $^{(3)}$Morphological classification from RC3. $^{(4)}$Total observed blue magnitude from RC3 except for NGC~3954, NGC 4682 and NGC 5530 (LEDA). $^{(5)}$Major-axis position angle from RC3. $^{(6)}$Inclination, derived as $\\cos^{2}{i}\\,=\\,(q^2-q_0^2)/(1-q_0^2)$. The observed axial ratio $q$ is taken from RC3 and the intrinsic flattening $q_0=0.11$ has been assumed following Guthrie (1992). $^{(7)}$Heliocentric velocity of the galaxy derived at the centre of symmetry of the rotation curve of the gas. $\\Delta V_\\odot = 10$ \\kms . $^{(8)}$Distance obtained as $V_0/H_0$ with $H_0=75$ \\kms\\ Mpc$^{-1}$ and $V_0$, the systemic velocity derived from $V_\\odot$ corrected for the motion of the Sun with respect to the Local Group, as in  RSA. $^{(10)}$Radius of the 25 $B$ mag arcsec$^{-2}$ isophote derived as $R_{25} = D_{25}/2$ with $D_{25}$ from RC3. $^{(11)}$Absolute total blue magnitude corrected for inclination and extinction from RC3. $^{(12)}$Radial extension of the stellar rotation curve in units of $R_{25}$ from this paper. $^{(13)}$Radial extension of the ionized gas rotation curve in units of $R_{25}$ from this paper. \\end{minipage} \\end{footnotesize} \\end{center} \\label{tab:sample_properties} \\end{table*} ",
        "conclusions": "\\label{sec:conclusions}  We have measured the ionized gas and stellar kinematics along the major axes of seventeen intermediate- to late-type spiral galaxies. The rotation curves and velocity dispersion profiles of ionized gas and star typically extend out to $\\approx0.9\\,R_{25}$ and $\\approx0.5\\,R_{25}$, respectively.  In most of cases the different kinematic behaviour of ionized gas and stars can be easily explained if gas is confined in the disc and supported by rotation while the stars belong mostly  to the bulge and are supported by dynamical pressure. However, kinematic peculiarities have been observed at least in NGC 3054 and NGC 7531.  In addition, we discussed the frequency of counterrotation in disc galaxies. We considered our seventeen spiral galaxies with those  studied in previous papers (Bertola et al.\\ 1996; Corsini et al.\\ 1999, 2003; Vega Beltr\\'an et al.\\ 2001) to build a qualified sample of 50 bright and nearby spirals, ranging from S0/a to Scd, for which the ionized gas and stellar kinematics have been measured along the major axis with the same analytical technique. We found that less than $12\\%$ and less than $8\\%$ (at the $95\\%$ confidence level) of these galaxies host a counterrotating gaseous and stellar disc, respectively. For comparison, we found that $\\sim30\\%$ of S0s host a counterrotating gaseous disc, and Kuijken et al.\\ (1996) estimated that less than $10\\%$ (at $95\\%$ confidence level) host a significant fraction of counterrotating stars.  To interpret the observed frequencies of the gaseous and stellar counterrotating components of disc galaxies, we suggest a scenario in which S0 and spiral galaxies are subject to external gas acquisition with equal probability.  The retrograde acquisition of small amounts of external gas gives rise to counterrotating gaseous discs only in S0 galaxies, since in spiral galaxies the acquired gas is swept away by the pre-existing gas. The formation of counterrotating gaseous discs is favoured in S0 galaxies since they gas-poor systems, while spiral discs host large amounts of gas (Roberts \\& Haynes 1994; Bettoni, Galletta \\& Garc\\`{\\i}a-Burillo 2003), which is corotating with the stellar component. When they acquire external gas in retrograde orbits, the gas clouds of the new retrograde and pre-existing prograde components collide, lose their centrifugal support, and accrete toward the galaxy centre. A counterrotating gaseous disc will be observed only if the mass of the newly supplied gas exceeds that of the pre-existing one (Lovelace \\& Chou 1996; Thakar \\& Ryden 1998). A counterrotating stellar disc is the end-result of star formation in the counterrotating gas disc. For this reason we observe a larger fraction of counterrotating gaseous disks in S0s than in spirals. This also explains why the mass of counterrotating gas in most S0 galaxies is small compared to that of the stellar counterrotating components (Kuijken et al.\\ 1996).  The fraction of S0s with a kinematically decoupled gas disc is consistent with the $50\\%$ that we expect if all the gas in S0s is of external origin (Bertola et al.\\ 1992).  Counterrotating gaseous and stellar discs in spirals are both the results of retrograde acquisition of large amounts of gas, and they are observed with the same frequency. In this framework stellar counterrotation is the end result of star formation in the a counterrotating gaseous disc."
    },
    "0404/astro-ph0404072_arXiv.txt": {
        "abstract": "By collecting optical and infrared photometry and low resolution spectroscopy, we have identified a large number of  low mass stars and brown dwarf candidates belonging to the young cluster ($\\sim$5 Myr) associated with the binary star $\\lambda$ Orionis. The lowest mass object found is a M8.5 with an estimated mass of 0.02 M$_\\odot$ ($\\sim$0.01 M$_\\odot$ for objects without spectroscopic confirmation). For those objects with spectroscopy, the measured strength of the H$\\alpha$ emission line follows a distribution similar to other clusters with the same age range, with larger equivalent widths for cooler spectral types. Three of the brown dwarfs have H$\\alpha$ emission equivalent widths of order 100 \\AA, suggestive that they may have accretion disks and thus are the substellar equivalent of Classical T Tauri stars. We have derived the Initial Mass Function for the cluster. For the substellar regime, the index of the mass spectrum is $\\alpha$=0.60$\\pm$0.06, very similar to other young associations. ",
        "introduction": "The Lambda Ori OB-T association, located at 400 pc (Murdin \\& Penston 1977) is a young stellar group which has not been studied so far in great depth. It is located  inside a fossil giant molecular cloud.   The O8 III star $\\lambda$$^1$ Ori, and to a lesser extent the 11 B stars near to it, excite the  H{\\sc II} region S\\,264.  Making use of the Infrared Astronomical Satellite (IRAS), Zhang et al. (1989) detected a dust ring with a diameter of 9 deg centered around the star $\\lambda$ Ori. This ring is complementary to a shell of neutral hydrogen discovered previously by Wade (1957, 1958). There are two nearby dark clouds within this ring, namely B35 and B30, separated from $\\lambda$ Ori by 2.2 and 2.7 deg, respectively. Based on a H$\\alpha$ emission survey, Duerr, Imhoff \\& Lada (1982), identified three  stellar clusters centered around B30, B35 and $\\lambda$$^1$ Ori, respectively. Those clusters were later confirmed from a statistical point of view by G\\'omez \\& Lada (1998).   Dolan \\& Mathieu (1999, 2001, 2002) collected moderately deep photometry (VRI filters) in an area about 8 sq.deg. centered on the OB association, discovering a significant population of low mass stellar members, and obtained  medium resolution multifiber spectroscopy for those candidates closest to the central star. Their derived distance of 450$\\pm$50 pc is larger than both the distance derived by Murdin \\& Penston (1977) and the value derived by Hipparcos (Perryman et al. 1997) for the five stars in the central area, 380$\\pm$30 pc.  According to Dolan and Mathieu (2002), the turn-off age for the massive stars is of order 6 Myr (see also Murdin \\& Penston 1977 for another age determination based on the masive stars, 4 Myr), although the star formation history might be more complex (Dolan \\& Mathieu 2001).   In this paper, we present additional, much deeper photometry, well beyond the hydrogen burning  limit at 0.072 M$_\\odot$ (Baraffe et al. 1998).  For some of the new candidate members, we have also obtained low resolution spectroscopy, which allow us to add additional clues about their membership and their substellar nature. In our study, we will assume: an age  of 5 Myr, a distance 400 pc --$(m-M)_0$=8.010-- and a reddening of E$(B-V)$=0.12 (Diplas \\& Savage 1994). Section 2 deals with the optical search and  the near infrared counterparts in the 2MASS All Sky Survey. Section 3 presents the analysis of these datasets, whereas the results are summarized in section 4. ",
        "conclusions": "We have collected deep optical photometry in about 0.3 sq.deg around the binary star $\\lambda$ Orionis, extending to well below the substellar boundary. The combination of this dataset set with near infrared photometry and low resolution spectroscopy --i.e., spectral types-- allow us to cull from the initial membership list the possible and probable low mass members of the cluster, both of stellar and substellar nature. We note, however, that additional work is required, in order to study other youth indicators such as low-gravity features and  the detection of lithium. We conclude that the pollution fraction due to interlopers is low (similar to 25 \\% for both the sample with or without spectroscopic information). The faintest object whose membership has been established is a brown dwarf with a mass slightly below 0.020 M$_\\odot$ (based on the Chabrier and Baraffe models) and a M8.5 spectral type. Moreover, H$\\alpha$ equivalent widths have been measured in the spectra.   A plot of the H$\\alpha$\\ equivalent widths as a function of spectral type shows a very similar distribution for Lambda Ori and for the similar age Sigma Ori clusters, with an increase on average for cooler spectral types. Some of the Lambda Orionis stars and brown dwarfs have W(H$\\alpha$) larger than the chromospheric saturation limit. By  analogy with Classical TTauri stars, they might have an accretion disk. We have also derived the Initial Mass Function in the range 4.7-0.02 M$_\\odot$, which shows different types of behavior when displayed as a mass spectrum. Across the stellar/substellar boundary, the index of a power law fit is $\\alpha$=+0.60$\\pm$0.06, quite similar to values recently derived for other young clusters in the same mass range."
    },
    "0404/astro-ph0404591_arXiv.txt": {
        "abstract": "Conditions for accelerated expansion of Friedmann--Robertson--Walker space--time are analyzed. Connection of this scenario with present--day observations are reviewed. It is explained how a scalar field could be responsible for cosmic acceleration observed in present times and predicted for the very early Universe. Ideas aimed at answering whether is that the actual case for our Universe are described. ",
        "introduction": "\\label{sec:intro}  The Standard Big Bang Model (SBB), based on a Friedmann--Robertson--Walker (FRW) Universe, evolves in time with essentially two phases. In the first one the energy related to relativistic matter (known as radiation) dominates over any other form of energy. During that period phase transitions described by particle physics took place to give rise to hadron formation, baryogenesis, nucleosynthesis and so on. Because in an expanding Universe radiation energy dilutes faster than the energy of pressureless matter, in the second phase the latter becomes dominant. Large scale structures (LSS) like galaxies and galaxy clusters formed during that period. In both phases the Universe expands in a decelerated fashion (for a short review on the SBB see the contribution to this book by J.\\ L.\\ Cervantes--Cota\\cite{Cervantes} and for more details the book \\cite{inflation}).  The SBB can easily accommodate phases of accelerated expansion of the Universe. According to cosmological observations, such a phase could correspond to the present state of the observable Universe and seems to be necessary in the very early Universe in order to solve several problems inherent to the SBB, particularly those problems related to the initial conditions.  For the Universe to expand acceleratingly, a very special kind of energy density is required to dominate over the remaining contributions to the total energy budget. This kind of energy is related to a negative pressure. One of the outstanding problems in modern cosmology is to find out what exactly this kind of matter is. It could be the case that there are different explanations to what causes the Universe to undergo cosmic acceleration in the present and in the very early phases of its evolution. As to the present era, a dominating vacuum energy is good enough to explain the observations but it introduces other problems which seem to be very difficult to solve. A good candidate is, instead, a single scalar field with dynamics dominated by its potential energy. This is also the favorite candidate to explain an era of cosmic acceleration in the very early Universe. In both cases, the problem is then to determine what the high-energy physics framework is where such scalar fields arise.  In the next section the condition for an accelerated FRW Universe will be derived. Then, in Sec.~\\ref{sec:scalfield} a scalar field will be described from the cosmological point of view, along with how it could be used to induce an accelerated expansion. Section \\ref{sec:Obs2} is devoted to explaining some ideas aimed at finding observational signatures in the data, allowing us to distinguish the origin of cosmic acceleration. If scalar fields are responsible for the observed and predicted eras of accelerated expansion, the observations should say something about high-energy physics that is outside the scope of Earth-based laboratories. Finally, conclusions are presented in section \\ref{conclusions}.  Throughout this contribution natural units are used, i.e., $c=\\hbar=1$. ",
        "conclusions": "\\label{conclusions}  Cosmic acceleration is a trivial solution to the Einstein equations for an isotropic and homogeneous Universe, as appears to be the one where we live. An accelerated expansion is typical of Universes filled with a kind of energy yielding a strong enough negative pressure. This is the case when the cosmic energy is dominated by the contribution of the vacuum.  Observational evidence strongly suggest that our Universe's evolution includes three well-defined epochs with regards to the increase of the cosmic volume. First, the very early Universe would undergo an accelerated expansion known as inflation. Then, a period of non-accelerated expansion would take place where most of the known kinds of matter and matter structures were formed. Finally, in recent times (with respect to cosmic scales) the Universe would enter a second epoch of accelerated expansion where the corresponding dominated matter-energy content is called dark energy.  A real scalar field is a good candidate for inducing cosmic acceleration. It may help to solve problems arising when a constant vacuum energy is used to explain inflation or the nature of the dark energy. For the required negative pressure, the scalar field dynamics must be dominated by its potential energy.  A hot question in cosmology is whether the observed (predicted) cosmic acceleration is (was) induced by a scalar field. If this is the case, the relevant question is to determine the origin and nature of the corresponding potential. This will open an important window into high energy physics.  Here, an idea was hinted at on how to differentiate between candidates for the dark energy. The proposal is to divide the evolution of the dark energy in periods where the corresponding equation of state could be approximated to be linear. The best--fit values for the corresponding slopes would indicate the favorite candidate. Encouraging results have been obtained in this direction.  It was also explained some of the difficulties that arise when deriving the inflationary potential from observations. It seems like the best that can be done is to indicate generic features of the potentials yielding perturbations spectra matching the measured data. It was emphasized that the use of data on the difference of the tensor and scalar indices of perturbations yields information on the scale--dependence of the tensor to scalar ratio of primordial perturbation amplitudes. This information may be very useful in classifying the inflationary potentials. Finally, a warning was issued about the possibility that the features of the inflaton potential drawn from the observational data could be biased by the order of the approximations used to derive the expressions underlying the calculations.  \\subsubsection"
    },
    "0404/astro-ph0404244_arXiv.txt": {
        "abstract": " ",
        "introduction": "The proper exploitation of astronomical observations requires one to deal with several problems to accurately describe the objects and processes under study. In particular, regarding the physics of the interstellar medium (ISM), it should be stressed that spectral emission data depends on the velocity field solely via its component along the line of sight, through Doppler shifts. Furthermore, this information is necessarily integrated along the line of sight, and radiative transfer leads to expressions in which the contributions of the density and velocity fields are mixed in a complex way~\\citep{hegmann2000}. Hence, to describe the physical conditions and processes in the ISM, one has to rely on a single number (e.g. antenna temperature) for any given direction in the plane of the sky and any given velocity along the line of sight. Although the comparison of antenna temperatures for various tracers helps, it is necessary to solve an inverse problem to have access to the three-dimensional properties of the medium, such as density and velocity, and compare them with various models, for instance to assess the roles of gravity, magnetic field or turbulence.  Indeed, it is now recognized that the different components of the ISM are subjected to turbulent motion. This has been observed in the ionized gas~\\citep{vanlangevelde92}, in {\\sc Hii} regions~\\citep{odell87} and in the neutral atomic phase~\\citep{spicker88,miville2003}, but molecular cloud studies are by far the most numerous~\\citep[see e.g.][]{kleiner85a,kitamura93,miesch94}. Estimates of Reynolds numbers from molecular viscosity in these clouds are of the order of $10^8$, consistent with turbulent flows~\\citep{chandrasekhar49}. Moreover, molecular lines in this phase exhibit suprathermal widths over a wide range of spatial scales, from a few km~s$^{-1}$ in small dark clouds to a few tens of km~s$^{-1}$ in giant complexes, while the thermal dispersion is only of about 0.3~km~s$^{-1}$ for molecular hydrogen at $T=10~$K. These linewidths scale as a power-law of the cloud's size~\\citep{larson81}, with an exponent close to that predicted by the classical theory of turbulence~\\citep{k41}. Such random motions within molecular clouds may in turn account for a number of other properties of spectral lines~\\citep{baker76}, as well as provide support against gravitational collapse, explaining the fact that the lifetime of molecular clouds is larger than their free-fall time~\\citep{scalo85} and that the star formation rate is therefore much smaller than predicted by gravitationally collapsing cloud models~\\citep{zuckerman74}.  Because of this interplay between random motion and many physical processes at work in the ISM in general and in the molecular phase in particular, one needs to accurately describe these turbulent flows~\\citep[see the review by][]{vazquez2000}. As noted earlier, this has to be done using the antenna temperature which represents the emission from a given direction and at a given velocity, and so a number of methods were devised to derive the statistical properties of the three-dimensional fields from those of the observational data. Among these, the velocity channel analysis (VCA) of~\\citet{lp00} is based on an analytical derivation of the properties of channel maps with varying velocity widths. It may however prove difficult to apply to actual observations, as shown by~\\citet{mamd2003}. The modified velocity centroids (MVC) of~\\citet{lazarian2003} are a recent promising attempt to reduce the influence of density fluctuations in the statistical properties of centroids, although it can be argued that they are only defined through their structure function. As a final example of velocity statistics retrieval methods, principal component analysis (PCA), which works on the full position-position-velocity cubes, is meant to decompose data onto an orthogonal basis and derive properties of the velocity field at each scale, as calibrated numerically by~\\citet{bruntetal2003}. The main objective of these works is to relate the scaling behaviour observed in the two-dimensional maps to scaling laws inferred for the three-dimensional fields. For instance,~\\citet{stutzki98} showed that for optically thin media, the spectral index of the integrated emission map is the same as that of the full density field, provided that the depth probed is larger than the transverse scales considered. In a recent work,~\\citet{mamd2003} used numerical simulations to show that the same is true for the normalized velocity centroid with respect to the three-dimensional velocity field. However, this latter result lacks theoretical support, and it is therefore the goal of this paper to present an analytical study aimed at clarifying the relationship between the velocity centroids and the velocity field, within a simple turbulent cloud model.  Given the observational data, one may derive moments of the antenna temperature profiles, each of these moments yielding a potentially informative two-dimensional map. For instance, the zeroth moment is the integrated emission, or intensity, while the first moment is the velocity centroid, which can also be normalized to the intensity~\\citep{munch58,kleiner85b,mamd2003}. For an optically thin line and uniform excitation conditions, the non-normalized centroid can be related to the cloud's total momentum, while the normalized centroid is a measure of average velocity within the medium. While density fluctuations may bias the description of the turbulent motion~\\citep{lazarian2003}, it is however commonly believed, and intuitively plausible, that their effects are somewhat compensated by normalization. To properly assess these, and following~\\citet{kleiner85b}~\\citep[see also][]{scalo84,kitamura93,miesch94}, we shall use autocorrelation functions of the moment maps and relate them to correlation functions of the underlying three-dimensional fields. To this end, we first describe the model we shall use and introduce the notations and assumptions in section~\\ref{sec_model}. We then present a brief summary of how moments of the line profile can be related to the density and velocity fields within the medium (section~\\ref{sec_moments}). Section~\\ref{sec_norm} contains a general study of the statistical properties of intensity, normalized and non-normalized velocity centroid maps as functions of the three-dimensional density and velocity fields' statistics. The equations obtained in the lowest order are then applied to the test case of fractional Brownian motion (fBm) density and velocity fields (section~\\ref{sec_fbm}). Section~\\ref{sec_test} presents a discussion of the various results obtained with respect to earlier works. Our concluding remarks are given in section~\\ref{sec_concl} and details on the calculations can be found in the appendix.  \\begin{center} \\begin{figure}[htbp] \\begin{picture}(16,7) \\put(0.8,0){\\includegraphics[width=16cm]{0139fig1.eps}} \\put(0.85,2.95){\\footnotesize{$x$}} \\put(2.3,4.7){\\footnotesize{$y$}} \\put(9.05,2.1){\\footnotesize{$z$}} \\put(2,2.1){\\footnotesize{$O$}} \\put(3.8,3.6){\\footnotesize{$\\mathbf{r}$}} \\put(2.4,3.4){\\footnotesize{$\\mathbf{x}_1$}} \\put(3.1,2.5){\\footnotesize{$\\mathbf{x}_2$}} \\put(7.5,3.6){\\footnotesize{$\\mathbf{R}$}} \\put(6.3,3.4){\\footnotesize{$\\mathbf{X}_1$}} \\put(7.8,2.4){\\footnotesize{$\\mathbf{X}_2$}} \\put(16.55,0.1){\\footnotesize{$z$}} \\put(16.55,3.05){\\footnotesize{$z$}} \\put(13.32,0.05){\\footnotesize{$0$}} \\put(11.32,0.05){\\footnotesize{$-D/2$}} \\put(14.8,0.05){\\footnotesize{$D/2$}} \\put(9.7,5.75){\\footnotesize{$\\rho$}} \\put(9.7,2.8){\\footnotesize{$v$}} \\put(15.6,1.65){\\footnotesize{$v_0$}} \\put(15.6,4.6){\\footnotesize{$\\rho_0$}} \\end{picture} \\caption{\\label{fig_1} Notations used in the paper (left) and schematics of the turbulent slab model (right). A lowercase boldface letter stands for a three-dimensional vector, while the corresponding uppercase represents its projection on the plane of the sky ($xOy$). The slab is infinite in the $x$ and $y$ directions and is limited to $I_D=[-D/2,D/2]$ in the $z$ direction. The density $\\rho$ and line-of-sight velocity $v$ distributions along a line of sight are shown on the right. The mean values $\\rho_0$ and $v_0$ are taken over the whole slab. } \\end{figure} \\end{center} ",
        "conclusions": "\\label{sec_concl}  An analytical study of a simple slightly compressible turbulent cloud model was presented, assuming homogeneity and isotropy of the turbulent flow. From the expressions of the antenna temperature for an optically thin spectral line and of its successive moments with respect to the line of sight velocity component, we computed the autocorrelation functions of the intensity and of both normalized and non-normalized velocity centroids, which involve averages, along the line of sight, of correlation functions of the three-dimensional density and velocity fields.  To the lowest order, the autocorrelation functions of the velocity centroids behave, with respect to the velocity field, as the autocorrelation function of the intensity with respect to the density field. This sheds light on the numerical result of~\\citet{mamd2003}, who found that, for fractional Brownian motion density and velocity fields, the spectral index of the normalized centroid is equal to that of the velocity field. We derived this result analytically, for separations across the sky much smaller than the cloud's depth, and in real space, while previous studies such as that of~\\citet{goldman2000} were performed in Fourier space. However, the result of~\\citet{mamd2003} holds for fields outside of the validity domain for our calculation.  Comparison of the expansions of the autocorrelation functions of both types of velocity centroids shows that normalization performs a correction of the first order in density fluctuations, although its magnitude remains to be assessed, a task for which numerical simulations are probably necessary. Numerical tests should also provide us with a robust comparison between normalized velocity centroids and the modified velocity centroids of~\\citet{lazarian2003}, which imply corrections of order two in density fluctuations. At present, this comparison has been performed only on simulations of highly compressible and magnetized turbulence, a case beyond the scope of our analytical study, and has shown that, in this particular case, modified velocity centroids provide a more reliable tool than normalized centroids.  In a forthcoming paper, we shall therefore present numerical simulations aimed at assessing the validity domain of our calculations and, beyond normalized and modified velocity centroids, pursuing the search for a better correction scheme able to retrieve the underlying velocity statistics from observational data. \\\\ I wish to acknowledge fruitful discussions with Alex Lazarian during his stay at the \\'Ecole normale sup\\'erieure. His suggestions, in the early stages of this work, proved very helpful. I also wish to thank Enrique V\\'azquez-Semadeni for his careful reading of the manuscript and insightful remarks."
    },
    "0404/astro-ph0404134_arXiv.txt": {
        "abstract": "Extraction of the CMB (Cosmic Microwave Background) angular power spectrum is a challenging task for current and future CMB experiments due to the large data sets involved. Here we describe an implementation of MASTER (Monte carlo Apodised Spherical Transform EstimatoR) described in \\citet{HIV02} which exploits the destriping technique as a map-making method. In this method a noise estimate based on destriped noise-only MC (Monte Carlo) simulations is subtracted from the pseudo angular power spectrum. As a working case we use realistic simulations of the {\\it PLANCK} LFI (Low Frequency Instrument). We found that the effect of destriping on a pure sky signal is minimal and requires no correction.  Instead we found an effect related to the distribution of detector pointings, which affects the high-$\\ell$ part of the power spectrum. We correct for this by subtracting a ``signal bias'' estimated by MC simulations.  We also give analytical estimates for this signal bias.  Our method is fast and accurate enough (the estimator is un-biased and errors are close to theoretical expectations for maximal accuracy) to estimate the CMB angular power spectra for current and future CMB space missions.  This study is related to {\\it PLANCK} LFI activities. ",
        "introduction": "\\label{sec:intro}  In the favoured model of structure formation driven by inflation, primordial fluctuations are expected to be Gaussian distributed. In this case all the statistical information encoded into CMB (Cosmic Microwave Background) anisotropies is completely described by their angular power spectrum $C_\\ell$. The main issue in this paper is the extraction of the $C_\\ell$ values starting from an observed CMB map. In recent years a maximum likelihood approach to the problem has been developed \\citep{GOR96} and successfully applied to the {\\it COBE} - DMR data \\citep{BEN96}. This process, however, involves a number of operations scaling as $\\sim N_{\\rm pix}^3$, where $N_{\\rm pix}$ is the number of pixels that cover the sky.  Recent ground-based and balloon-borne experiments have improved our knowledge of the CMB anisotropy due to the better angular resolution and higher sensitivity of these experiments. Due to the sizes of their data sets they have already posed challenges in extracting an unbiased estimate of $C_\\ell$. This will become more demanding with current and future satellite experiments like {\\it WMAP} \\citep{BEN03} and {\\it PLANCK} \\citep{TAU00} which produce maps with $N_{\\rm pix} \\simeq {\\rm few}\\times 10^6$. Therefore, brute force maximum likelihood estimation of the CMB angular power spectrum is not feasible without approximations or assumptions specific to particular experimental or observing strategies \\citep{OHS99,WAN03}.  New methods are under study and some of them have already produced interesting results. \\citet{HIV02} introduced the MASTER technique used to construct the BOOMERanG angular power spectrum in \\citet{NET02}. An extension of MASTER is being prepared to treat {\\it PLANCK}-like data sets. The core of the MASTER technique is {\\it i)} a high-pass filter applied to the TOD (Time Ordered Data) as part of a naive map-making process in order to reduce instrumental non-white noise, and {\\it ii)} MC (Monte Carlo) simulations of pure signal, pure noise, and signal plus noise to correct for data filtering and instrumental noise, and for estimating the error bars of $C_\\ell$. The approach has been shown to be unbiased and it gives error bars close to theoretical ones when cosmic variance and instrumental noise properties are concerned \\citep{HIV02}.  \\citet{BAL02} follow a MASTER-like approach but instead of using a high-pass filter they consider an IGLS (Iterative Generalized Least Square) map-making algorithm already developed to treat {\\it PLANCK} data \\citep{NAT01}. Furthermore, they estimate noise properties needed for proper MC simulations directly from the data \\citep{NAT02}. The extracted $C_\\ell$ are unbiased and have close to optimal error bars.  In this paper we present an approach similar to that of \\citet{HIV02} but exploiting the destriping algorithm \\citep{BUR97,DEL98,MAI99a,MAI99b} in the map-making process. Although destriping algorithms are considered approximations of proper IGLS map-making in the sense that they will not necessarily produce the minimum variance map, they are able to provide estimates of various systematic effects, remove drifts from TOD and return cleaned TOD (e.g., see \\citealt{MEN02} for an application with periodic fluctuations). Furthermore, destriping makes no assumptions on the beam shape. It has been demonstrated \\citep{MAI02} that the instrument noise is the driver of the destriping performance regardless of the beam shape (including sidelobes).  A generalized maximum likelihood approach to the destriping method has been implemented \\citep{KEI03} which is able to fit different sets of base functions (in addition to the simple constant baseline) and, in principle, could better remove the contributions of different systematic effects from the TODs. All these properties make destriping attractive. Additionally, it is fast and needs no prior information on the instrumental noise. Note, however, that as we combine destriping with the MASTER approach for the $C_\\ell$ estimation, this part utilizes information on noise properties. Both iterative and non-iterative methods to estimate the noise characteristics directly from the data have been proposed \\citep{DOR01,NAT02}. These methods have been applied in the $C_\\ell$ estimation \\citep[e.g.][]{NET02,BAL02}.  This paper is organised as follows. In Section~\\ref{sec:destri} we describe the destriping technique. In Section~\\ref{sec:MASTER} we review the MASTER approach for the extraction of $C_\\ell$ from a map. We have then applied our combination of destriping with the MASTER approach to simulations of one {\\it PLANCK} LFI 100 GHz detector. The applied scanning strategy and sky coverages are described in Section~\\ref{sec:scanning_strategy}. In Section~\\ref{sec:filter_function} we explain our findings concerning the filter function and signal bias that are possible means to model the effects caused by our map-making method on the spectrum estimates. The analysis pipeline and the CPU times required to run it are explained in Section~\\ref{sec:times}. Section~\\ref{sec:simul} presents the simulations and the simulation results. We draw our conclusions in Section~\\ref{sec:conclusion}. In Appendix~\\ref{sec:pointing_distribution} we describe in more details the signal bias approach for modelling the effect of the distribution of detector pointings in output map pixels. ",
        "conclusions": "\\label{sec:conclusion}  We have demonstrated that the combination of destriping technique and our MASTER approach can tackle the extraction of the CMB angular power spectra. The approach was found to work well yielding accurate estimates of the true power spectrum of the sky. As a practical example, we considered the 100~GHz channels of the LFI instrument of the {\\it PLANCK} satellite.  Destriping as a map-making method is general in a sense that it requires no prior information either on the instrument noise properties or on the actual beam shape. We can expect that even in the case of non-stationary noise we are able to recover the baseline magnitudes accurately. The MASTER technique requires knowledge of the noise characteristics. In realistic CMB experiments the noise parameters need to be estimated from the measured data. However, we did not address the noise estimation techniques in this study. Thus the correct instrument noise model was assumed throughout this paper.   We found that the effect of the sky signal on destriping does not cause a significant distortion on the power spectrum estimates.  Therefore no filter function is needed to compensate for it.  Instead we discovered an effect related to the clustering of the detector pointings.  Because of the high stability of the motion of the {\\it PLANCK} satellite the centre points of the measurements taken during the 1~h period between satellite spin-axis repointings form tight clusters of 60. In our idealized simulation these were all assumed to fall on a single point, which exaggerates the effect somewhat. This leads to a noticeable effect at high $\\ell$ (for $\\ell > 700$ in our simulations).  We were able to account for this effect, and remove it, by introducing the concept of signal bias in the $C_\\ell$ estimate, and determining it by signal-only MC simulations.  We also derived a reasonably accurate analytical estimate for the signal bias.  It was shown that the signal bias will be reduced when data from several detectors are combined, since that increases the density of detector pointings. In addition, due to satellite spin axis nutation (amplitudes up to 1.5~arcmin, \\citealt{LEE02}) and spin rate variation the detector pointings from successive scan circles will not fall exactly on top of each other, in reality. Therefore, in the case of the real {\\it PLANCK} experiment this signal bias may be small enough to be ignored.  An ensemble of angular power spectrum estimates were produced by applying our estimation method to the maps of the signal+noise MC realizations. Our results showed that, after removing the signal and noise biases, any remaining bias was buried under the MC variations; if any bias exists its relative level was not higher than 0.1~per cent.  There was a good match between the $\\pm 1\\sigma$ error bars obtained from the MC realizations and derived from the analytical model of the pseudo-$C_\\ell$ estimators \\citep{EFS03}. This indicates that the implementation related performance losses of our method are small."
    },
    "0404/astro-ph0404302_arXiv.txt": {
        "abstract": "We investigate environmental effects on evolution of bright cluster galaxies in a $\\Lambda$-dominated cold dark matter universe using a combination of dissipationless $N$-body simulations and a semi-analytic galaxy formation model. We incorporate effects of ram-pressure stripping (RPS) and minor merger-induced small starburst (minor burst) into our model. By considering minor burst, observed morphology-radius relation is successfully reproduced. When we do not consider minor burst, the RPS hardly increases the intermediate B/T population. In addition, the RPS and minor burst are not important for colours or star formation rates of galaxies in the cluster core if star formation time-scale is properly chosen, because the star formation  is sufficiently suppressed by consumption of the cold gas. We also find that SF in bulge-dominated galaxies is mainly terminated by starburst induced by major mergers in all environments. ",
        "introduction": "It has been found that galaxy morphology is a function of environment (Dressler 1980; Whitmore et al. 1993) and redshift (Dressler et al. 1997). Colours and star formation rate (SFRs) of galaxies also show similar dependence on environment and redshift (e.g. Butcher \\& Oemler 1984; Lewis et al. 2002). To account for these observational trends, several mechanisms that may suppress the star formation (SF) and transform one morphological type into another have been proposed.   Interaction between galaxies is one possible process to promote morphological transformation. Numerical simulations confirmed that major mergers produce galaxies resembling ellipticals as merger remnants (Barnes 1996) and that accretion of small satellites onto their host spiral lead a host spiral to S0 type (Walker et al. 1996). Because the galaxy merger triggers starburst, the cold gas contained in original galaxies is exhausted in a very short time.  The second mechanism is removal of hot reservoir. In denser environment such as clusters, diffuse hot gas reservoir that is originally confined in haloes of non-central galaxies becomes part of intracluster medium. A galaxy whose hot gas reservoir is removed slowly exhausts its cold gas in its SF time-scale (strangulation).  Above 2 processes, i.e. major merger-induced starburst and strangulation, has been incorporated in most of semi-analytic (SA) models which have failed to reproduce intermediate $B/T$ population (Okamoto \\& Nagashima 2001; Diaferio et al. 2001). In this paper, we introduce two additional processes in order to solve this problem. One is ram-pressure stripping (RPS) of the cold gas from galactic disks (Gunn \\& Gott 1972). Another is minor merger-induced small starburst. Hydrodynamic simulations showed some fraction of disk gas is fuelled to the galactic centre by a minor merger (Walker et al. 1996). By using a combination of $N$-body simulations with a SA model, we study roles of above 4 processes in evolution of galaxies and what process is responsible for the formation of intermediate $B/T$ galaxies. ",
        "conclusions": "We have investigated the effects of major mergers, strangulation, RPS of cold disk gas, and minor mergers on the evolution of bright cluster galaxies. We have used a combination of the cosmological $N$-body simulations and the SA galaxy formation model. This method enables us to study above environmental effects in a fully cosmological context.  We have determined the model parameters of the reference model to reproduce galaxy properties at $z = 0$.  Our results are summarised as follows. \\begin{enumerate} \\item The process that terminates SF in early-type galaxies ($B/T > 0.4$) in all environments is starburst trigger by a major merger. \\item If we adopt appropriate SF time-scale, so as to reproduce the observed cold gas mass fraction in the field, the dominant process that determines colours of galaxies in the cluster core is the strangulation. \\item Since the strangulation sufficiently suppresses SF in cluster cores, the effect of the RPS is hardly observed. \\item Minor burst does not affect galaxy properties except for morphology. \\item Without minor bursts the fraction of intermediate $B/T$ galaxies in clusters becomes too small. The model with the minor burst can reproduce the observed morphology-radius relations. We conclude that the minor burst is essential process to form intermediate $B/T$ galaxies. \\item The RPS hardly increases intermediate population in the cluster without minor bursts. When the minor burst is taken into account, the RPS increases the intermediate galaxy fraction in the cluster cores. \\end{enumerate}"
    },
    "0404/astro-ph0404187_arXiv.txt": {
        "abstract": "We present an analysis of the 2--10\\,keV {\\it XMM-Newton}/EPIC-pn spectrum of the Seyfert-1 galaxy NGC~4593.  Apart from the presence of two narrow emission lines corresponding to the K$\\alpha$ lines of cold and hydrogen-like iron, this spectrum possesses a power-law form to within $\\sim 3-5\\%$.  There is a marked lack of spectral features from the relativistic regions of the black hole accretion disk.  We show that the data are, however, consistent with the presence of a radiatively-efficient accretion disk extending right down to the radius of marginal stability if it possesses low iron abundance, an appropriately ionized surface, a very high inclination, or a very centrally concentrated emission pattern (as has been observed during the Deep Minimum State of the Seyfert galaxy MCG--6-30-15).  Deeper observations of this source are required in order to validate or reject these models. ",
        "introduction": "The fluorescent K$\\alpha$ emission line of iron is currently the best probe we have to study strong-field gravitational effects in the vicinity of black holes.  This line, together with an associated backscattered continuum, is readily formed when a hard X-ray continuum source irradiates the surface of a relatively cold and optically-thick slab of gas (Basko 1978; Guilbert \\& Rees 1988; Lightman \\& White 1988; George \\& Fabian 1991; Matt et al. 1991).  Nowadays, the hard X-ray source is usually identified with thermal Comptonization from an accretion disk corona, and the optically-thick structure as the accretion disk itself (see Reynolds \\& Nowak 2003 for a recent review).  The diagnostic power of these spectral features lies in investigations of their Doppler broadening and gravitational redshifting (Fabian et al. 1989; Laor 1991).  The best example to date of using these features to probe strong-field gravity is the Seyfert-1.2 galaxy MCG--6-30-15 (Tanaka et al. 1995; Wilms et al. 2001; Fabian et al. 2002; Reynolds et al. 2004).  This object displays a highly-broadened and skewed iron line that is strongly suggestive of emission from an accretion disk reaching down to near the radius of marginal stability for a rapidly-rotating black hole.  As yet, there is no competing model that can explain, in detail, the iron line feature in MCG--6-30-15.  However, it is an open question whether these relativistic spectral features are generic in the spectra of various classes of active galactic nuclei (AGN).  Nandra et al. (1997a) used data from the {\\it ASCA} observatory to conclude that relativistically-broadened iron lines are very common features in the X-ray spectra of Seyfert-1 nuclei.  On the other hand, {\\it ASCA} found that these emission lines were often weaker and/or narrower in the X-ray spectra of low-luminosity AGN (e.g., Reynolds, Nowak \\& Maloney 2000; Terashima et al. 2002), high-luminosity AGN (Nandra et al. 1997b) and radio-loud AGN (Sambruna, Eracleous \\& Mushotzky 1999).  Recent results from the European Photon Imaging Camera (EPIC) on broad the {\\it XMM-Newton} observatory have painted a more complex picture. While {\\it XMM-Newton} has, indeed, found undisputed cases of broad iron lines in the Seyfert galaxies MCG--6-30-15 (Wilms et al. 2001, Fabian et al. 2002), MCG--5-23-16 (Dewangan, Griffiths \\& Schurch 2003), NGC~3516 (Turner et al. 2002), Mrk335 (Gondoin et al.  2002), and Mrk766 (Pounds et al. 2003a), other Seyfert-1 galaxies appear to show an absence of such features, with the best example to date being NGC~5548 (Pounds et al. 2003b).  Clearly, the presence or absence of a relativistic iron line depends upon currently unknown factors and is not a simple function of AGN class.  Progress must be made by careful analysis of as many AGN X-ray spectra as possible.  With this motivation, this {\\it Paper} presents a careful analysis of the hard-band {\\it XMM-Newton} EPIC-pn spectrum of the Seyfert-1 galaxy NGC~4593 ($z=0.009$).  Section~2 describes in brief our observation and data reduction techniques.  Section~3 presents an analysis of the 2--10\\,keV spectrum of NGC~4593 and demonstrates a marked absence of a ``standard'' relativistic iron line, a result which is discussed in more detail in Section~4. Section~5 summarizes our principal conclusions. ",
        "conclusions": "In this paper, we have analyzed the 2--10\\,keV {\\it XMM-Newton} EPIC-pn spectrum of NGC~4593, with particular emphasis on the implications of this spectrum for the nature of the accretion flow in this Seyfert-1 nucleus. Only two spectral features are detected which can be identified with narrow K$\\alpha$ emission lines of cold and hydrogen-like iron at 6.40\\,keV and 6.97\\,keV, respectively.  Once these have been modelled, the spectrum has a power-law form (to within a 3-4\\% accuracy) across the 2--10\\,keV band.  We fail to detect any X-ray reflection features from a relativistic accretion disk.  However, we show that, even if a radiatively-efficient geometrically-thin complete accretion disk exists, its X-ray reflection signatures would be buried in the noise if the accretion disk has either a very centrally concentrated irradiation profile or an appropriately ionized surface.  Either of these models can be tested by longer EPIC observations which would be sensitive to the subtle features displayed by a highly blurred or ionized reflection spectrum.  If longer observations still fail to detect any accretion disk signatures, we are forced to consider other scenarios.  Firstly, the inner disk may be very hot and optically-thin, thereby being incapable of producing any X-ray reflection features. Secondly, the observed X-ray continuum might be highly anisotropic, being strong beamed away from the disk (and towards the observer) thereby rendering any disk features undetectable."
    },
    "0404/astro-ph0404378_arXiv.txt": {
        "abstract": "The conceptual difficulties associated with a cosmological constant have led to the investigation of alternative models in which the equation of state parameter, $w=p/\\rho$, of the dark energy evolves with time. We show that combining the  supernova type Ia observations {\\it with the constraints from WMAP observations} restricts large variation of $\\rho(z)$ at low redshifts. The combination of these two observational constraints is stronger than either one. The results are completely consistent with the cosmological constant as the source of dark energy. ",
        "introduction": "Well before the data from the high redshift supernova project became available, several independent constraints indicated the existence of a cosmological constant \\citep{crisis1,crisis2,crisis3}. In the last decade, observational evidence for an accelerating universe has become conclusive with almost all other possibilities being ruled out by observations of high redshift supernovae \\citep{nova_data1,nova_data2,nova_data3} and the cosmic microwave background radiation (CMBR) \\citep{boomerang,wmap_params}. The accelerated expansion of the universe requires either a cosmological constant or some form of dark energy \\citep{review1,review2,review3,review4} to drive the acceleration, with $w\\equiv p/\\rho < -1/3$. Although a cosmological constant is the simplest solution from a phenomenological point of view (requiring just one fine tuned parameter), there is no natural explanation of the small observed value. This has led theorists to develop models in which a field, typically a scalar field, provides the source of dark energy, e.g., quintessence \\citep{quint1,quint2,quint3,quint4,quint5,quint6,quint7,quint8}, k-essence \\citep{k-essence1,k-essence2,k-essence3,k-essence4,k-essence5}, tachyons \\citep{tachyon1,tachyon2,tachyon3,tachyon4,tachyon5,tachyon6,tachyon7,tachyon8,tachyon9,tachyon10,tachyon11,tachyon12}, phantom fields \\citep{phantom1,phantom2,phantom3,phantom4,phantom5,phantom6,phantom7,phantom8,phantom9,phantom10,phantom11}, branes \\citep{brane1,brane2,brane3,brane4,brane5}, etc. In the absence of significant {\\it spatial} variation in the dark energy, the key difference between such models and the one with the cosmological constant is that, in general, $w$ is a function of redshift $z$ in the former. Perturbations in dark energy also lead to observable signatures, though these can easily be confused with other physical effects \\citep{perturb1,perturb2,perturb3}. There have also been a few proposals for unified dark matter and dark energy \\citep{unified_dedm1,unified_dedm2,unified_dedm3} but these models are yet to be  developed in sufficient detail to allow direct comparison with observations in a fruitful manner. If the current observations had excluded $w=-1$ then one could have immediately ruled out cosmological constant as a candidate; but since this is not the case, direct exploration of $w(z)$ at different redshifts, in order to check possible dependence of $w$ on  $z$, is of importance. Observations constrain the entire suite of parameters that describe cosmological parameters and while it is possible to choose other cosmological parameters so that $\\Lambda$CDM is not allowed, these models have not been ruled out so far.  It has been known for some time that supernova observations and constraints from structure formation can be combined to put stringent limits on models for dark energy \\citep{lss_de}. Several attempts have been made in the past to constrain the equation of state for dark energy, along with  other cosmological parameters, using the observations of galaxy clustering, temperature anisotropies in the CMBR and the high redshift supernovae \\citep{lss_de,constraints_1,constraints_2,constraints_3,constraints_4,constraints_5,constraints_7,constraints_8,constraints_9,constraints_10}.  This work, which is in the same spirit, focuses exclusively on constraining the variation of equation of state for dark energy by using a combination of such observations. We wish to constrain the variation of dark energy while keeping most of the other cosmological parameters fixed around their favoured values. In particular, we study the effect of a varying equation of state on the angular power spectrum of the  CMBR fluctuations. Dark energy is not expected to be dynamically significant  at the time of decoupling and --- in fact --- models in which this is not true are plagued by slow growth of density perturbations \\citep{quint_growth1,quint_growth2,quint_growth3}. Nevertheless, evolving dark energy will affect the features of temperature anisotropies in the CMBR in at least two ways: (1)~The angular scale of features in temperature anisotropy, like the peaks, will change since the angular diameter distance depends on the form of $w(z)$. (2)~The integrated Sachs-Wolfe (ISW) effect will also depend on the nature of dark energy and its evolution, this effect is more relevant at small $l$. Thus observations of temperature anisotropies in the CMBR can be used to constrain the {\\it evolution} of dark energy. Combined with the supernova observations, this allows us to put tight constraints on the equation of state of dark energy and its evolution. Our approach here is to use the full WMAP angular power spectrum in order to ensure that both the effects mentioned above are captured in the analysis. For lower multipoles ($l \\leq 20$) signals from the ISW effect, perturbations in dark energy and reionisation need to be disentangled, however the relative importance of this part of the angular power spectrum is limited as we use the full angular power spectrum from WMAP. Our aim is to demonstrate that the combination of WMAP observations and high redshift supernova observations is a very powerful constraint on variations in dark energy, certainly more powerful that either of the observations used in isolation. {\\it As far as we know, most previous attempts to constrain the dark energy sector using WMAP data have not used the full angular power spectrum. } ",
        "conclusions": "The observations we have considered constrain distances and the main constraint from CMB observations is on a reduced quantity $w_{eff}$ which is the integrated effective value for the angular diameter distance. If $w(0) < w_{eff}$ ($w(0) > w_{eff}$) then $w'(0) > 0$ ($w'(0) < 0$) is expected, indeed this explains the orientation of the allowed region in figure~\\ref{fig:w1w0}. There is a strong degeneracy between $w_{eff}$ and $\\Omega_{nr}$, $\\partial w_{eff}/\\partial \\Omega_{nr} < 0 $ (see Figure \\ref{fig:w0omega}). The cosmological constant itself is ruled out at $99\\%$ confidence level for $\\Omega_{nr} > 0.375$ and only phantom models survive beyond this, thus if other observational constraints were to rule out $\\Omega_{nr} \\leq 0.375$ then we will be forced to work with phantom models. (Note that the Cosmological constant and other non-phantom models are not ruled out by supernova observations for these values of $\\Omega_{nr}$.) But with current observations, such a constraint does not exist and the cosmological constant model is allowed. Hence observations do not {\\it require} varying dark energy or phantom models even though such models are consistent with observations. The detailed dependence on $\\Omega_{nr}$ and other parameters will be explored in detail in a later work \\citep{jbp2}.  In our analysis we have not taken into account any perturbations in dark energy. This effect becomes increasingly important as $w$ approaches zero, leading to suppression of perturbations in matter \\citep{quint_growth1,quint_growth2,quint_growth3}. This effect is relevant only at $l \\leq 20$ in most models and hence its relative importance is small as we are using the full WMAP angular power spectrum here. Perturbations in dark energy also have a non-adiabatic component, but this cannot be incorporated in our analysis as it requires a detailed model for dark energy whereas we are working with parameterised variations. We are carrying out a more detailed analysis where we allow parameters like $n$, $h$, $\\Omega_B$, $\\tau$, etc. to vary. Preliminary results suggest that after marginalizing over other parameters the region in the $w(0)-w'(0)$ allowed by CMB observations increases, mainly on the lower side of the region shown in figure~\\ref{fig:w1w0}. Supernova observations then provide stronger constraints for smaller $w'(0)$ while the CMB observations continue to provide much stronger constraints on large $w'(0)$. Adding other observational constraints from structure formation reduces the allowed region significantly as extreme variations of parameters allowed by WMAP data predict too much or too little structure formation \\citep{jbp2}. These observations are not very relevant in the present work as we are keeping most parameters fixed.  Our analysis of models and observations shows that the allowed variation of $\\rho^{DE}(z)$ is strongly constrained by a combination of WMAP data and observations of high redshift supernovae. Indeed, given the allowed window as seen in Fig.~\\ref{fig:wenvelope}, the cosmological constant seems to be the most attractive dark energy candidate."
    },
    "0404/astro-ph0404008_arXiv.txt": {
        "abstract": "We extend our earlier work on ambipolar diffusion induced formation of protostellar cores in isothermal sheet-like magnetic interstellar clouds, by studying nonaxisymmetric collapse for the physically interesting regime of magnetically critical and supercritical model clouds ($\\mui \\geq 1$, where $\\mui$ is the initial mass-to-magnetic flux ratio in units of the critical value for gravitational collapse). Cores that form in model simulations are effectively triaxial, with shapes that are typically closer to being oblate, rather than prolate. Infall velocities in the critical model ($\\mui = 1$) are subsonic; in contrast, a supercritical model ($\\mui = 2$) has extended supersonic infall that may be excluded by observations. For the magnetically critical model, ambipolar diffusion forms cores that are supercritical ($\\muc > 1$) and embedded within subcritical envelopes ($\\muenv < 1$). Cores in our models have density profiles that eventually merge into a near-uniform background, which is suggestive of observed properties of cloud cores. ",
        "introduction": "Magnetic fields play an important role in star formation, especially in the early stages of core formation and collapse; measured mass-to-flux ratios of molecular clouds yield an average that is $\\sim 1 - 2$ times greater than the critical value for collapse (Crutcher 1999). However, observational biases tend to push toward higher values of measured mass-to-flux ratio (Crutcher 2003, private communication), so that moderately subcritical cloud regions are not ruled out. Dense cores within molecular clouds are the sites of star formation, with detected infall up to $\\approx 0.5 \\, \\cs \\approx 0.1 \\kms$ on scales $\\lesssim 0.1 \\pc$, (e.g., in L1544, Tafalla et al. 1998; Williams et al. 1999), where $\\cs$ is the isothermal sound speed. Ciolek \\& Basu (2000) have fit the main features of the observed velocity and density profiles in L1544, modeling it as an axisymmetric supercritical core embedded in a moderately subcritical envelope.  Axisymmetry is clearly an idealization to real cores. Observations suggest a typical projected axis ratio of 0.5 for cores (Myers et al. 1991), and deprojections of the distribution of shapes imply intrinsically triaxial but nearly oblate cores (Jones, Basu, \\& Dubinski 2001). Polarized emission measurements from dense cores also imply triaxiality (Basu 2000). More generally, detailed submillimeter maps of star-forming regions reveal significant irregular structure and multiple cores (Motte, Andr\\'e, \\& Neri 1998). Theoretical nonaxisymmetric magnetic models of the collapse and fragmentation of a single core were presented by Nakamura \\& Hanawa (1997) and Nakamura \\& Li (2002), without and with ambipolar diffusion, respectively, using the magnetic thin-disk approximation (Ciolek \\& Mouschovias 1993, hereafter CM93). The early stages of core formation in a nonaxisymmetric infinitesimally thin subcritical sheet (including the effects of magnetic tension but ignoring magnetic pressure) were studied by Indebetouw \\& Zweibel (2001). Here, we also study a planar cloud that is perpendicular to the mean magnetic field, focusing on the case of either exactly critical or decidedly supercritical cores; these cases are shown to lead to observationally distinguishable outcomes. We again use the thin-disk approximation, which allows for finite thickness effects, and explicitly includes the effects of both magnetic pressure and tension. \\vspace{-4ex} ",
        "conclusions": "Ambipolar diffusion leads to a nonuniform distribution of mass-to-flux ratio, in a natural extension of the process described by Mouschovias (1978). Stars form preferentially in the most supercritical regions. Core shapes are somewhat triaxial, and usually more nearly oblate than prolate. The core column density eventually merges into a near-uniform background value. In the critical ($\\mui = 1$) model, a surrounding region is established which is mildly magnetically subcritical (due to flux redistribution) and infall motions both inside and outside cores are subsonic. Conversely, the supercritical ($\\mui = 2$) model exhibits supersonic motions within cores, and extended rapid motions outside them. The critical model requires a significantly longer time to develop gravitational instability; however, we caution that the growth time for both models are likely upper limits due to the possibility of nonlinear perturbations in more realistic situations. We also note that all motions in these models are fundamentally gravitationally driven; the neutral speeds are typically greater than those of the ions."
    },
    "0404/hep-ph0404198_arXiv.txt": {
        "abstract": "\\PRE{\\vspace*{.3in}} Dark matter may be composed of superWIMPs, superweakly-interacting massive particles produced in the late decays of other particles.  We focus on the case of gravitinos produced in the late decays of sleptons or sneutrinos and assume they are produced in sufficient numbers to constitute all of non-baryonic dark matter. At leading order, these late decays are two-body and the accompanying energy is electromagnetic.  For natural weak-scale parameters, these decays have been shown to satisfy bounds from Big Bang nucleosynthesis and the cosmic microwave background.  However, sleptons and sneutrinos may also decay to three-body final states, producing hadronic energy, which is subject to even more stringent nucleosynthesis bounds.  We determine the three-body branching fractions and the resulting hadronic energy release.  We find that superWIMP gravitino dark matter is viable and determine the gravitino and slepton/sneutrino masses preferred by this solution to the dark matter problem.  In passing, we note that hadronic constraints disfavor the possibility of superWIMPs produced by neutralino decays unless the neutralino is photino-like. ",
        "introduction": "\\label{sec:introduction}  SuperWIMPs, superweakly interacting massive particles produced in the late decays of weakly-interacting massive particles (WIMPs), are promising non-baryonic dark matter candidates~\\cite{Feng:2003xh,Feng:2003uy}. Well-motivated superWIMP candidates are the gravitino in supersymmetric models~\\cite{Feng:2003xh,Feng:2003uy,Ellis:1984er,Bolz:1998ek,% Bolz:2000fu,Ellis:2003dn} and the first excited graviton in universal extra dimension models~\\cite{Feng:2003xh,Feng:2003nr}.  The supersymmetric possibility is realized naturally in supergravity with a gravitino lightest supersymmetric particle (LSP) and a slepton or sneutrino next-to-lightest supersymmetric particle (NLSP). (Throughout this paper, ``slepton'' denotes a charged slepton.) Both the gravitino and the NLSP have weak-scale masses.  The NLSP freezes out as usual with a relic density near the observed value.  However, after time $t \\sim 10^4 - 10^8~\\s$, it decays through \\begin{eqnarray} \\tilde{l} &\\to& l \\, \\tilde{G} \\nonumber \\\\ \\tilde{\\nu} &\\to& \\nu \\, \\tilde{G} \\ . \\label{decay} \\end{eqnarray} The gravitino then inherits much of the relic density of the slepton or sneutrino~\\cite{Covi:1999ty}, and its relic density is naturally of the right magnitude without the introduction of new scales.  This is in contrast to other gravitino dark matter scenarios, where the gravitino is a thermal relic and the desired density is obtained by an appropriately chosen gravitino mass $m_{\\Gravitino} \\sim \\kev$~\\cite{Pagels:ke}, or the gravitino is produced during reheating~\\cite{Moroi:1993mb,Bolz:2000fu}, where the desired dark matter density is obtained only for a fine-tuned reheat temperature $T_R \\sim 10^{10}~\\gev$.  The decays of \\eqref{decay} occur well after Big Bang nucleosynthesis (BBN).  An immediate concern, therefore, is that they might destroy the successful light element abundance predictions of BBN.  In fact, BBN is not the only worry --- the Planckian spectrum of the cosmic microwave background (CMB), the diffuse photon background, and bounds on late time entropy production from the coincidence between BBN and CMB baryometry all impose significant constraints.  As demonstrated in Refs.~\\cite{Feng:2003xh,Feng:2003uy}, however, these constraints on the electromagnetic (EM) energy released in the decays of \\eqref{decay} exclude some of the weak scale parameter space, but leave much of it intact.  Of course, at the border between the excluded and allowed regions, slight deviations from standard cosmological predictions are expected, providing possible signals in future observations.  In fact, the existing anomaly in the $^7$Li abundance prediction of standard BBN may already be interpreted as a signal of superWIMP dark matter~\\cite{Feng:2003uy}.  Such signals are particularly welcome, since superWIMP dark matter is so weakly interacting that it is impossible to detect through conventional direct and indirect dark matter searches.  The previous work, however, neglected hadronic energy produced in WIMP decays.  This was natural, as the WIMP decays of \\eqref{decay} contribute only to electromagnetic energy, as we will discuss below. However, the three-body decays \\begin{eqnarray} \\tilde{l} &\\to& l Z \\tilde{G} \\ , \\ \\nu W \\tilde{G} \\nonumber \\\\ \\tilde{\\nu} &\\to& \\nu Z \\tilde{G} \\ , \\ l W \\tilde{G} \\ , \\end{eqnarray} produce hadronic energy when the $Z$ or $W$ decays hadronically. Hadronic energy release is very severely constrained by the observed primordial light element abundances~\\cite{Reno:1987qw,% Dimopoulos:1987fz,Dimopoulos:1988ue,Kohri:2001jx,Kawasaki:2004yh}, and so even subdominant hadronic decays could, in principle, provide stringent constraints.  These three-body decays may be kinematically suppressed when $m_{\\tilde{l}, \\tilde{\\nu}} - m_{\\tilde{G}} < m_W, m_Z$, but even in this case, four-body decays, such as $\\tilde{l} \\to l \\gamma^* \\tilde{G} \\to l q\\bar{q} \\tilde{G}$, contribute to hadronic cascades and may be important.  In this paper, we determine the hadronic branching fractions and compare them to BBN constraints on hadronic energy release.  Although we focus on the supersymmetric case, our results may be extended straightforwardly with minor numerical modifications to the case of graviton superWIMPs in extra dimensions.  In evaluating the constraints, there are two possible approaches.  As the superWIMP relic abundance is automatically in the right range, a natural assumption is that superWIMP gravitinos make up all of the non-baryonic dark matter, with \\begin{equation} \\OmegaSWIMP \\simeq 0.23 \\ . \\label{omega} \\end{equation} This is the approach taken here.  Note that $\\OmegaSWIMP$ need not be identical to $\\Omega_{\\gravitino}$ --- not all relic gravitinos need be produced by NLSP decays.  Our assumption therefore requires that the other sources of gravitinos be insignificant.  For the supersymmetric case, this typically implies an upper bound on reheat temperatures of $T_R \\alt 10^{10}~\\gev$~\\cite{Moroi:1993mb,Bolz:2000fu}. Higher reheat temperatures may also be allowed~\\cite{Buchmuller:2003is,Kolb:2003ke} and are desirable, for example, to accommodate leptogenesis.  For the universal extra dimension case, the requirement of insignificant Kaluza-Klein graviton production during reheating implies $T_R \\alt 1$ to $10^2~\\tev$, depending, in part, on the number of extra dimensions~\\cite{Feng:2003nr}.  On the other hand, the Universe has proven to be remarkably baroque, and there is no guarantee that dark matter is composed of only one component.  One might therefore relax the constraint of \\eqref{omega} on the superWIMP energy density and assume, for example, that the NLSP freezes out with its thermal relic density $\\Omega_{\\text{NLSP}}^{\\text{th}}$.  The superWIMP gravitino density is then $\\OmegaSWIMP = (m_{\\gravitino}/m_{\\text{NLSP}}) \\Omega_{\\text{NLSP}}^{\\text{th}}$.  In this approach, the gravitino density may be low and even insignficant cosmologically.  {}From a particle physics viewpoint, however, the viability of the gravitino LSP scenario is still worth investigation, as it has strong implications for the superpartner spectrum and collider signatures, independent of its cosmological importance.  These approaches differ significantly, not only in the viewpoint taken, but also in their implications.  Suppose, for example, that the gravitino and NLSP masses are both parametrized by a mass scale $m_{\\text{SUSY}}$.  The NLSP number density then scales as $1/m_{\\text{SUSY}}$ if one assumes \\eqref{omega}, but scales as $m_{\\text{SUSY}}$ if one assumes a thermal relic NLSP density, since $\\Omega_{\\text{NLSP}}^{\\text{th}} \\propto \\langle \\sigma v \\rangle ^{-1} \\propto m_{\\text{SUSY}}^2$, where $\\langle \\sigma v \\rangle$ is the thermally-averaged NLSP annihilation cross section.  Low masses are excluded in the former case, while high masses are disfavored in the latter.  We consider both approaches to be worthwhile, but defer discussion of the thermal relic density approach, along with its very different consequences for the superpartner spectrum and implications for collider physics, to a separate study~\\cite{Feng:2004}.  Assuming a fixed $\\OmegaSWIMP$ here, we find that gravitino superWIMPs provide a viable solution to the dark matter problem.  We determine the allowed masses for the NLSP and gravitino.  At the same time, we find that the hadronic energy constraint is significant and does in fact provide the leading constraint in natural regions of parameter space. We conclude that the analysis below must be done in any scenario with similar late decays.  Given any standard cosmology, a significant thermal relic abundance of NLSPs will be generated, these NLSPs will eventually decay to superWIMPs, and the hadronic constraints we discuss below must be analyzed before the scenario may be considered viable.  In passing, we note that another possibility is that the NLSP is the lightest neutralino.  For a general neutralino, the decays $\\chi \\to Z \\tilde{G}, h \\tilde{G}$ produce large hadronic energy release.  These two-body modes are absent for photino-like neutralinos, where the only two-body decay is $\\tilde{\\gamma}\\to\\gamma \\tilde{G}$. However, even in this case, one cannot avoid hadronic activity from decay through a virtual photon, $\\tilde{\\gamma} \\to \\gamma^* \\tilde{G} \\to q\\bar{q} \\tilde{G}$, with a branching fraction of ${\\cal O} (10^{-2})$. Given the stringency of constraints on hadronic energy release, the neutralino NLSP case is highly constrained.  Assuming that $\\OmegaSWIMP$ makes up a significant fraction of the dark matter, natural regions of parameter space are excluded~\\cite{Feng:2004}.  In the current paper, we therefore focus on the more promising possibilities of slepton and sneutrino NLSPs.  The paper is organized as follows.  In Sec.~\\ref{sec:bbn}, we summarize what is known about constraints on EM and hadronic energy release in the early Universe.  In Sec.~\\ref{sec:swimp}, we give a detailed description of the gravitino superWIMP scenario with a slepton or sneutrino NLSP.  In Sec.~\\ref{sec:had}, we determine the hadronic energy release from slepton and sneutrino decays and find the allowed mass regions.  We summarize in Sec.~\\ref{sec:conclusion}.  The relevant gravitino interactions, decay amplitudes, and three-body hadronic decay width formulae are collected in the Appendix. ",
        "conclusions": "\\label{sec:conclusion}  In this paper, we explored the possibility of superWIMPs as candidates for dark matter.  Examples of superWIMPs are the gravitino in supersymmetric models and the lightest Kaluza-Klein graviton in universal extra dimension models.  SuperWIMPs obtain their relic density through late decays: WIMP $\\to$ superWIMP $+$ SM particles.  The decay usually occurs between $10^4$ s and $10^8$ s, which has important cosmological implications.  Such late decays release EM and hadronic energy into the Universe, which may affect BBN predictions for the primordial abundance of the light elements.  The constraints on EM injection have been studied in detail in Ref.~\\cite{Feng:2003uy}.  In this paper, we have analyzed the hadronic BBN constraints.  We have taken the lightest gravitino as a concrete example of a superWIMP and have focused on slepton and sneutrino NLSPs as the most promising WIMP candidates.  We have determined the hadronic energy release by calculating three-body decay widths in detail.  For the cases of $\\tilde{\\tau}_R$ and $\\tilde{\\nu}$ NLSPs, the hadronic decay branching fraction is below the level of $10^{-3}$ when $m_{\\tilde{\\tau}_R} \\alt 1~\\tev$ and $m_{\\tilde{\\nu}} \\alt 300~\\gev$, respectively, or when the decay lifetime is $\\tau \\agt 10^6~\\s$.  We identified the allowed and disfavored regions of the $(m_{\\tilde{G}}, \\delta m)$ plane, imposing CMB constraints, EM BBN constraints, and hadronic BBN constraints. For the sneutrino NLSP case, $\\delta m \\alt 60 - 200~\\gev$ is allowed for a large range of $m_{\\tilde{G}}$.  For the stau NLSP, the hadronic constraints are weaker: $\\delta m \\alt 200-1200~\\gev$.  However, additional constraints from CMB and EM energy injection apply. Combining all the constraints for the right-handed stau NLSP, the allowed window is $300~\\gev \\alt m_{\\tilde{G}} \\alt 1~\\tev$, $200~\\gev \\alt \\delta m \\alt 1200~\\gev$, corresponding to $m_{\\tilde{\\tau}_R} \\agt 500~\\gev$.  We have assumed $\\OmegaSWIMP = 0.23$ in our analysis.  The constraints would be relaxed if superWIMPs are only part of the present dark matter.  In addition, there are still ambiguities in the BBN constraints if the decay time lies in the region where EM and hadronic effects are comparable.  In particular, as we discussed, their effects on D might cancel, and the allowed parameter space could be enlarged.  Progress in firming up BBN constraints in this region, and for other elements, such as $^3$He and $^6$Li may have crucial implications for the superWIMP dark matter scenario.  Although the superWIMP itself would escape all direct and indirect dark matter searches, the slepton/sneutrino NLSP will have rich implications for collider phenomenology.  Such metastable sleptons will not decay inside the detector, resulting in signals of highly ionizing tracks for sleptons and missing energy signals for sneutrinos.  Discussion of the collider phenomenology, combined with thermal relic density calculations for the NLSP in supergravity models will be presented in Ref.~\\cite{Feng:2004}.  Although we have focused on the SUSY scenario in this paper, superWIMPs could in general be any gravitationally interacting particle that obtains its relic density through the late decay of a WIMP.  The discussion and results for late decay of a leptonic WIMP in other models will be qualitatively similar to the discussion and results presented here.  This scenario also suggests that a particle that appears to play the role of dark matter at late times, even after BBN, could very well be different from the particle that constitutes dark matter now.  The very late decay of ``would be'' dark matter particles may have important cosmological implications, for example, affecting small scale structure.  This feature may provide qualitatively new possibilities for explaining puzzling cosmological observations~\\cite{Chen:2003gz,Sigurdson:2003vy}.  We have assumed $R$-parity conservation in our discussion.  In the case of $R$-parity violation (RPV), the gravitino could still constitute dark matter as long as its decay lifetime is comparable to or longer than the age of Universe~\\cite{Takayama:2000uz,Moreau:2001sr}.  Gravitino dark matter in an RPV scenario could be distinguished from the superWIMP scenario by both cosmological observations and collider experiments.  For example, the decay of even a tiny amount of gravitino dark matter into SM particles in an RPV scenario could be seen in the diffuse photon flux.  In addition, if the RPV is not extremely small, the collider signatures could be different from those in the superWIMP scenario. Further work is needed to study how to distinguish these two scenarios."
    },
    "0404/astro-ph0404265_arXiv.txt": {
        "abstract": "Five new objects with proper motions between 1.0\\arcsec/yr and 2.6\\arcsec/yr have been discovered via a new RECONS search for high proper motion stars utilizing the SuperCOSMOS Sky Survey.  The first portion of the search, discussed here, is centered on the south celestial pole and covers declinations $-$90$^\\circ$ to $-$57.5$^\\circ$.  Photographic photometry from SuperCOSMOS and $JHKs$ near-infrared photometry from 2MASS for stars nearer than 10 pc are combined to provide a suite of new M$_{Ks}$-color relations useful for estimating distances to main sequence stars.  These relations are then used to derive distances to the new proper motion objects as well as previously known stars with $\\mu$ $\\ge$ 1.0\\arcsec/yr (many of which have no trigonometric parallaxes) recovered during this phase of the survey.  Four of the five new stars have red dwarf colors, while one is a nearby white dwarf.  Two of the red dwarfs are likely to be within the RECONS 10 pc sample, and the white dwarf probably lies between 15 and 25 pc.  Among the 23 known stars recovered during the search, there are three additional candidates for the RECONS sample that have no trigonometric parallaxes. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/gr-qc0404014_arXiv.txt": {
        "abstract": "Using numerical techniques, we study the collapse of a scalar field configuration in the Newtonian limit of the spherically symmetric Einstein--Klein--Gordon (EKG) system, which results in the so called Schr\\\"odinger--Newton (SN) set of equations. We present the numerical code developed to evolve the SN system and topics related, like equilibrium configurations and boundary conditions. Also, we analyze the evolution of different initial configurations and the physical quantities associated to them. In particular, we readdress the issue of the gravitational cooling mechanism for Newtonian systems and find that all systems settle down onto a $0$--node equilibrium configuration. ",
        "introduction": "\\label{sec:introduction} In a previous paper of ours\\cite{fsglau}, we studied the formation of a gravitationally bounded object comprised of scalar particles, under the influence of Newtonian gravity. The dynamics of the system is described by the coupled Schr\\\"odinger--Newton (SN) system of equations, which is nothing but the weak field limit of its general relativistic counterpart, the Einstein--Klein--Gordon (EKG) system.  As at the moment we have no hints to finding an analytic solution for the evolution, we found necessary to develop numerical techniques to study the formation process of the scalar objects. The study of the dynamical properties of the fully time-dependent SN system has been done before in the literature\\cite{seidel90,seidel94,schmethod,hu}, but more is needed in order to understand the gravitational collapse of a weakly gravitating scalar field.  The main aim of this paper is to perform an exhaustive numerical study of the collapse and evolution of a spherically symmetric scalar object in the Newtonian regime. Here, we develop a numerical strategy to evolve the SN system, and study important issues like the stability and the formation process of gravitationally bound scalar systems, a topic that has recently become attractive in Cosmology \\cite{fsglau,schmethod,hu,phi2,dmota,jcerv,fuchs}.  A summary of the paper is as follows. In Sec.~\\ref{mathematics}, we present the relativistic EKG and Newtonian SN equations that describe the evolution of a self-gravitating scalar field in the spherically symmetric case. Correspondingly, it is described how the EKG and the SN equations are solved in order to obtain regular and asymptotically flat solutions. These solutions are known as \\textit{boson stars} and \\textit{oscillatons}, respectively, for complex and real scalar fields. However, we focus our attention on their corresponding weak field limit, SN system and its properties.  In Sec.~\\ref{sec:numerics}, we present an appropriate numerical code to evolve the SN system, providing details about the algorithms used. The issues concerning the physical boundary conditions imposed on the SN system and the accuracy of the code are of particular importance.  The results of the numerical evolutions are given in Sec.~\\ref{results}. We systematically test the boundary conditions, the convergence of the numerical solutions and how the latter reproduces the expected results for the equilibrium configurations of the SN system.  Sec.~\\ref{sec:cooling} is devoted to the study of the gravitational cooling mechanism, first described in\\cite{seidel94}. Finally, the conclusions are collected in Sec.~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions} In this paper, we have studied the evolution of a self-gravitating scalar field in the Newtonian regime using numerical methods. The latter were systematically tested for their accuracy, convergence and reliability. In all cases, the numerical code gave the expected correct results of the $0$--node equilibrium configurations and its perturbations. Also, the numerical code preserved the scaling invariance of the SN system all along the time of the evolutions.  An important point was the boundary condition imposed on the SN system. We found that the implementation of a sponge by adding an imaginary potential is an appropriate boundary condition. It allowed us to maintain under control the amount of scalar matter reflected by the numerical boundary.  Our results imply that $0$--node equilibrium configurations are intrinsically stable against radial perturbations, and that they play the important role of final states arbitrary scalar systems settle down onto in the Newtonian regime. On the contrary, excited equilibrium configurations were found to be intrinsically unstable configurations. Even though they are initially virialized, they evolve towards a $0$-node solution.  An important point here is that, due to the scaling properties of the SN system, the study of the whole space of possible equilibrium configurations was reduced to the analysis of some properly sized configurations. Moreover, the same was done to study non-equilibrium configurations to give simplified and accurate simulations.  Last but not least, we retook the analysis of the gravitational cooling within the Newtonian regime of scalar fields. The main result is that, in the weak field limit, the gravitational cooling is a very efficient mechanism, which allows any initial configurations to decay into a $0$--node Newtonian scalar soliton. So far, we have not found any evidence for systems that disperse away by ejecting all their mass.  We expect that the results presented here would provide useful information about structure formation in the universe, not only regarding models of scalar field dark matter as in\\cite{fsglau,sin,hu,arbey,gal}, but also about other models whose dynamics is governed by the SN system beyond spherical symmetry as in\\cite{schmethod}, case about which we expect to report in the near future."
    },
    "0404/astro-ph0404053_arXiv.txt": {
        "abstract": "When the polarimetric sensitivity and the angular resolution exceed a threshold, magnetic fields show up almost everywhere on the solar surface. Here I revise the observational properties of the weakest polarization signals, which appear in the InterNetwork (IN) regions. We already have some information on the magnetic field strengths and inclinations, mass motions, lifetimes, magnetic fluxes, magnetic energies, etc. Since the IN covers a substantial faction of the solar surface, it may account for most of the unsigned magnetic flux and energy existing on the solar surface at any given time. This fact makes IN fields potentially important to understand the global magnetic properties of the Sun (e.g., the structure of the quiet solar corona, an issue briefly addressed here). The spectropolarimeters on board of Solar-B have the resolution and sensitivity to routinely detect these IN fields. ",
        "introduction": "Most of the solar surface appears non-magnetic when it is observed in routine synoptic magnetograms. However, magnetic fields are detected almost everywhere, also in InterNetwork (IN) regions, when the polarimetric sensitivity and the angular resolution are high enough. These magnetic fields are now accessible to many spectropolarimeters. They cover much of the solar surface and, therefore, they may account for most of the unsigned magnetic flux and energy existing on the solar surface at any given time. The contribution summarizes the main observational properties of the IN fields, as deduced from these recent measurements. In addition, it discusses the origin of the IN magnetism, and why it is a natural target for Solar-B. ",
        "conclusions": "The quiet Sun is the component of the solar surface magnetism that seems to account for most of the magnetic flux and magnetic energy (\\S~\\ref{energy}). This fact makes it potentially important to understand the global magnetic properties of the Sun (solar dynamo, coronal heating, origin of the solar wind, and so on). However, its influence has been neglected so far. It produces very weak polarization signals which hardly show up in conventional magnetograms. The situation is slowly turning around,  and it will dramatically change with the advent of space-borne polarimeters like those of Solar-B. For example, Fe~{\\sc i}~6302~\\AA\\ is expected to produce a circular polarization of the order of $5\\times 10^{-3}$ when the angular resolution reaches 0\\farcs 5 (\\citealt{dom03b}). These signals are well above the noise level in the normal mapping model of the SOT spectropolarimeter (\\citealt{shi04}). In other words, Solar-B is expected to routinely detect the quiet Sun magnetic fields (at least that fraction having kG field strengths, see \\S~\\ref{mfs}).    The true role of the quiet Sun magnetism is still unknown. Only preliminary steps to figure it out have been given. Let me point out a recent work by \\citet{sch03b} where they study the influence of the quiet Sun magnetic fields on the extrapolation of the photospheric field to the corona. They conclude that an important modification of the network-rooted field lines is induced by the presence of the IN, implying that a significant part of this disorganized IN photospheric field does indeed reaches the quiet corona."
    },
    "0404/astro-ph0404579_arXiv.txt": {
        "abstract": "A simple and general description  of the dynamics of a narrow-eccentric ring is presented. We view an eccentric ring which precesses uniformly at a slow rate as exhibiting a global $m=1$ mode, which can be seen as originating from a standing wave superposed on an axisymmetric background.  We  adopt a continuum description using the language of fluid dynamics which gives equivalent results for the secular dynamics of thin rings  as the the well known description in terms of a set of discrete elliptical streamlines formulated by Goldreich and Tremaine (1979). We use this to discuss the non linear mode interactions that  appear in the ring  through  the excitation of higher $m$ modes because of the coupling of the $m=1$ mode with an external satellite potential, showing  that they   can lead to the excitation of the $m=1$ mode through a feedback process. In addition to the external perturbations by neighboring satellites, our model includes effects due to inelastic inter-particle collisions. Two main conditions for the ring to be able to maintain a steady $m=1$ normal mode are obtained. One can be expressed as an integral condition for the normal mode pattern to precess uniformly, which requires the correct balance between the differential precession induced by the oblateness of the central planet, self-gravity and collisional effects is the continuum form of that obtained from the $N$ streamline model of Goldreich and Tremaine (1979). The other condition, not before examined in detail, is for the steady maintenance of the non-zero radial action that the ring   contains because of its finite normal mode.  This requires a balance   between injection due to eccentric resonances arising from external satellites and additional collisional damping associated with the presence of the $m=1$ mode.  We estimate that such a balance  can occur in the $\\epsilon-$ring of Uranus, given its currently observed physical and orbital parameters. ",
        "introduction": "The  nature of the dynamical mechanism that maintains the apse alignment of narrow-eccentric planetary rings is one of the most interesting and challenging problems of Celestial Mechanics.  According to  the leading model (Goldreich and Tremaine 1979) the self-gravity of the ring counter-acts the differential precession induced by the oblateness of the central planet. Using this hypothesis, a prediction of the total mass of the ring can be made, which, in general, is not in good agreement with the inferred mass of the observed eccentric rings in the Uranus system (Tyler\\etal 1986, Graps\\etal 1995, Goldreich and Porco 1987, Porco and Goldreich 1987). This led to the consideration of other factors that  might play an important role in the dynamics. In particular, at their narrowest point, the ring particles are `close-packed'.  In such a situation particle interaction or pressure effects may affect the precession of particle orbits. A simple model where the {\\it pinch} locks the differential precession, was introduced by Dermott and Murray (1980). A more global picture, including the effect of  stresses due to particle interactions and neighboring satellite perturbations, which offered a better agreement with the observations, has been produced by Borderies\\etal (1983). Their dynamical model is described in terms of mutually interacting {\\it streamlines} and the satellite interactions (see Goldreich  and Tremaine 1981) are computed using a resonance-continuum approximation. The standard self-gravity model was later revisited  by Chiang and Goldreich (2000), who  considered the effects of collisions near the edges, proposing that a sharp increase of an order of magnitude in the surface density should be observed within the last few hundred meters of the ring  edges. More recently, employing a pressure term that describes close-packing, Mosqueira and Estrada (2002) obtained surface-density solutions that agree well with the currently available mass estimates.  However, several questions remain, such as how steady global $m=1$ modes are maintained by intercations with a satellite and mode couplings against dissipative processes. In particular if such modes are associated with  particle close packing at certain orbital phases, this is likely to produce enhanced dissipation that has to be made up through the action of satellite torques. Here we consider the issue as to how the torque rate input available from satellites can counteract such collisional dissipation.  In this work we build, from first principles, a simple general continuum or fluid like model of a narrow-eccentric ring. The eccentric pattern in the ring can be described as being generated by a normal mode of oscillation of wave-number $m=1$ which may be  considered to be a standing wave. Dissipation can be allowed to occur due to inter-particle collisions leading to a viscosity which would lead to damping of the mode. However, this global $m=1$ mode can  also be perturbed by neighboring-shepherd satellites which can inject energy and angular momentum through resonances. In this way losses due to particle collisions may be balanced. It is that process that is the focus of this paper. Other possible mechanisms, such as  mode excitation through self-excitation through viscous overstability, that could arise with  an appropriate dependence of viscosity on physical state variables (see Papaloizou and Lin 1988, Longaretti \\& Rappaport 1995), are beyond the scope of this paper and accordingly  not investigated here.  To describe the ring perturbations and   the $m=1$ mode we use the Lagrangian-displacement of the particle orbits from their unperturbed circular  ones (for a similar treatment applied to very diverse problems see for example Lebovitz 1967, Lynden-Bell and Ostriker 1967, Friedman and Schutz 1978, Shu\\etal 1985).  In section~\\ref{LD} we set up the equations for the Lagrangian variations starting from the  equations of motion in a 2D flat disk approximation. We also compute the Lagrangian variation of the satellite potential as seen by a particle as a consequence of the existence of the $m=1$ mode, and we discuss what non linear couplings appear in the ring as a result including the excitation of the $m=1$ mode through a feedback process.  In section~\\ref{m1m} we derive the radial equation of motion for the  2D Lagrangian displacement under the assumption that the ring is primarily in a $m=1$ normal mode. The definition of the radial action in terms of the Lagrangian displacement for small eccentricities is given in section~\\ref{RA}, and its rate of change is obtained.  We can then determine a condition for the steady  maintenance of the amplitude  or eccentricity  associated with the $m=1$ mode, which requires the external satellite input to balance the dissipative effects due to collisions.  \\noindent In this paper, the application we consider is to the $\\epsilon$ ring of Uranus. As there are no effective  corotation resonances in this ring if the eccentricity of the perturbing satellites is neglected as is done here,  we shall consider only Lindblad resonances and  defer consideration of corotation resonances to future work.  The satellite torque is obtained in section~\\ref{SATT}. In section~\\ref{RAST} we show that for the narrow rings considered here, the satellite contribution to the rate of change of the radial action is just a fraction of the corresponding satellite torque dependent only on the relevant azimuthal mode number, $m.$  The additional condition for the existence of the normal mode, i.e. the condition of uniform-precession, is derived in section~\\ref{UP}. The self-gravity term appearing in this condition is computed in section~\\ref{SGterm}. In section~\\ref{vq} we show that the eccentricity gradient is necessarily positive in a narrow-ring which in which  uniform precession is mainly maintained by self-gravity and we estimate its value in the linear regime.  In section~\\ref{discu} we discuss our results and by  considering the total ring radial action using a very simple N-body approach, we illustrate the global functioning of a narrow-eccentric ring. We are able to obtain  both ring spreading and the balance of satellite torques and  collisional dissipation required to maintain ring eccentricity in various limiting cases of ring  evolution. Finally, we apply our results to the $\\epsilon-$ring of Uranus estimating that the balance between the satellite torque input and the dissipative effects necessary to maintain its eccentricity can be established in this case. ",
        "conclusions": ""
    },
    "0404/astro-ph0404323_arXiv.txt": {
        "abstract": "In order to understand the nature of magnetic reconnection in ``free space'' which is free from any influence of external circumstances, I have studied the structure of spontaneous reconnection outflow using a shock tube approximation. The reconnection system of this case continues to expand self-similarly. This work aims 1) to solve the structure of reconnection outflow and 2) to clarify the determination mechanism of reconnection rate of the ``self-similar evolution model'' of fast reconnection.  Many cases of reconnection in astrophysical phenomena are characterized by a huge dynamic range of expansion in size ($\\sim 10^7$ for typical solar flares). Although such reconnection is intrinsically time dependent, a specialized model underlying the situation has not been established yet.  The theoretical contribution of this paper is in obtaining a solution for outflow structure which is absent in our previous papers proposing the above new model. The outflow has a shock tube-like structure, i.e., forward slow shock, reverse fast shock and contact discontinuity between them. By solving the structure in a sufficiently wide range of plasma-$\\beta$: $0.001 \\leq \\beta \\leq 100$, we obtain an almost constant reconnection rate ($\\sim 0.05$: this value is the maximum for spontaneous reconnection and is consistent with previous models) and a boundary value along the edge of the outflow (good agreement with our simulation result) which is important to solve the inflow region. Note that everything, including the reconnection rate, is spontaneously determined by the reconnection system itself in our model. ",
        "introduction": "\\label{sec:Int}  It is widely accepted that magnetic reconnection very commonly plays an important role as a powerful energy converter in astrophysical plasma systems. However, there still remain many open questions, not only regarding the microscopic physics of the anomalous resistivity, but also the macroscopic magnetohydrodynamical (MHD) structure. In this work, our attention is focused on the macroscopic evolutionary properties of magnetic reconnection, especially on the structure of the reconnection outflow.   We consider reconnection in an anti-parallel magnetic configuration called a ``current sheet system''. In this system, two similarly uniform magnetized regions are set in contact, divided by a boundary. We assume that the directions of the magnetic field on both sides are anti-parallel to clarify the essence of the problem. In this case, the boundary carries a strong electric current, hence, we call it the current sheet. If resistivity is enhanced in this current sheet, it plays an important role in energy conversion from the magnetic form to others. In resistive plasmas, there are two fundamental processes of magnetic energy conversion. One is magnetic diffusion, and the other is magnetic reconnection.  The best known process of energy conversion in resistive plasmas from magnetic to other forms is Ohmic dissipation (magnetic diffusion). We must note, however, that plasmas are highly conductive in most astrophysical problems. Since the resistivity is very small, energy conversion by global magnetic diffusion takes a very long time, and is not applicable to many astrophysical phenomena with very violent energy releases, e.g., solar flares. Even in such a case, the majority believe that magnetic reconnection can convert the magnetic energy very quickly. This is why we must study magnetic reconnection.  Such energy conversion in the current sheet system is very important in many astrophysical plasma systems. The most famous example is the relation between solar activity and geomagnetospheric activity, called solar flares and geomagnetospheric substorms. In this case, at least three current sheet systems are included. The first one is in the solar corona. The second and third are on the day-side and the night-side of the geomagnetosphere, respectively. Recent observations and numerical simulations suggest that similar phenomena, seemingly activated by magnetic reconnection, are universal, e.g., flares in accretion disks of young stellar objects (YSOs, see Koyama et al. 1994, Hayashi et al. 1999), galactic ridge X-ray emissions (GRXE, see Koyama et al. 1986, Tanuma et al. 1999).  We must note that many cases of actual magnetic reconnection in astrophysical systems usually grow over a huge dynamic range in their spatial dimensions. For example, the initial scale of the reconnection system can be defined by the initial current sheet thickness, but this is too small to be observed in typical solar flares. We do not have any convincing estimate of the scale, but if we estimate it to be of the order of the ion Larmor radius, it is extremely small ($\\sim 10^0$ [m] in the solar corona). Finally, the reconnection system develops to a scale of the order of the initial curvature radius of the magnetic field lines ($\\sim 10^7$ [m] $\\sim$ 1.5\\% of the solar radius for typical solar flares). The dynamic range of the spatial scale is obviously huge ($\\sim 10^7$ for solar flares). For geomagnetospheric substorms, their dynamic range of growth is also large ($\\sim 10^4$ for substorms). Such a very wide dynamic range of growth suggests that the evolution of the magnetic reconnection should be treated as a development in ``free space'', and that external circumstances do not affect the evolutionary process of magnetic reconnection, at least at the expanding stage just after the onset of reconnection.  Even if a system is not completely free from the influence of external circumstances, we can approximately treat it as a spontaneously evolving system if the evolution timescale (which is estimated as Alfv\\'{e}n transit time $\\tau_A$ for the final system scale because the spontaneous expansion speed of the reconnection system is equivalent to the fast-mode propagation speed, see Nitta et al. 2001; hereafter paper 1) is much smaller than the timescale imposed by the external circumstances (e.g., convection timescale). For typical substorms, the evolution timescale of the reconnection is of the order of seconds or minutes while the convection timescale to compress the current sheet is of the order of hours. In such cases, the external circumstances simply play the role of triggering the onset of reconnection and the evolution itself is approximately free from the influence of the external circumstances. Though, of course, exceptional cases in which the external circumstances intrinsically influence the evolution can arise in particular situations (e.g., cases in which very fast convection drives the reconnection), the author believes that it is worth establishing a new model applicable to the cases which are free from external influences.   However, no reconnection model evolving in a free space has been studied. We should note that most previous theoretical and numerical works on reconnection treated it as a boundary problem strongly influenced by external circumstances.  In actual numerical studies, there is a serious and inevitable difficulty: For magnetic reconnection to be properly studied, the thickness of the current sheet must be sufficiently resolved by simulation mesh size. Usually, the thickness of the current sheet is much smaller than the entire system. Hence, in previous works, we have been forced to cut a finite volume out of actual large-scale reconnection systems for numerical studies.  Of course, we wish to realize that the boundary conditions of this type of finite simulation box reproduce the evolution in unbounded space. However, we should note that, in actual simulations, the boundary of such a simulation box necessarily affects the evolution inside the box, even if we use so-called ``free'' boundary conditions. This has an adverse influence on the evolutionary process of the numerically simulated magnetic reconnection: When the reconnection proceeds and physical signals propagating from the inner region cross the boundaries of the simulation box, the subsequent evolution is necessarily affected by the boundary conditions. This is because the information propagated through the boundary is completely lost and such an artificially cut-out simulation box will never receive the proper response of the outer regions. Additionally, numerical and unphysical signals emitted from the boundary may disturb the evolution. Such artificially affected evolution is obviously unnatural, and the resultant stationary state should differ from actual reconnection occurring in a free space.  Even in previous numerical studies aimed at clarification of the time evolution of reconnection (for example, series of studies originated by Ugai \\& Tsuda 1977 or Sato \\& Hayashi 1979), the evolution could not be followed for a long time. This is mainly owing to restriction of the size of the simulation box, hence, application of the results was limited to spatial scales typically, say, a hundred times the spatial scale of the diffusion region. We are interested here in an evolutionary process in free space without any influence of external circumstances. In such a system, the evolution and resultant structure would be quite different from these previous numerical models.  The same problem also appears in previous theoretical works. We must realize that, in many cases, astrophysical magnetic reconnection is essentially a non-stationary process because it grows in a huge spatial dynamic range. Most of these works, however, for mathematical simplicity, treat a stationary state of the reconnection in a finite volume (for example, see Priest \\& Forbes 1986). These are obviously boundary problems and solutions must be influenced by boundary conditions. The problem is that we do not know how we should set the boundary condition in order to simulate external influences. In general, it is impossible to set. Hence, the situations argued in previous works are rather artificial and unnatural in many cases of astrophysical application.  As will be discussed later (see section \\ref{sec:stat}), if we cut the central region out from the entire expanding system, the central region tends to be a stationary state and is very similar to the well-known Petschek model, i.e., the system is characterized by the fast-mode rarefaction dominated inflow and the figure-X-shaped slow shock. The author thinks that stationary solutions are meaningful as the interior solution of the entire evolving system in free-space though such stationary models are frequently treated as externally ``driven'' processes by boundary conditions.  We demonstrated the above process of self-similar evolution model of fast reconnection by numerical simulation (see paper 1) and semi-analytic study (see Nitta et al. 2002; hereafter paper 2). Let us overview such an evolutionary process in free-space.  We suppose a two-dimensional equilibrium state with an anti-parallel magnetic field distribution, as in the Harris solution. When magnetic diffusion takes place in the current sheet by some localized resistivity, magnetic reconnection will occur, and a pair of reconnection jets is ejected along the current sheet. This causes a decrease in total pressure near the reconnection point. Such information propagates outward as a rarefaction wave. In a low-$\\beta$ plasma ($\\beta \\ll 1$ in the region very distant from the current sheet [asymptotic region]; as typically encountered in astrophysical problems), the propagation speed of the fast-magnetosonic wave is isotropic, and is much larger than that of other wave modes. Thus, information about the decreasing total pressure propagates almost isotropically as a fast-mode rarefaction wave (hereafter FRW) with a speed almost equal to the Alfv\\'{e}n speed $V_{A0}$ in the asymptotic region. Hence, the wave front of the FRW (hereafter FRWF) has a cylindrical shape except near the point where the FRWF intersects with the current sheet. When the FRWF sufficiently expands, the initial thickness of the current sheet becomes negligible compared with the system size $V_{A0} t$, where $t$ is the time from the onset of reconnection. In such a case, there is only one characteristic scale, i.e., the radius of the FRWF ($V_{A0} t$), which linearly increases as time proceeds. This is just the condition for self-similar growth.   In our semi-analytic work (paper 2), the boundary condition along the edge of the outflow is artificially imposed approximating the result of our numerical simulation in paper 1. We set the boundary condition at $y=0$ for simplicity because the outflow is very thin. This boundary condition denotes the junction condition to the reconnection outflow. Hence, an important question ``how is this boundary condition spontaneously determined by the reconnection system'' still remains. Our motivation of this work is to clarify the mechanism to determine this boundary condition.  The boundary condition at $y=0$ determines the inflow region. Hence, the reconnection rate which is usually denoted by the normalized inflow speed toward the diffusion region is also determined. We must note that, in our self-similar evolution model, everything including the reconnection rate is spontaneously determined by the reconnection system itself. We try to clarify how the system regulates the reconnection.  This paper is organized as follows. We state a model of reconnection outflow as a kind of shock tube problem in section \\ref{sec:model}. Basic equations for outflow structure and the numerical procedure to solve the equations are listed in section \\ref {sec:b-eqs}. Properties of the result, e.g., the reconnection rate and boundary condition along the slow shock are shown in section \\ref{eq:res}. In section \\ref{sec:sum-dis}, we summarize our study and discuss spontaneous structure formation in the reconnection outflow. ",
        "conclusions": "\\label{sec:sum-dis}  \\subsection{Summary} \\label{sec:sum}  We here briefly summarize the structure of the reconnection system obtained in our numerical simulation in paper 1. The feature of the inner region is very similar to fast reconnection regimes (Priest \\& Forbes 1986) of which the Petschek model is one member, i.e., we can find the fast-mode rarefaction dominated inflow and X-shaped slow shock. However, the entire structure shows properties of time-dependent reconnection as follows. The reconnection outflow has a finite span which is increasing in proportion to the time from the onset. A V-shaped forward slow shock forms in the vicinity of the spearhead of the outflow. In region 3, plasmas are compressed by the reconnection jet (piston effect), and squeezed plasmas are gushing out to make a vertical outflow. These are properly included in evolutionary process.  In this paper, the structure of reconnection outflow in the self-similar evolution model has been studied in a shock tube approximation. We have obtained the reconnection rate and the required boundary condition to solve the inflow region from the structure of the reconnection outflow.    \\subsection{Self-Similarity of the reconnection outflow}  We have solved the structure of the reconnection outflow for the self-similar evolution model of fast magnetic reconnection. The structure is spontaneously determined by the reconnection system itself as approximated by a kind of shock tube problem. The reconnection outflow is divided into several regions by discontinuities, e.g., forward shock, reverse shock, and contact discontinuity. We solved this shock tube problem by an iterative (Newton-Raphson) method. The moving speeds of these discontinuities are constant in time in our result. Thus, as expected from the well-known shock tube problem, the reconnection outflow extends self-similarly.  First, we argue for self-similarity of the evolution of the reconnection system in free space. In the discussion of  paper 2, we need an {\\it a priori} statement to authorize a self-similar solution for the inflow region so that the reconnection outflow (in other words, the boundary condition along the slow shock) also evolves self-similarly. The validity of this assumption was, however, unclear in previous papers. We must note that because we have found the self-similarly evolving solution for the reconnection outflow, we can ensure self-similar evolution of the entire reconnection system.  Comparison of the outflow structure with a shock tube model is also argued in Abe \\& Hoshino (2001). They obtained a similar structure to our result, but there is a critical difference. The V-shaped discontinuity of the forward shock is identified as an intermediate shock in their result instead of as a slow shock in our case. The author cannot explain the reason for the difference, but notes that their attention is focused on a much earlier stage than our case. The scale of the plasmoid is the order of the current sheet thickness in their case, but we are interested in a much larger scale structure at the self-similar stage. Actually, the numerical result of Ugai 1999 (see figure 9 of that paper) shows a similar structure to our result. The later evolution of the simulation of Abe \\& Hoshino (2001) tends also to coincide with our result.    \\subsection{$\\beta$-dependence of the reconnection rate}  Second, our attention is focused on the $\\beta$-dependence of the reconnection rate $R$. In our result, $R$ is almost constant in the range $0.001 \\leq \\beta \\leq 100$ (see figure \\ref{fig:rec-rate}). We could foresee this behavior of the reconnection rate from the following intuition.  We consider the low $\\beta$ limit in the asymptotic region. When the plasma $\\beta$ at the asymptotic region far outside the current sheet becomes smaller, the Petschek type X-slow shock must be stronger. This is explained as follows. The gas pressure in the up-stream region is very small relative to the magnetic pressure for a low-$\\beta$ plasma, while in the down-stream region, the gas pressure should be very high to keep total pressure equilibrium with the up-stream region. Hence, the total pressure equilibrium imposes a large gas pressure jump at the X-slow shock. This leads to a large jump in the mass density.  Next we consider the high $\\beta$ limit in the asymptotic region. Both sides of the Petschek type X-slow shock are filled with gas-pressure dominated plasmas in this case, and balanced with each other mainly by gas pressure. This results in a very small density jump at the slow shock ($X \\rightarrow 1$). Note that this does not mean weak shock limit because the jump of the tangential component of velocity is significant ($v_1 \\sim 1$ while $v_{xp}=0$). The $\\beta$-dependence of the compression ratio $X$ of the X-slow shock is shown in figure \\ref{fig:comp} and is consistent with the above discussion.  The reconnection rate $R$ is defined as $R \\equiv -v_{yp} B_{xp}$. First we discuss the low $\\beta$ limit. The inflow speed $-v_{yp}$ will increase as the slow shock strengthens ($\\beta$ decreases) for the following reason. The ejection speed of the reconnection jet is almost constant ($\\sim V_{A0}$) while the mass density of the jet increases as the slow shock strengthens (the compression ratio $X$ increases as $\\beta$ decreases). Thus, the inhalation of mass flux toward the slow shock is also strengthened. This results in an increase of $-v_{yp}$ (see figure \\ref{fig:vin-Bin}) as $\\beta$ decreases because the mass density at the pre-slow shock region is almost constant (though weakly rarefied by the fast-mode rarefaction wave). As the inflow speed $-v_{yp}$ increases, bending of the magnetic field line becomes remarkable, and $B_{xp}$ decreases (see figure \\ref{fig:vin-Bin}). We can also understand the result for the high $\\beta$ limit in the same way ({\\it vice versa} for the high $\\beta$ limit). Consequently, the product of $-v_{yp}$ and $B_{xp}$, and hence the reconnection rate $R$ is almost constant independent of $\\beta$.  Dependence of the reconnection rate on the plasma-$\\beta$ is also considered in other works, e.g., Ugai \\& Kondoh 2001. Their attention is focused on an implicit contribution of $\\beta$ via a current-driven anomalous resistivity model. The aim of our study is intrinsically different from theirs. Needless to say, anomalous resistivity is very important to confirm the validity of a localization of resistivity which is necessary to establish fast reconnection (see Yokoyama \\& Shibata 1994). We have concentrated on a macroscopic process which can be treated in the ideal MHD regime and paid scant attention to how we should introduce the anomalous resistivity. We should note that we have no definite information on how to make a relevant macroscopic (MHD) model of anomalous resistivity which should be treated in a microscopic regime.      \\subsection{Relation with magnetic Reynolds number} \\label{sec:Rey-indep}  One may feel that our result is strange, because, while the reconnection needs electric resistivity, we treat only ideal MHD and do not consider the dependence on the magnetic Reynolds number. We here discuss a regulation process of the energy conversion speed. As discussed in section \\ref{sec:model}, we assume the magnetic diffusion speed is fully adjustable with the spontaneous inhalation speed $-v_{yp}$ which is determined by the above shock tube problem. This assumption is valid for the case of sufficiently large electric resistivity (small magnetic Reynolds number) in which the spontaneous inhalation speed is smaller than the maximum diffusion speed. Such a large resistivity may be realized by anomalous resistivity. We need an assumption for the anomalous resistivity because we do not have a convincing model for the mechanism and estimated value of anomalous resistivity (Coppi \\& Friedland 1971, Ji et al. 1998, Shinohara et al. 2001). Note that the diffusion speed is adjustable by a change of the diffusion region thickness even if the resistivity is fixed to be constant, but if there is a lower limit for the thickness (e.g., simulation mesh size or the ion Larmor radius), the diffusion speed is bound by an upper limit.  In paper 1, we suppose the Reynolds number to be 24.5. One may think that this is an incredibly large resistivity. We should be careful of what denotes the length in the definition of the magnetic Reynolds number. Usually speaking, the magnetic Reynolds number according to Spitzer resistivity for typical solar flares is extremely large: \\begin{equation} R_m \\sim 10^{13} \\frac{(L/10^7 \\mbox{[m]})(T/10^6 \\mbox{[K]})^{3/2}(B/10^{-3} \\mbox{[T]})}{n/10^{15} [\\mbox{m}^{-3}]} \\end{equation} where $L$, $T$, $B$ and $n$ are the dimension of the entire system, the temperature, the magnetic field strength and the plasma number density,  respectively. We must note that ``the effective Reynolds number'' which is based not on the dimension of the entire system but on that of the diffusion region ($L \\sim 10^0 \\mbox{[m]}$: the ion Larmor radius for solar corona) is important. Hence, the effective Reynolds number $R_e$ for Spitzer resistivity is $R_e \\sim 10^6$. If anomalous resistivity which is much larger than the Spitzer resistivity (typically assumed as $\\sim 10^{5-6} \\times $[Spitzer resistivity] in many numerical MHD simulations) takes place, the effective Reynolds number is estimated as $R_e \\sim 10^{0-1}$. The magnetic Reynolds number adopted in paper 1 (which is the effective Reynolds number: $R_e=24.5$) may be relevant for the case of anomalous resistivity.  A localized resistivity is assumed as a model of anomalous resistivity at the center of the reconnection system. We put the thickness of the initial current sheet $D$ as in paper 1. Here, $D$ is supposed to be of the order of the ion Larmor radius, hence, it is close to the minimum value of the thickness. The dimension of the resistive region is set larger than $D$ (in paper 1, we set $2 D$) in order to deal with change of the diffusion region size. Thus, the effective thickness of the diffusion region can vary between $D$ and $2 D$. The maximum diffusion speed is achieved when the thickness becomes $D$.  The smallest effective magnetic Reynolds number $R_{e min}$ is defined as \\begin{equation} R_{e min} \\equiv \\frac{V_{A0}}{\\eta/D} \\end{equation} where $\\eta$ is the magnetic diffusivity. For the value $R_{emin}=24.5$ in paper 1, the maximum diffusion speed into the diffusion region is determined by the Sweet-Parker model (Sweet 1958, Parker 1963): $v_{dmax} \\sim V_{A0}/\\sqrt{R_{e min}} \\sim 0.2$, while the spontaneous inhalation speed determined by the shock tube approximation is $-v_{yp} \\sim 0.05$. The diffusion speed does not limit the inflow speed because $-v_{yp} < v_{dmax}$. This is why the reconnection rate of our model does not depend on the magnetic Reynolds number.  If the Reynolds number is not so small (small resistivity), say, $R_{e min} \\geq 400$, the diffusion speed is smaller than the spontaneous inhalation speed, and may limit the inflow speed. In such cases, the reconnection rate will be suppressed by the small diffusion speed and depends on the magnetic Reynolds number in ways discussed in the previous literature. This problem is very interesting; however, detailed discussion for such cases is beyond the scope of this work.  In this meaning, the value of $R$ plotted in figure \\ref{fig:rec-rate} will be the maximum value for spontaneous reconnection as a function of $\\beta$ when the resistivity is sufficiently large ($R_{e min} \\leq 400$). The system cannot exceed this value by spontaneous inhalation of the inflow plasma without an external injection.  A similar discussion about spontaneous reconnection was made by Ugai in his series of works (e.g., Ugai \\& Tsuda 1979 and Ugai 1983). He obtained a result supporting Petschek's prediction, i.e., $R \\propto (\\log R_e)^{-1}$. We here compare with the result of Ugai 1983. His interest is focused on the steady state in a finite simulation box (several times the initial current sheet thickness). This is realized in a very late stage from the onset. At that stage, waves emitted from the central region propagate beyond the entire simulation box (Ugai paid attention to a stage later than roughly 7 times the Alfv\\'{e}n transit-time with respect to the entire size of the simulation region). We should note that this is completely different from our case. His resultant reconnection rate logarithmically varies in a range of 0.1-0.3 for $2.5 \\leq R_e \\leq 200$, and the inflow speed varies in a range of 0.1-0.3 $V_{A0}$ while the maximum diffusion speed (defined by the minimum mesh size) varies between $20 \\geq v_{dmax} \\geq 0.3$. The author cannot explain why such a very large inflow speed is attained (e.g., inflow speed $-v_{yp} \\sim 0.05$ in our case), but we can see that their inflow speed is not so small compared with the maximum diffusion speed for $R_e > 10^2$. Thus, the inflow speed might be restricted by the diffusion speed and the reconnection rate might depend on the magnetic Reynolds number in Ugai's result. The author thinks such a situation is unnatural for astrophysical reconnection as discussed our works (e.g., section \\ref{sec:Int} of this paper). Since there is no essential difference between Ugai's numerical procedure and our procedure, if one set a sufficiently wider simulation box than Ugai, he would obtain a similar result to ours.     \\subsection{Boundary condition along the slow shock} \\label{sec:bc}   Finally, we discuss the boundary condition for the inflow region which remained as an open question in paper 2. We obtain the distribution of magnetic vector potential $A_1'$ for the perturbed magnetic field along the edge of the reconnection outflow (see figure \\ref{fig:bc}).  The result is roughly consistent with the result of our numerical simulation, but it cannot explain the numerical result quantitatively. The major difference is in the region $x > x_f$. This is mainly due to the two-dimensional effects of the actual outflow. We here assume a quasi-one-dimensional model for the reconnection outflow for simplicity, but the actual outflow clearly has a two-dimensional structure. For example, the magnetic field lines in this region have a round-shaped upward convexity. If we make a precise model of the reconnection outflow and take account of this two-dimensional shape, such a discrepancy will be reduced. Note that, in our experience, the physical quantities in region 1 or 2 which are important to determine the reconnection rate are insensitive to the model of the slow shock shape, and so we do not need to be too concerned about this problem. Obviously, such further discussion cannot proceed without the help of a computer simulation, and is beyond the scope of this work.    \\subsection{Relation to stationary models} \\label{sec:stat}  As noted in summary section \\ref{sec:sum}, the central region of the entire expanding system tends to be a stationary state which is similar to the Petschek model. The author thinks that the Petschek-like stationary state is unique as the inner solution of spontaneously evolving reconnection systems in free-space with a locally enhanced resistivity (see paper 1). Even when an external circumstance strongly influences the evolution, such the Petschek-like central structure will hold for a long duration of the order of a hundred times Alfv\\'{e}n transit time with respect to the proper scale of the external circumstance (e.g., diameter of a magnetic flux tube in case of solar flares). This is because the induced inflow speed is of the order of $10^{-2} V_{A0}$ (see paper 1), and it needs a long time to change the inflow region according to the influence from the external circumstance (see paper 1).  Other types of stationary states may occur in the following situations: 1) very fast plasma flow ($\\sim V_{A0}$: similar to or faster than the FRW propagation) is injected through the boundary, 2) proper scale of the external circumstance is not much larger than the initial current sheet thickness (e.g., in case of laboratory plasma), 3) resistive region is not so localized, etc. However, our self-similar evolution model may be applicable to many cases in astrophysics. \\\\     The author would like to thank Syuniti Tanuma (Kyoto University), Kazunari Shibata (Kyoto University) and Kiyoshi Maezawa (Japan Aerospace Exploration Agency), who are collaborators on a series of previous papers, for their fruitful comments; Takahiro Kudoh (University of Western Ontario) who gave me a lot of opportunities for useful discussion at the initial stage of this work even in my dream and successively encouraged me. Discussions with anonymous referee were very helpful to improve this paper to be more persuasive. My thanks are due to Mike Kryshak (SOKENDAI) for valuable advice on technical expression in English. I also thank Naoko Kato (SOKENDAI) for her advice on English usage and successive encouragement.     \\appendix"
    },
    "0404/astro-ph0404609_arXiv.txt": {
        "abstract": "We report the discovery of two damped Ly$\\alpha$ absorption-line systems (DLAs) near redshift $z=1$ along a single quasar sightline (Q1727+5302) with neutral hydrogen column densities of $N_{HI} = (1.45\\pm0.15)\\times10^{21}$ and $(2.60\\pm0.20)\\times10^{21}$ atoms cm$^{-2}$.  Their sightline velocity difference of 13,000 km s$^{-1}$ corresponds to a proper separation of 106$h_{70}^{-1}$ Mpc if interpreted as the Hubble flow ($\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$). The random probability of such an occurrence is significantly less than 3\\%. Follow-up spectroscopy reveals neutral gas-phase Zn abundances of [Zn/H] = $-0.58\\pm0.15$ (26.5\\% solar) and $-1.32\\pm0.28$ (4.7\\% solar), respectively. The corresponding Cr abundances are [Cr/H] = $-1.26\\pm0.15$ (5.5\\% solar) and $-1.77\\pm0.28$ (1.7\\% solar), respectively, which is evidence for depletion onto grains.  Follow-up IR images show the two most likely DLA galaxy candidates to have impact parameters of $\\approx22h_{70}^{-1}$ kpc and $\\approx32h_{70}^{-1}$ kpc if near $z=1$. They are significantly underluminous relative to the galaxy population at $z=1$.  To investigate the possibility of additional high-$N_{HI}$ absorbers we have searched the SDSS database for $z>1$ quasars within 30 arcmin of the original sightline.  Five were found, and two show strong \\ion{Mg}{2}-\\ion{Fe}{2} absorption near $z=1$, consistent with classical DLA absorption $\\approx 37$\\% of the time, but almost always $N_{HI} > 10^{19}$ atoms cm$^{-2}$. Consequently, this rare configuration of four high-$N_{HI}$ absorbers with a total sightline velocity extent of 30,600 km s$^{-1}$ may represent a large filament-like structure stretching over a proper distance of 241$h_{70}^{-1}$ Mpc along our sightline, and a region in space capable of harboring excessive amounts of neutral gas. Future studies of this region of the sky are encouraged. ",
        "introduction": "Intervening damped Ly$\\alpha$ absorption-line systems (DLAs) in quasar spectra are very rare, with an incidence of $\\approx 0.17$ per unit redshift at $z=1$ (Rao, Turnshek, \\& Nestor 2004, hereafter RTN2004). Consequently, unless DLAs are correlated, the appearance of two DLAs along any single quasar sightline (``double-damped'') represents a very unlikely event. As such, the discovery of any double-damped absorption  warrants a closer investigation.  Here we report the discovery of double-damped absorption near $z=1$ in the Sloan Digital Sky Survey (SDSS) quasar Q1727+5302 during our most recent {\\it Hubble Space Telescope} UV spectroscopic survey for DLAs (RTN2004).  The purpose of the present paper is to report initial results pertaining to this discovery, and thereby encourage future studies of this region of the sky.  The velocity separation of the absorption is 13,000 km s$^{-1}$, which corresponds to a proper radial distance  of $106h^{-1}_{70}$ Mpc if interpreted as due to the Hubble flow.\\footnote{To calculate physical quantities in this paper we adopt a cosmology with $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$, and $H_0=70$ km s$^{-1}$ Mpc$^{-1}$ ($h_{70} = H_0/70$).} We speculate that this configuration may represent a neutral hydrogen gas filament with a large cosmological extent along our sightline. In fact, the comoving size of this putative filament would be larger than anything previously reported.  As discussed in earlier contributions (e.g., Rao \\& Turnshek 2000, hereafter RT2000, and references therein), DLAs are excellent tracers of the bulk of the neutral hydrogen gas in the Universe, and the aim of our most recent DLA UV survey has been to improve our knowledge of the incidence and cosmological mass density of DLAs at redshifts $z<1.65$.  The new sample which led to the discovery of the double-damped absorption was derived from the SDSS Early Data Release (EDR) (Schneider et al. 2002). We applied a strong \\ion{Mg}{2}-\\ion{Fe}{2} rest equivalent width (REW) selection criterion (RT2000) to optical spectra in order to identify candidate DLA absorption lines ($N_{HI} \\ge 2\\times10^{20}$ atoms cm$^{-2}$), and then we obtained HST STIS UV spectra to confirm or refute their presence.  The current overall success rate for identifying DLAs with this method is $\\approx 37$\\%.  Since the \\ion{Mg}{2}$\\lambda\\lambda$2796,2803 absorption lines are saturated, the REW of the absorption is most closely tied to kinematic spread, not column density.  Recently, Nestor et al. (2003) have discussed evidence for a correlation between kinematic spread and metallicity.  The paper is organized as follows.  In \\S2 we present the HST discovery spectrum for the double-damped absorption, a follow-up MMT spectrum used to determine neutral-gas-phase metal abundances, and IRTF imaging data used to search for galaxies associated with the double-damped absorbers along the quasar sightline. In \\S3 we summarize evidence that exists for strong \\ion{Mg}{2} absorption systems near $z=1$ in other SDSS quasars in the same region of the sky. A brief summary and discussion of the results is presented in \\S4. ",
        "conclusions": "Investigations of the possibility of clustering between DLAs and Lyman break galaxies (LBG) have been made at high redshift (e.g., Gawiser et al. 2001, Aldelberger et al. 2003, Bouch\\'e \\& Lowenthal 2003).  Gawiser et al. (2001) and Aldelberger et al. (2003) find no significant evidence for clustering at $z\\approx4$ and $z\\approx3.2$, respectively, while Bouch\\'e \\& Lowenthal (2003) present weak evidence ($2.6\\sigma$ significance) for clustering at $z\\approx3$ on a size scale up to $\\approx1.8h^{-1}_{70}$ Mpc ($\\approx7h^{-1}_{70}$ Mpc comoving).  Here we have reported the discovery of an apparently non-random (see below) structure on the sky near redshift $z=1$ (\\S2.1) which is much larger than the scales explored in high-redshift DLA-LBG clustering studies and much larger than normal galaxy clustering. The structure consists of two extremely high-$N_{HI}$ (even by DLA standards) DLA absorbers separated by 13,000 km s$^{-1}$ along a single quasar sightline. If interpreted as a cosmological redshift separation, this double-damped structure has a radial proper distance of $106h_{70}^{-1}$ Mpc. The incidence of DLAs above a survey threshold of $N_{HI} \\ge 2\\times10^{20}$ atoms cm$^{-2}$ is $\\approx 0.17$ per unit redshift at $z=1$ (RTN2004).  Given this incidence, the probability of observing a second DLA within $\\Delta z \\approx \\pm0.09$ of another DLA is $<$3\\%. We quote this probability as an upper limit because the incidence of DLAs is significantly smaller for higher column density systems, and the column densities of the two systems included in the double-damped absorption are factors of $\\approx7$ and $\\approx13$ times larger than the survey threshold.$\\footnote{Of the $\\approx50$ DLAs presently known to us at $z<1.65$, the two which make up the double-damped absorption rank as the 8th and 14th highest $N_{HI}$ systems.}$ Therefore, the double-damped absorption may be the result of correlated DLAs and be caused by a cosmologically extended filament of neutral gas along our sightline.   In addition to the identification of this remarkable structure, we have made metal abundance determinations for the two systems which make up the double-damped absorption. We find them to have [Zn/H] = $-0.58\\pm0.15$ (26.5\\% solar) and $-1.29\\pm0.27$ (4.7\\% solar), with evidence for some depletion onto grains in both cases (\\S2.1). Infrared imaging indicates that the two most likely DLA candidate galaxies are relatively faint in relation to the galaxy population at $z=1$, with $K$-band luminosities that are $\\approx0.06L_K^*(z=1)$ and $\\approx0.15L_K^*(z=1)$ (\\S2.2).  These add to the list of underluminous galaxies that have been identified as being responsible for DLA absorption (Rao et al. 2003). The results indicate that the presence of luminous galaxies relative to the local population evidently is not a requirement for the presence of large concentrations of neutral gas. We have also identified two new candidate DLA systems in this same region of the sky, separated from the original sightline by 10 arcmin and 25 arcmin (\\S3). The discovery of these two additional systems increases the probability that this is a non-random structure.  If the two new redshift systems are included, the structure stretches $30,600$ km s$^{-1}$ along the sightline, corresponding to a radial proper distance of $241h_{70}^{-1}$ Mpc.  Filaments nearly as large as that implied by the double-damped absorption have been seen in mock redshift surveys and cold dark matter simulations of structure formation. For example, Faltenbacher et al. (2002) report correlations in cluster orientations with respect to one another and find alignments of galaxy clusters' major axes on comoving scales of $\\approx 140h_{70}^{-1}$ Mpc, corresponding to a proper distance of $\\approx 70h_{70}^{-1}$ Mpc at $z=1$.  The large scale structure simulations of Eisenstein, Loeb, \\& Turner (1997) show similar alignments.  Recently, Palunas et al. (2004) have reported evidence for a structure with a proper size of $\\approx25h^{-1}_{70}$ Mpc ($\\approx80h^{-1}_{70}$ Mpc comoving), which they found during a search for Ly$\\alpha$-emitting galaxies at $z=2.38$.  Miller et al. (2004) have reported evidence for structures on even larger scales based on an analysis of the QSO distribution in the 2dF redshift survey.  For this case of double-damped absorption, the size of the putative neutral hydrogen gas filament would be larger than any claimed so far.  Therefore, our interpretation should be considered speculative pending future studies. Nevertheless, the properties of the double-damped absorption are of course relevant to studies of DLAs in general.  The present-day cosmological model (approximately 73\\% dark energy, 24\\% dark matter and 5\\% ordinary matter, e.g., Spergel et al. 2003) is one in which large-scale structures can form early in time. However, a specific set of cosmological parameters may also indicate that it is highly improbable for certain structures to grow from initial Gaussian perturbations.  Thus, surveys to find evidence for extreme large-scale structure at high redshift have the potential to result in important cosmological constraints. Based on the numbers of DLAs discovered at low redshift so far and the DLA column density distribution, the existence of this double-damped absorption along one sightline represents evidence for a non-random distribution of DLAs which should be further investigated.  \\bigskip  \\centerline{\\bf Acknowledgments}  We thank members of the SDSS collaboration who made the SDSS project a success and who made the EDR spectra available. We acknowledge support from NASA-STScI, NASA-LTSA, and NSF. HST-UV spectroscopy made the $N_{HI}$ determinations possible, while follow-up metal abundance measurements (MMT) and imaging (IRTF) were among the aims of our LTSA and NSF programs. Funding for creation and distribution of the SDSS Archive has been provided by the Alfred P. Sloan Foundation, Participating Institutions, NASA, NSF, DOE, the Japanese Monbukagakusho, and the Max Planck Society. The SDSS Web site is www.sdss.org. The SDSS is managed by the Astrophysical Research Consortium for the Participating Institutions: University of Chicago, Fermilab, Institute for Advanced Study, the Japan Participation Group, 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 University of Washington."
    },
    "0404/astro-ph0404359_arXiv.txt": {
        "abstract": "We have imaged $\\sim$ 1 deg$^{2}$ in the V-band in the direction of the Hercules Cluster (Abell 2151). The data are used to compute for the first time the luminosity function (LF) of galaxies in the cluster down to the dwarf regime (M$_{lim}$ $\\sim$ -13.85 ). The global LF is well described by a Schechter function (\\cite{schechter76}) with best-fit parameters $\\alpha$ = -1.30 $\\pm$ 0.06 and M$_V$$^*$ = -21.25 $\\pm$ 0.25. The radial dependence of the LF has also been studied, finding that it turns out to be almost constant within the errors even further away than the virial radius. Given the presence of significant substructure within the cluster, we have analized the LFs in different regions. While two of the subclusters present LFs consistent with each other and with the global one, the southernmost one exhibits a somewhat steeper faint-end slope. ",
        "introduction": "The luminosity function of galaxies (LF), that is, the probability density of galaxies in a certain population having a given luminosity, is one of the key observable quantities for galaxy evolution theories. The comparison between the LF of galaxies in clusters and in the field should clarify the role played by the environment in regulating galaxy evolution. Recent results by large surveys (\\cite{depropris03}) point towards the composite LF of cluster galaxies having a steeper faint-end slope than field galaxies, revealing that this influence indeed exists. A2151 is a nearby (z = 0.0367), irregular and spiral-rich cluster ($\\sim$ 50\\%). There are strong evidences (from optical and X-ray studies) suggesting that the cluster is still on the process of collapse: the lack of hydrogen deficiency in the spiral population (\\cite{GH85}), the bumpy distribution of the hot intracluster gas and its low X-ray flux (\\cite{HS96}), and the presence of at least three distinct subclusters (\\cite{BDB95}, BDB hereafter) point towards A2151 being a young and relatively unevolved cluster, thus making it an excellent target for the study of the LF and stablishing comparisons with more evolved systems. ",
        "conclusions": ""
    },
    "0404/astro-ph0404390_arXiv.txt": {
        "abstract": "s{The thermal relic abundance of species critically depends on the assumed underlying cosmological model. In the case of neutralinos, freeze-out takes place long before Big Bang Nucleosynthesis, which provides the strongest constraint on the evolution of the Hubble parameter in the Early Universe. We show that non-standard cosmologies, such as models featuring a quintessential scalar field or primordial anisotropies, can lead to large enhancements in the neutralino relic abundance, up to six orders of magnitudes. Within these scenarios, supersymmetric models with large neutralino annihilation cross sections may account for the whole inferred amount of cold dark matter, yielding on the other hand large indirect detection rates.} ",
        "introduction": "Observational cosmology is accumulating evidences for a picture of the Universe dominated by a {\\em dark sector} made up of a non-luminous and non-baryonic component ({\\em dark matter}) and by a gravitationally non-clustering form of energy with negative pressure ({\\em dark energy})~\\cite{Spergel:2003cb}. While the standard cosmological lore, often dubbed {\\em Concordance} or $\\Lambda$CDM {\\em Model} nicely fits most of the observational features~\\cite{giavalisco}, particle physics and theoretical cosmology face the compelling questions related to the nature and the amount of both dark matter and dark energy.  The minimal supersymmetric extension of the Standard Model provides a suitable candidate for cold dark matter, namely the lightest neutralino. The thermal relic abundance of neutralinos $\\Omega_\\chi h^2$ depends both on the particle physics setup, and therefore on the supersymmetric (SUSY) model, and on the underlying cosmological model, which affects the freeze-out of species in the Early Universe. Models where the r\\^ole of dark energy is played by a dynamical scalar field ({\\em Quintessence}), or scenarios where extra energy density components are present at the time of neutralino freeze-out, may yield significant modifications in the relic abundance of neutralinos~\\cite{Profumo:2003hq}. This in turns translates in the intriguing possibility that SUSY models with large neutralino annihilation cross sections, and hence with copious dark matter detection rates, give rise to the required amount of cold dark matter. ",
        "conclusions": ""
    },
    "0404/astro-ph0404445_arXiv.txt": {
        "abstract": "We performed N--body + hydrodynamical simulations of the formation and evolution of galaxy groups and clusters in a $\\Lambda$CDM cosmology. The simulations invoke star formation, chemical evolution with non-instantaneous recycling, metal dependent radiative cooling, strong starbursts and (optionally) AGN driven galactic super winds, effects of a meta-galactic UV field and thermal conduction. The properties of the galaxy populations in two clusters, one Virgo-like ($T$$\\sim$ 3 keV) and one (sub) Coma-like ($T$$\\sim$6 keV), are discussed. The global star formation rates of the cluster galaxies are found to decrease very significantly from redshift $z$=2 to 0, in agreement with observations. The total K-band luminosity of the cluster galaxies correlates tightly with total cluster mass, and for models without additional AGN feedback, the zero point of the relation matches the observed one fairly well. Compared to the observed galaxy luminosity function, the simulations nicely match the number of intermediate--mass galaxies {\\mbox{(--20$\\la M_{\\rm{B}} \\la$--17,}} smaller galaxies being affected by resolution limits) but they show a deficiency of bright galaxies in favour of an overgrown central cD. High resolution tests indicate that this deficiency is {\\it not} simply due to numerical ``over--merging''.  The redshift evolution of the luminosity functions from $z$=1 to 0 is mainly driven by luminosity evolution, but also by merging of bright galaxies with the cD.  The colour--magnitude relation of the cluster galaxies matches the observed ``red sequence'', though with a large scatter, and on average galaxy metallicity increases with luminosity. As the brighter galaxies are essentially coeval, the colour--magnitude relation results from metallicity rather than age effects, as observed. On the whole, a top-heavy IMF appears to be preferably required to reproduce also the observed colours and metallicities of the stellar populations. ",
        "introduction": "Clusters of galaxies are of great interest both as cosmological probes and as ``laboratories'' for studying galaxy formation. The mass function and number density of galaxy clusters as a function of redshift is a powerful diagnostic for the determination of cosmological parameters (see Voit 2004 for a recent comprehensive review). Besides, clusters represent higher than average concentrations of galaxies, with active interaction and exchange of material between them and their environment, as testified by the non-primordial composition of the hot surrounding gas (e.g.\\ Matteucci \\& Vettolani 1988; Arnaud {\\it et~al.}\\ 1992; Renzini {\\it et~al.}\\ 1993; Finoguenov {\\it et~al.}\\ 2000, 2001; De Grandi {\\it et~al.}\\ 2004; Baumgartner {\\it et~al.}\\ 2005).  There are observational and theoretical arguments indicating that clusters are not fair samples of the average global properties of the Universe: the morphological mixture of galaxies in clusters is significantly skewed toward earlier types with respect to the field population, implying star formation histories peaking at higher redshifts than is typical in the field (Dressler 1980; Goto {\\it et~al.}\\ 2003; Kodama \\& Bower 2001; see also Section~3); this is qualitatively in line with the expectation that high density regions such as clusters, in a hierarchical bottom--up cosmological scenario, evolve at an ``accelerated'' pace with respect to the rest of the Universe (Bower 1991; Diaferio {\\it et~al.}\\ 2001; Benson {\\it et~al.}\\ 2001). Although clusters are somewhat biased sites of galaxy formation, they present the advantage of being bound structures with deep potential wells, likely to retain all the matter that falls within their gravitational influence; henceforth, they represent well-defined, self--contained ``pools'', where one can aim at keeping full track of the process of galaxy formation and evolution, and of the global interplay between galaxies and their environment.  The physics of clusters of galaxies has thus received increasing attention in the past decade, benefiting from a number of X--ray missions measuring the emission of the hot intra--cluster gas (e.g., {\\it ASCA, ROSAT, XMM, Chandra}) as well as from extensive optical/NIR surveys probing the distribution of galaxies and their properties (e.g.\\ MORPHS, SDSS, 2MASS). Understanding cluster physics is also crucial to reconstruct, from the observed X--ray luminosity function and temperature distribution, the intrinsic mass function of clusters as a function of redshift, which is a quantity of profound cosmological interest (Voit 2004).  The baryonic mass in clusters is largely in the form of a hot intra--cluster medium (ICM), which dominates by a factor of 5--10 over the stellar mass (Arnaud {\\it et~al.}\\ 1992; Lin, Mohr \\& Stanford 2003). Consequently, early theoretical work and numerical simulations concentrated on pure gas dynamics when modelling clusters. Recently however, attention has focused also on the role of galaxy and star formation, and related effects. Star formation locks--up low entropy gas, and supplies thermal and kinetic energy to the surrounding medium via supernova explosion and shell expansion; both processes likely contribute to the observed ``entropy floor'' in low--mass clusters, and the corresponding breaking of the scaling relations expected from pure gravitational collapse physics (Voit {\\it et~al.}\\ 2003). Besides, star formation is accompanied by the production of new metals and the chemical enrichment of the environment; the considerable (about 1/3 solar) % metallicity of the ICM indicates that a significant fraction of the metals produced --- comparable or even larger than the fraction remaining within the galaxies --- is dispersed into the intergalactic medium (Renzini 2004), affecting the cooling rates of the intra--cluster gas. It is thus clear that the hydrodynamical evolution of the hot ICM is intimately connected to the formation and evolution of cluster galaxies. Only recently, however, due to advances in computing capabilities as well as in detailed physical modelling, cluster simulations have reached a level of sophistication adequate to trace star formation and related effects in individual galaxies, and the chemical enrichment of the ICM by galactic winds, in a reasonably realistic way (Valdarnini 2003; Tornatore {\\it et~al.}\\ 2004).  Indeed so far theoretical predictions of the properties of cluster galaxy populations within a fully cosmological context, have been mainly derived by means of semi--analytical models. High resolution, purely N--body cosmological simulations of the evolution of the collisionless dark matter component, are combined with semi--analytical recipes describing galaxy formation and related physics (such as chemical enrichment, stellar feedback and exchange of gas and metals between galaxies and their environment; see e.g.\\  De Lucia, Kauffmann \\& White 2004 for a recent reference); with such schemes, the evolution of galaxies is ``painted'' on top of that of the simulated dark matter haloes and sub--haloes. The advantage of this technique is that very high resolution can be attained, since pure N--body simulations can handle larger numbers of particles than hydrodynamical simulations; besides, a wide range of parameters can be explored for baryonic physics (e.g.\\ star formation and feedback efficiency, Initial Mass Function, etc.).  In this paper, we present for the first time (to our knowledge) an analysis of the properties of the galaxy population of clusters as predicted directly from cosmological simulations including detailed baryonic physics, gas dynamics and galaxy formation and evolution. The resolution for N--body + hydrodynamical simulations cannot reach the level of the purely N--body simulations that constitute the backbone of semi--analytical models, so we cannot resolve galaxies at the faint end of the luminosity function ($M_B$$\\ga$--16). On the other hand, our simulations have the advantage of describing the actual hydrodynamical response of the ICM to star formation, stellar feedback and chemical enrichment. Although some uncertain physical processes still necessarily rely on parameters (like the star formation efficiency and the feedback strength, see Section~2), once these are chosen, the interplay between cluster galaxies and their environment follows in a realistic fashion, as part of the global cosmological evolution of the cluster.  Moreover, the intimate relation between stellar Initial Mass Function (IMF), stellar luminosity, chemical enrichment, supernova energy input, returned gas fraction and gas flows out of/into the galaxies is included self--consistently in the simulations (while sometimes these are treated as adjustable, {\\it independent} parameters in semi--analytical models). The properties we obtain for the galaxy populations (notably, global star formation rates, luminosity functions and colour--magnitude relations) are then the end result of {\\it ab--initio} simulations, with a far minor degree of parameter calibration than in semi--analytical schemes.  In a standard $\\Lambda$CDM cosmology, we have performed N-body + hydrodynamical (SPH) simulations of the formation and evolution of clusters of different mass, on scales of groups to moderately rich clusters (emission-weighted temperature from 1 to 6~keV). In Paper~I of this series (Romeo {\\it et~al.}\\ 2004, in preparation) we analyze the properties and distribution of the hot ICM in the simulated clusters, and discuss the effects that star formation and related baryonic physics have, on the predicted X--ray properties of the hot gas. Several sets of simulations have been carried out, assuming different IMFs and feedback prescriptions (see Paper~I and Section~2). The chemical and X--ray properties of the ICM are best reproduced assuming a fairly top--heavy IMF and a high, though not extreme, feedback (super-wind) efficiency (the simulations marked, hereafter, as AY-SW; see Paper~I and Table 1).  In this Paper~II we focus instead on the properties of the galaxy population in our simulated richer clusters (with temperatures between 3 and 6~keV), where the number of identified galaxies is statistically significant. We will mainly discuss the results from the AY-SW simulations, favoured by the resulting properties of the ICM; results from simulations with different input physics (IMF, wind efficiency, preheating) are also discussed for comparison, where relevant.  In Section~2 we briefly introduce the code and the simulations (full details are given in Paper~I), as well as the procedure to identify cluster galaxies in the simulations. In Section~3 we discuss the global star formation histories of cluster galaxies, and in Section~4 we determine global luminosities of the clusters simulated with different prescriptions for baryonic physics. In Section~5 and~6 we discuss luminosity functions and colour--magnitude relations of the galaxy population in our clusters, and, finally, in Section~7 we summarize our results. ",
        "conclusions": "In this paper we have presented for the first time and analyzed the properties of the galaxy population in clusters, as predicted from full {\\it ab initio} cosmological + hydrodynamical simulations.  Our results are based on cosmological simulations of galaxy clusters including self-consistently metal-dependent atomic radiative cooling, star formation, supernova and (optionally) AGN driven galactic super-winds, non-instantaneous chemical evolution, effects of a meta-galactic, redshift dependent UV field and thermal conduction. In relation to modelling the properties of cluster galaxies this is an important step forward with respect to previous theoretical works on the subject, e.g. based on semi--analytical recipes super--imposed on N--body only simulations.  The global star formation rates of the ``Virgo'' and ``Coma'' cluster galaxies are found to decrease very significantly with time from redshift $z$=2 to 0, in agreement with what is inferred from observations of the inner parts of rich clusters (e.g., Kodama \\& Bower 2001).  We have determined galaxy luminosity functions for the ``Virgo'' and ``Coma'' clusters in the $B, V, R, H$ and $K$ bands; the comparison to observed galaxy luminosity functions reveals a deficiency of bright galaxies ($M_{\\rm{B}}$$\\la$--20). We carried out a test simulation of ``Virgo'' at eight times higher mass resolution and two times higher force resolution; the results of this test, still running at the present, indicate that the above mentioned deficiency of bright galaxies is {\\it not} due to ``over--merging''; higher resolution simulations of ``Coma'' clusters are in progress as well to further check this point. From a suite of simulations for the ``Virgo'' cluster with different input physics, we find that the deficiency of bright galaxies becomes less prominent with decreasing super-wind strength, in particular for models invoking a Salpeter IMF and only early feedback; in fact more mass can be accumulated in stars and galaxies, with low feedback. Such models, however, present various drawbacks: the cold fraction is too high and the metal production is too low, as seen in the too blue colours of the galaxies and/or in the low metallicity of the ICM, which can hardly be enriched to the observed level of about 1/3 of solar abundance; the latter point is discussed in detail in Paper~I, but it also follows from more general arguments (Portinari {\\it et~al.} 2004). The bright galaxy deficiency might be explained as a selection effect, in the sense that we have selected cluster haloes for the TreeSPH re-simulations, which are ``too relaxed'' compared to an average cluster halo, so that the brightest galaxies have by now merged into the central cD by dynamical friction; we shall return to this in a forthcoming paper.  The redshift evolution of the luminosity functions from redshift $z$=1 to 0 is mainly driven by passive luminosity evolution of the stellar populations, but also by the above mentioned merging of bright galaxies into the cD.  The slope of the colour--magnitude relation of the simulated galaxies is in good agreement with the observed one, however the scatter is larger than observed, partly due to poissoinian noise within the fainter galaxies which are formed by small numbers of star particles. Such internal dispersion in stellar ages is also responsible for a luminosity dimming between $z=1$ and $z=0$, faster than indicated by the observed evolution of the Fundamental Plane. The typical average galaxy colours are best matched when adopting a top-heavy IMF (as originally suggested by Arimoto \\& Yoshii 1987), while Salpeter IMF simulations yield too blue colours. Moreover we find that the average metallicity of the simulated galaxies increases with luminosity, and that the brightest galaxies are essentially coeval. Hence, the colour--magnitude relation results from metallicity rather than age effects, as concluded by Kodama \\& Arimoto (1997) on the basis of its observed evolution."
    },
    "0404/astro-ph0404029_arXiv.txt": {
        "abstract": "We present the analysis of the coadded rest-frame UV spectrum ($1200<z<2000$ \\AA) of five $K$-luminous galaxies at $z\\sim 2$ from the K20 survey. The composite spectrum is characterized by strong absorption lines over the UV continuum from C, N, O, Al, Si, and Fe in various ionization stages. While some of these lines are interstellar, several among the strongest absorptions are identified with stellar photospheric lines. Most of the photospheric and interstellar features are stronger in the $K$-luminous composite spectrum than in LBGs at $z\\sim 3$. This suggests higher metallicity and possibly also larger interstellar velocity dispersion caused by macroscopic motions. The absorption lines and the slope of the UV continuum is well matched by the spectrum of the nearby luminous infrared galaxy NGC~6090, which is in the process of merging. A metallicity higher than solar is suggested by comparing the pure photospheric lines (SiIII, CIII, FeV) with starburst models. The evidence of high metallicity, together with the high masses, high star-formation rates, and possibly strong clustering, well qualify these galaxies as progenitors of local massive ellipticals. ",
        "introduction": "Despite that as much as 75\\% of the stellar mass in the universe is in early-type galaxies (e.g. Fukugita et al. 1998), the epoch at which these massive galaxies have assembled and acquired their high metallicity is still unknown. Although the progenitors of these massive galaxies should be seen at high redshifts, connecting them with their high redshift counterparts remains a difficult task. The currently best known high-$z$ galaxy population is that of Lyman Break Galaxies (LBGs, Steidel et al. 1999). Like in local starbursts, the strongest features in the rest-frame UV spectrum of LBGs are interstellar and photospheric absorption lines of C, N, O, Si, and Fe, typical of hot massive stars (e.g. Shapley et al. 2003, S03 hereafter). They have sub-solar metallicity (Pettini et al. 2001), very young stellar population ($<$0.2 Gyr), low masses ($\\sim10^{10}M_\\odot$). Moreover, it appears that the spatial clustering of $z=3$ LBGs is too low for them to be direct progenitors of $z\\sim1$ Extremely Red Objects (EROs) and local massive galaxies (Daddi et al. 2001).  While these UV-selected LBGs are certainly an important component of the high-$z$ galaxy population, it is still not clear how representative of the entire population at high-$z$ they are. There might be other objects at high-$z$ that would not be picked up by this selection. In fact, there are evidences that $K$-selected galaxies at $2<z<4$  are much more strongly clustered than UV-selected ones (Daddi et al. 2003). For instance, the $K$-selected galaxies which are in the so-called {\\em Redshift Desert} at $1.5<z_{\\rm spec}<2.5$ (Daddi et al. 2004, D04 hereafter) appear to be an important class of objects, with overall physical properties quite distinct from those of LBGs. Comparing the properties of LBGs at $z=3$ from S03 with the sample investigated in this paper taken from D04, it appears that $K$-band luminous galaxies at $z\\sim2$ have higher reddening ($E(B-V)\\sim0.3$--0.4 vs $\\sim0.10$--0.17), redder UV slopes ($-\\beta\\sim0.54$  vs $\\sim0.73$--1.09), higher SFRs ($\\sim$100--500 M$_\\odot$/yr vs $\\sim25$--50 M$_\\odot$/yr), older ages ($\\sim700$ Myr vs $\\sim100$--300 Myr) and, therefore, higher masses ($\\sim10^{11}M_\\odot$ vs $\\sim10^{10}$$M_\\odot$). These differences are likely arising as a consequence of the $K$-selection, as opposed to the UV selection for LBGs. Given their observational properties, it has been argued that these $K$-luminous galaxies may represent better candidates than LBGs as star forming progenitors of massive early-type galaxies (D04). If these objects are indeed the progenitors of local early-type galaxies one would expect them to harbor more metal rich stellar populations and interstellar medium, compared to LBGs. In this Letter we revisit these $K$-luminous galaxies by analyzing their composite spectrum and comparing it with those of various templates, including LBGs and starburst galaxies, in order to gather new insight on their metal enrichment, on their nature and on their role in galaxy evolution. ",
        "conclusions": "The most striking feature of the near-UV spectrum of $K$-band luminous galaxies at $z\\sim 2$ is the presence of strong stellar absorption features, with strengths not previously reported at high redshifts that suggest a high metallicity for this class of high-$z$ galaxies. Stellar photospheric lines are important diagnostic of the stellar content of galaxies and their detection also demonstrate an UV continuum without a dominant AGN contribution. These stellar features are typical of B-type stars (de Mello et al. 2000), which are dominating the UV flux in the explored spectral region, rather than short-lived very massive O stars. This can be well expected given that the strong Balmer breaks detected in the near-IR imply high mass contents and continuous SFR for $\\sim 700$ Myr (see D04). The fact that our composite spectrum was generated with a small number of galaxies may affect these results, as a few peculiar objects could be present. Coaddition of a larger number of spectra is desirable for the future in order to smooth out the contribution of individual objects.  As these $K$-selected starbursts are found to have redder UV slopes and higher reddening than UV-selected galaxies at $1.5<z<3.5$ (the UV selection requires blue and flat slopes), it seems reasonable to find that they have also a higher metal content. It is not clear whether these galaxies are evolutionary descendants of $z\\simgt3$ LBGs. The higher metallicity and dust content could be a consequence of the continuous star formation which occurred between the two epochs. However, if the D04 tentative estimate of a correlation length of $r_0>7$ \\h1 Mpc for the $K$-selected galaxies at $z\\sim 2$ is confirmed, they seem to be more clustered than the LBGs at z$\\sim 3$. This suggests a different evolutionary path for the two populations. Sawicki \\& Yee (1998) suggested that LBGs are either the progenitors of present-day sub-L* galaxies, or may form luminous galaxies through mergers as they evolve. However, one would expect a stronger clustering in the latter case (Daddi et al. 2001).  The signatures of their high metallicity, the confirmation of very high SFRs, together with the evidences for high masses and possibly strong clustering reinforce the suggestion that these $K$-luminous objects at $z\\sim 2$ well qualify as progenitors of local early-type galaxies, caught while still in the act of actively forming stars. Indeed, with their masses being estimated above 10$^{11}$M$_\\odot$ they would most likely become early-type galaxies at $z=0$."
    },
    "0404/astro-ph0404503_arXiv.txt": {
        "abstract": "We summarize the detection rates at wavelengths other than optical for $\\sim$99,000 galaxies from the Sloan Digital Sky Survey (SDSS) Data Release 1 ``main'' spectroscopic sample. The analysis is based on positional cross-correlation with source catalogs from ROSAT, 2MASS, IRAS, GB6, FIRST, NVSS and WENSS surveys. We find that the rest-frame UV-IR broad-band galaxy SEDs form a remarkably uniform, nearly one parameter, family. As an example, the SDSS $u$ and $r$ band data, supplemented with redshift, can be used to predict $K$ band magnitudes measured by 2MASS with an rms scatter of only 0.2 mag; when measurement uncertainties are taken into account, the astrophysical scatter appears not larger than $\\sim$0.1 mag. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404517_arXiv.txt": {
        "abstract": "We present  CS(3-2) and  DCO$^+$(2-1) observations of 94 starless cores and compare the results with previous CS(2-1) and $\\rm N_2H^+$(1-0) observations to study inward motions in starless cores. Eighty-four cores were detected in both CS and DCO$^+$ lines. A significant number of CS(3-2) profiles  and a small number of $\\rm DCO^+$(2-1) lines show the classical  ``infall asymmetry'' similar to that seen in CS(2-1) observations. The $\\rm DCO^+$(2-1) lines, however, usually show a single Gaussian peak. The integrated intensity of $\\rm N_2H^+$ correlates well with that of $\\rm DCO^+$(2-1), but poorly with that of CS(2-1) and CS(3-2), suggesting that CS suffers significantly more depletion onto grains than do either  $\\rm DCO^+$ or $\\rm N_2H^+$. Despite these depletion effects, there is evidently enough optical depth for the CS(3-2) and CS(2-1) spectral lines to exhibit infall asymmetries.  The velocity shifts of the CS(3-2) and (2-1) lines with respect to $\\rm N_2H^+$ correlate well with each other and have similar distributions. This implies that, in many cores, systematic inward motions of gaseous material may occur over a range of density of at least a factor $\\sim 4$.  We identify 18 infall candidates based on observations of CS(3-2), CS(2-1), $\\rm DCO^+$ (2-1) and $\\rm N_2H^+$ (1-0).  The eight best candidates, L1355, L1498, L1521F, L1544, L158, L492, L694-2, and L1155C-1, each show at least four indications of infall asymmetry and no counter-indications.  Fits of the spectra to a 2-layer radiative transfer model in ten infall candidates suggest that the median effective line-of-sight speed of the inward-moving gas is $\\sim 0.07$ $\\rm km~s^{-1}$ for CS (3-2) and $\\rm \\sim 0.04~km~s^{-1}$ for CS(2-1). Considering that the optical depth obtained from the fits is usually smaller in CS(3-2) than in (2-1) line, this may imply that CS(3-2) usually traces inner denser gas in higher inward motions than CS(2-1). However, it is also possible that this conclusion is not representative of all starless core infall candidates, due to the statistically small number analyzed here. Further line observations will be useful to test this conclusion. ",
        "introduction": "Starless cores are dense ($\\rm \\ga 10^4~cm^{-3}$) condensations in molecular clouds without any embedded Young Stellar Objects (YSOs) (e.g., Lee \\& Myers 1999-LM99). They are the best known candidates to possibly form stars (Tafalla et al. 1998; Lee, Myers, \\& Tafalla 1999-LMT99; Jijina, Myers, \\& Adams 1999; Lee, Myers, \\& Tafalla 2001-LMT01) and, therefore, provide an opportunity to study  the initial conditions of star formation (Benson \\& Myers 1989; Ward-Thompson et al. 1994; Shirley et al. 2000; Andr\\'e, Ward-Thompson, \\& Barsony 2000).  The inward motions in starless cores are one of the essential elements needed to understand the onset of star formation. These motions can be  studied by observing two kinds of molecular lines, optically thin and thick tracers, to detect the spectral ``{\\it infall asymmetry}'' : a double peaked profile in which the blue component is brighter than the red one, or a skewed single blue peak in an optically thick tracer (e.g., CS2-1 and $\\rm H_2CO$2-1), and a symmetric single peak in an optically thin tracer (e.g., $\\rm C^{18}O$1-0 and $\\rm N_2H^+$1-0) (Hummer \\& Rybicki 1968; Leung \\& Brown 1977).  This method has been used  quite successfully to observe some general features of inward motions in starless cores (e.g., LMT99, LMT01). The inward motions in starless cores  are found to be slow (typically 0.05  - 0.1 $\\rm km~s^{^-1}$), and spatially extended (typically 0.1 pc) (LMT01) and can be partly explained by ambipolar diffusion in a super-critical core, or by turbulent dissipation (e.g., Nakano 1998; Myers \\& Lazarian 1998; Ciolek \\& Basu 2000). However, the physics of inward motions in starless cores is still not fully understood.  Different molecular species and various transitions of specific molecules are sensitive to different chemical and physical conditions  and, therefore, can probe different physical regions of the infalling cores. Therefore, observations of many different species and transitions would help to widen our present knowledge of inward motions in starless cores (e.g., Onishi et al. 1999; Caselli et al. 1999;  Gregersen \\& Evans 2000; Lee et al. 2003).  This study is a continuation of our previous systematic infall studies, which used CS(2-1) and $\\rm N_2H^+$(1-0) by LMT99 and LMT01. In this study we use two $\\it new$ tracers: CS(3-2) and $\\rm DCO^+$(2-1). The CS(3-2) line has higher critical density ($\\rm n_{cr}\\approx 1.3\\times10^{6}~cm^{-3}$) than CS(2-1) by  a factor of about 4 (Evans 1999) and so is likely a much better tracer of dense gas. Thus, the CS(3-2) line will be a better probe of the kinematics closer to the nucleus of the core unless the optical depth of the CS(3-2) is larger than CS(2-1).  DCO$^+$(2-1)  is  similar to $\\rm N_2H^+$(1-0); it is also a  high density tracer and is usually optically thin. However, the DCO$^+$(2-1) has high enough opacity and is bright enough to make extensive surveys feasible. In this paper we present new survey results with CS(3-2) and DCO$^+$(2-1) toward starless cores and, by combining these with previous CS(2-1) and $\\rm N_2H^+$(1-0) data, discuss how inward motions in starless cores occur.  In $\\S$2 we explain details of the observations such as the observational equipment and methods, target selection, frequency determination of observing lines, and data reduction. Detection statistics, spectral features of CS and DCO$^+$, $\\rm \\delta V$ analysis of the observing lines, and selection of infall candidates  are described in $\\S$3. Implications of the observational results are discussed in $\\S$4. In $\\S$5 we summarize the main conclusions of this study. ",
        "conclusions": "\\subsection{Infall speeds} We derive infall speeds for a sub-sample of cores via fits to their CS spectra using a two-layer model consisting of  a cool (2.7 K) absorbing front screen moving away from us and a warm emitting rear layer approaching us with the same speed (see LMT01). Our non-linear least squares fitting routine has 5 free parameters, peak optical depth, infall speed, velocity dispersion, excitation temperature of the rear layer and LSR velocity. In the model, both layers are assumed to have the same velocity dispersion and move with the same speed but in opposite direction. Note that the number of fit parameters is more than the 3 required to fit any line with a Gaussian.  However, the two additional parameters are needed to specify the peak optical depth and the infall speed. Thus, this model has the least number of parameters possible to describe a line with infall asymmetry due to self-absorption and internal motion. Indeed, each parameter reflects different aspects of the line profile and so can be well constrained by the fitting procedure. The line width, LSR velocity, and the excitation temperature are very well determined by the width, centroid velocity, and intensity scale of the observed spectrum. The  peak optical depth is sensitive to the depth of the self-absorption dip and so fits to spectra with a clear dip or red shoulder can easily constrain the optical depth. In addition, since the first four parameters can be constrained by fits to the spectra, the last parameter - infall speed - can be reliably determined with good precision. This is the main reason that, to derive infall speed,  we fit only spectra with clear double peaks, or a blue peak and a clear red shoulder.  Fits were conducted for the CS(3-2) profiles of 12 sources showing such infall asymmetry. Fig. 6 shows that our model easily reproduces the observed spectral features. Note that in the figure we did not fit a model profile  to the CS(2-1) spectra of two sources (L1355 and L1155C-1) although they do show infall asymmetry. This is because they show neither clear double peaks nor a red shoulder in their spectra and so the fits to their spectra would result in large errors in the derived infall speeds. Typical 1-D infall speeds for the sources modeled are $\\sim$ $\\rm 0.02 - 0.13~km~s^{-1}$.  We also obtained the infall speeds for a sub-sample of 8 sources using fits to the CS(2-1) spectra which show similar asymmetry to that of CS(3-2). The resulting infall speeds derived from our model fits to both the CS(3-2) and CS(2-1) observations are compared in Fig. 7. An interesting feature in Fig. 7 is that the CS(3-2) lines generally have higher infall speeds than CS(2-1); typically $\\sim 0.07$ $\\rm km~s^{-1}$ for CS (3-2) and $\\rm \\sim 0.04~km~s^{-1}$ for CS (2-1).  We believe that this is not dominated by fitting errors to the infall speeds since the difference between the CS(2-1) and CS(3-2) speeds is much larger than the fitting error. As discussed in LMT01, we tried to estimate the fitting error of the infall speed by taking two best-fit model spectra of CS(2-1) and CS(3-2) for L1521F, adding random noise with the same rms as the observed noise, and re-fitting a model to the spectra. Fig. 8-(a) plots the distribution of infall speeds obtained from the best-fit for twenty synthetic CS(2-1) and (3-2) spectra in L1521F. This figure shows that the two distributions of infall speeds are clearly separated. The estimated 1 sigma uncertainties of the distribution of infall speeds are about $\\sim 0.001 \\rm ~km~s^{-1}$ for the CS(2-1) and $\\sim 0.003 \\rm ~km~s^{-1}$ for the CS(3-2), which are negligibly small compared to the difference ($\\sim 0.03 \\rm ~km~s^{-1}$) between the infall speeds for two lines. Note that most of spectra used in the fitting procedure are as good as the L1521F spectra. Thus, their fit results are thought to be also as good as for L1521F, suggesting that the tendency for CS(3-2) line to trace higher inward speeds than CS(2-1) is real.  One might question how the infall speed can be determined with precision as high as indicated here. We note that the infall speed, like the centroid velocity, is sensitive to emission from the entire spectrum, so with sufficient signal to noise ratio one can determine the infall speed to a fraction of a channel width. Note that this precision, due to random uncertainties, is not the same as the accuracy, due to systematic uncertainties. Detailed comparisons show that the accuracy of these estimates of two-layer infall speed, with respect to the values from a more sophisticated Monte Carlo model, is likely to be a few 0.01 $\\rm km~s^{-1}$, for spectra with resolution and signal to noise ratio similar to those considered here (De Vries and Myers 2004, in preparation). Furthermore the infall speed from the two-layer model usually increases monotonically with the infall speed from the detailed models. Thus even if the two-layer infall speed is systematically incorrect by a few 0.01 $\\rm km~s^{-1}$, the sign of its gradient with optical depth, which we derive in this paper, is probably unaffected.  Figure 8-(b) and (c) are another test to demonstrate how different infall speeds are required to  produce the  difference between the CS(2-1) and CS(3-2) line profiles.  The solid line in Figure 8-(b) shows the best  $\\chi^2$ fit to the CS(2-1) spectrum in L1521F, which produces an infall speed of $0.014 \\rm ~km~s^{-1}$ and a good match to the observed spectral line profile. Similarly, the solid line in  Figure 8-(c) shows the best  $\\chi^2$ fit to the CS(3-2) spectrum in L1521F, producing an infall speed of $0.045 \\rm ~km~s^{-1}$ and a  good fit to the observed spectral line profile.   However, the dashed line in Figure 8-(b) shows the spectrum obtained if we attempt to fit the CS(2-1) spectrum with the CS(3-2) velocity (and vice-versa in Figure 8-(c)).  Casual inspection of these fits, as well as the $\\chi^2$ values, reveals that small changes in the infall velocity result in poor fits to the spectra. This strongly suggests that the difference between the CS(2-1) and CS(3-2) infall velocities is statistically significant, and that the CS(3-2) line traces higher infall speeds than does the CS(2-1) line.  This analysis with the 2-layer model indicates that the CS 3-2 profiles require slightly but distinctly higher speeds than the 2-1 profiles. However it still remains to be determined whether these higher speeds correspond to denser gas in the inner regions of the cores. Higher speeds associated with denser gas if the line formation were dominated by excitation effects, since the centrally  condensed structure of the core implies that the excitation temperature decreases outward, and since the 3-2 line has a higher critical density than the 2-1 line by a factor $\\sim 4$ (Evans 1999).  But the higher speeds could also correspond to less dense gas, if the line formation were influenced more by optical depth effects, and if the optical depth of the 3-2 line exceeds significantly that of the 2-1 line (e.g. Lucas 1976). The other possible problem in our analysis is resolution difference between CS(2-1) ($\\sim 25''$) and (3-2) data ($\\sim 43''$) which may affect the profile shapes that we observe. To deal with these problems, we use the FCRAO data for CS(2-1) (LMT01) which have  similar beam resolution ($\\sim 52''$) to that of CS (3-2) data. Using the same two layers model, we derive the infall speed and optical depths for FCRAO CS(2-1) lines. Table 3 lists the parameters obtained from the fits for all the data, including 12m data (CS 3-2), Haystack and FCRAO data (CS 2-1). The table clearly indicates that, regardless of the different beam resolution of the data, the CS(3-2) line tends to trace higher infall speed than the CS(2-1) line, and that the CS(3-2) lines have lower optical depth than CS(2-1) lines for all sources except for L158 and L1445.  We also believe that, as the case for the infall speeds, the optical depth derived from the fits is not dominated by fitting errors. In performing the  error analysis for the infall speed in L1521F (see above), the optical depths for each of the twenty CS(2-1) and (3-2) spectra were obtained at the same time. Fig. 9-(a) shows the distribution of such optical depths, indicating  that the two distributions of the optical depths are clearly separated. The estimated 1 sigma uncertainties of the distribution are about $\\sim 0.035$ for the CS(2-1) and $\\sim 0.03$ for the CS(3-2), which are very small compared to the difference ($\\sim 0.24$) between the optical depths. Fig. 9-(b) and (c) also confirm that the difference is  statistically significant as it was for the infall speeds (Fig. 8-b and c).  Fig. 10 visualizes the distribution of the sources in infall speed difference versus optical depth difference between two CS(3-2) and (2-1) lines. The data consist of seven points for FCRAO Data and eight points for Haystack data. Data points from the same core are joined by a line. As shown, nearly all of the data points (11 of 15) and majority of the cores (6 1/2 out of 10) lie in the upper right quadrant, the regime where infall speed increases inward, indicating that the 3-2 line  usually traces gas which is closer to the core center and which has faster infall speed than does the 2-1 line.  This means  that the infall speed generally increases inward for this sample and these lines.  Note that this trend is evident in both the FCRAO and Haystack data, and so apparently the change in resolution between these telescopes does not reverse the tendency for points to occupy the upper right quadrant in Fig. 10.  In summary the conclusion that CS(3-2) traces faster and inner gas than CS(2-1) seems now reliable at least in infall candidates considered among our sample. However, we are still cautious of whether this is necessarily true of most starless cores, or even of most infall candidates, because of the small number of objects involved in our analysis. Further detailed observations of more infall candidates with the same lines are needed to test whether the conclusion generally apply to starless cores or most infall candidates. Furthermore, our conclusion may depend on many factors such as observing lines and should be further tested.  For example infall speeds from these CS observations will be able to be compared with those from high-S/N observations of ``low-depletion'' species like  $\\rm N_2H^+$ and DCO$^+$ (see the section 4.3 ) if they show enough infall asymmetry. Further such line observations are necessarily required.  \\subsection{Implication of the similar $\\rm \\delta V$ distribution for CS(3-2) and CS(2-1)} Fig. 3 and 4 provide important clues as to how inward motions occur. Given the expectation that CS(3-2) is a better probe of the denser inner region, one might expect a more skewed $\\rm \\delta V$ distribution for CS(3-2) than for CS(2-1). However, the $\\rm \\delta V$ distribution for CS(3-2) is found to be statistically very similar to that for CS(2-1) (Fig. 3).  What does this imply about inward motions in cores? We saw in the previous section that CS(3-2) traces faster infalling gas than CS(2-1) by typically $\\rm \\sim 0.03~km~s^{-1}$. The increase in infall velocity would result in a similar increase in line velocity which was used to compute $\\rm \\delta V$. Note, however, that this increase is not that great compared to the $\\rm N_2H^+$ line width, typically 0.3 $\\rm km~s^{-1}$.  Therefore the typical increase in $\\rm \\delta V$ is about 0.03/0.3, or 0.1, or only about half of a bin size in the histograms in Fig. 3 which explains why the histograms look similar despite the higher critical density and higher infall speed for  CS(3-2) compared to (2-1). The similar $\\rm \\delta V$ distribution between CS(3-2) and CS(2-1) may suggest that, in many cores,  inward-moving gas extends over a substantial range of core densities. Furthermore, the good correlation between $\\rm \\delta V_{CS21}$ and $\\rm \\delta V_{CS32}$ (Fig. 4) suggests that the lower density and higher density gas probed by the 2-1 and 3-2 lines respectively move together, in a systematic rather than random fashion.  \\subsection{Depletion of CS and $\\rm DCO^+$} The CS molecule is known to easily deplete out of the gas phase by adsorbing onto dust in  cold dense cores (e.g., Bergin \\& Langer 1997). Therefore, it can be argued that CS cannot be used to probe the kinematics in the nuclei of infalling cores. $\\rm N_2H^+$, however, does not seem to suffer from the effects of depletion as readily as CS (Tafalla et al. 2002; Bergin \\& Langer 1997). Recent mapping surveys of several starless cores have shown that the $\\rm DCO^+$ emission usually has a similar intensity distribution to that of $\\rm N_2H^+$ (Lee et al. 2004 and Bourke et al. 2004 in prep.), suggesting that $\\rm DCO^+$ also does not readily deplete onto grains. Fig. 11 supports this idea. The figure compares the integrated intensities of $\\rm N_2H^+$, CS(2-1), and CS(3-2) of the sources. Note that, in constructing  Fig. 11, we dropped all sources with clear self-absorption features in CS(2-1) (17: number of the dropped sources), CS(3-2) (21) and DCO+(2-1) (14) to minimize any optical depth effects. The middle and bottom panels of Fig. 11 compare the $\\rm N_2H^+$ integrated intensities with that of CS(3-2) and CS(2-1) and show weak to negligible correlations [$r=$0.54  between $\\rm N_2H^+$ and CS(3-2), and $r=$0.40 for $\\rm N_2H^+$ and CS(2-1) ], suggesting the possibility that CS is depleting out onto grains. However, we should note that despite the strong likelihood of CS depletion in the cores we observe, there is evidently still enough optical depth, excitation gradient, and relative motion to give the  CS (2-1) and (3-2) spectral lines an  infall asymmetry. It will need a more detailed chemical evolution models that self-consistently solve for radiative transfer effects to specify the extent and degree of this depletion.  On the other hand, the upper panel of Fig. 11 compares the $\\rm N_2H^+$  and $\\rm DCO^+$ integrated intensities and shows a stronger correlation ($r=$0.75) which suggests that the $\\rm DCO^+$ does not suffer from significant depletion (or, at least, only depletes as much as $\\rm N_2H^+$)."
    },
    "0404/astro-ph0404498_arXiv.txt": {
        "abstract": "We present the preliminary results of 235 MHz, 327 MHz and 610 MHz observations of the galaxy cluster A3562 in the core of the Shapley Concentration. The purpose of these observations, carried out with the Giant Metrewave Radio Telescope (GMRT, Pune, India), was to study the radio halo located at the centre of A3562 and determine the shape of its radio spectrum at low frequencies, in order to understand the origin of this source. In the framework of the re--acceleration model, the preliminary analysis of the halo spectrum suggests that we are observing a young source (few $10^8$ yrs) at the beginning of the re--acceleration phase. ",
        "introduction": "A number of X--ray luminous galaxy clusters show large--scale synchrotron radio emission associated to the intracluster medium. These diffuse sources are know as radio halos, when they are located at the centre of the hosting cluster and show low or negligible polarization, and radio relics, when found in the cluster outskirts and highly polarized.  \\noindent Both halos and relics have low surface brightness, large linear size (from $\\sim$0.5 Mpc to more than 1 Mpc) and steep integrated spectra, i.e. $\\alpha >1$ (S $\\propto$ $\\nu^{-\\alpha}$).  \\noindent Radio halos and relics represent the most striking evidence for $\\mu$G magnetic fields on cluster scale and relativistic electrons diffused within the whole cluster volume (for a recent review see Giovannini \\& Feretti 2002).  \\noindent The existence of this class of radio sources is believed to be connected to cluster mergers, since thus far they have only been found in clusters with significant signs of a current or recent merging event. In particular the leading hypothesis for the origin of the observed radio emission from these objects is a re--acceleration process, probably via turbolence powered by cluster mergers, of a population of relatively low energy ($\\gamma \\sim 10^3$) electrons, initially injected in the intracluster medium (two-phase model, Brunetti et al. 2001). ",
        "conclusions": ""
    },
    "0404/astro-ph0404451_arXiv.txt": {
        "abstract": "Turbulence affects the structure and motions of nearly all temperature and density regimes in the interstellar gas. This two-part review summarizes the observations, theory, and simulations of interstellar turbulence and their implications for many fields of astrophysics. The first part begins with diagnostics for turbulence that have been applied to the cool interstellar medium, and highlights their main results. The energy sources for interstellar turbulence are then summarized along with numerical estimates for their power input. Supernovae and superbubbles dominate the total power, but many other sources spanning a large range of scales, from swing amplified gravitational instabilities to cosmic ray streaming, all contribute in some way. Turbulence theory is considered in detail, including the basic fluid equations, solenoidal and compressible modes, global inviscid quadratic invariants, scaling arguments for the power spectrum, phenomenological models for the scaling of higher order structure functions, the direction and locality of energy transfer and cascade, velocity probability distributions, and turbulent pressure.  We emphasize expected differences between incompressible and compressible turbulence. Theories of magnetic turbulence on scales smaller than the collision mean free path are included, as are theories of magnetohydrodynamic turbulence and their various proposals for power spectra. Numerical simulations of interstellar turbulence are reviewed. Models have reproduced the basic features of the observed scaling relations, predicted fast decay rates for supersonic MHD turbulence, and derived probability distribution functions for density. Thermal instabilities and thermal phases have a new interpretation in a supersonically turbulent medium. Large-scale models with various combinations of self-gravity, magnetic fields, supernovae, and star formation are beginning to resemble the observed interstellar medium in morphology and statistical properties. The role of self-gravity in turbulent gas evolution is clarified, leading to new paradigms for the formation of star clusters, the stellar mass function, the origin of stellar rotation and binary stars, and the effects of magnetic fields. The review ends with a reflection on the progress that has been made in our understanding of the interstellar medium, and offers a list of outstanding problems. ",
        "introduction": "\\label{sect:intro}  In 1951, von Weizs\\\"acker (1951) outlined a theory for interstellar matter (ISM) that is similar to what we believe today: cloudy objects with a hierarchy of structures form in interacting shock waves by supersonic turbulence that is stirred on the largest scale by differential galactic rotation and dissipated on small scales by atomic viscosity. The ``clouds'' disperse quickly because of turbulent motions, and on the largest scales they produce the flocculent spiral structures observed in galaxies. In the same year, von Hoerner (1951) noticed that rms differences in emission-line velocities of the Orion nebula increased with projected separation as a power law with a power $\\alpha$ between 0.25 and 0.5, leading him to suggest that the gas was turbulent with a Kolmogorov energy cascade (for which $\\alpha$ would be $0.33$; Section 4.6). Wilson et al. (1959) later got a steeper function, $\\alpha\\sim0.66$, using better data, and proposed it resulted from compressible turbulence. Correlated motions with a Kolmogorov structure function (Section \\ref{sect:diag}) in optical absorption lines were observed by Kaplan (1958). One of the first statistical models of a continuous and correlated gas distribution was by Chandrasekhar \\& M\\\"unch (1952), who applied it to extinction fluctuations in the Milky Way surface brightness. Minkowski (1955) called the ISM ``an entirely chaotic mass ... of all possible shapes and sizes ... broken up into numerous irregular details.''  These early proposals regarding pervasive turbulence failed to catch on. Interstellar absorption and emission lines looked too smooth to come from an irregular network of structures -- a problem that is still with us today (Section 2). The extinction globules studied by Bok \\& Reilly (1947) looked too uniform and round, suggesting force equilibrium. Oort \\& Spitzer (1955) did not believe von Weizs\\\"acker's model because they thought galactic rotational energy could not cascade down to the scale of cloud linewidths without severe dissipation in individual cloud collisions. Similar concerns about dissipation continue to be discussed (Sections \\ref{sect:heating}, \\ref{sect:decay}). Oort and Spitzer also noted that the ISM morphology appeared wrong for turbulence: ``instead of more or less continuous vortices, we find concentrated clouds that are often separated by much larger spaces of negligible density.'' They expected turbulence to resemble the model of the time, with space-filling vortices in an incompressible fluid, rather than today's model with most of the mass compressed to a small fraction of the volume in shocks fronts. When a reddening survey by Scheffler (1967) used structure functions to infer power-law correlated structures up to 5$^\\circ$ in the sky, the data were characterized by saying only that there were two basic cloud types, large (70 pc) and small (3 pc), the same categories popularized by Spitzer (1968) in his textbook.   Most of the interesting physical processes that could be studied theoretically at the time, such as the expansion of ionized nebulae and supernovae (SNe) and the collapse of gas into stars, could be modeled well enough with a uniform isothermal medium. Away from these sources, the ISM was viewed as mostly static, with discrete clouds moving ballistically. The discovery of broad emission lines and narrow absorption lines in {\\it H I} at 21 cm reinforced this picture by suggesting a warm intercloud medium in thermal pressure balance with the cool clouds (Clark 1965). ISM models with approximate force equilibrium allowed an ease of calculation and conceptualization that was not present with turbulence. Supernovae were supposed to account for the energy, but mostly by heating and ionizing the diffuse phases (McCray \\& Snow 1979). Even after the discoveries of the hot intercloud (Bunner et al. 1971, Jenkins \\& Meloy 1974) and cold molecular media (Wilson et al. 1970), the observation of a continuous distribution of neutral hydrogen temperature (Dickey, Salpeter, \\& Terzian 1977), and the attribution of gas motions to supernovae (e.g., McKee \\& Ostriker 1977), there was no compelling reason to dismiss the basic cloud-intercloud model in favor of widespread turbulence. Instead, the list of ISM equilibrium ``phases'' was simply enlarged. Supersonic linewidths, long known from {\\it H I} (e.g., McGee, Milton, \\& Wolfe 1966, Heiles 1970) and optical (e.g., Hobbs 1974) studies and also discovered in molecular regions at this time (see Zuckerman \\& Palmer 1974), were thought to represent magnetic waves in a uniform cloud (Arons \\& Max 1975), even though turbulence was discussed as another possibility in spite of problems with the rapid decay rate (Goldreich \\& Kwan 1974, Zuckerman \\& Evans 1974). A lone study by Baker (1973) found large-scale correlations in {\\it H I} emission and presented them the context of ISM turbulence, deriving the number of ``turbulent cells in the line of sight,'' instead of the number of ``clouds.'' Mebold, Hachenberg \\& Laury-Micoulaut (1974) followed this with another statistical analysis of the {\\it H I} emission. However, there was no theoretical context in which the Baker and Mebold et al. papers could flourish given the pervasive models about discrete clouds and two or three-phase equilibrium.  The presence of turbulence was more widely accepted for very small scales. Observations of interstellar scintillation at radio wavelengths implied there were correlated structures (Rickett 1970), possibly related to turbulence (Little \\& Matheson 1973), in the ionized gas at scales down to $10^9$ cm or lower (Salpeter 1969; {\\it Interstellar Turbulence II,} next chapter, this volume). This is the same scale at which cosmic rays ({\\it Interstellar Turbulence II}) were supposed to excite magnetic turbulence by streaming instabilities (Wentzel 1968a). However, there was (and still is) little understanding of the physical connection between these small-scale fluctuations and the larger-scale motions in the cool neutral gas.  Dense structures on resolvable scales began to look more like turbulence after Larson (1981) found power-law correlations between molecular cloud sizes and linewidths that were reminiscent of the Kolmogorov scaling law. Larson's work was soon followed by more homogeneous observations that showed similar correlations (Myers 1983, Dame et al. 1986, Solomon et al. 1987). These motions were believed to be turbulent because of their power-law nature, despite continued concern with decay times, but there was little recognition that turbulence on larger scales could also form the same structures in which the linewidths were measured. Several reviews during this time reflect the pending transition (Dickey 1985, Dickman 1985, Scalo 1987, Dickey \\& Lockman 1990).  Perhaps the most widespread change in perception came when the Infrared Astronomical Satellite (IRAS) observed interstellar ``cirrus'' and other clouds in emission at 100 $\\mu$ (Low et al. 1984). The cirrus clouds are mostly transparent at optical wavelengths, so they should be in the diffuse cloud category, but they were seen to be filamentary and criss-crossed, with little resemblance to ``standard'' clouds. Equally complex structures were present even in IRAS maps of ``dark clouds,'' like Taurus, and they were observed in maps of molecular clouds, such as the Orion region (Bally et al. 1987). The wide field of view and good dynamic range of these new surveys finally allowed the diffuse and molecular clouds to reveal their full structural complexity, just as the optical nebulae and dark clouds did two decades earlier. Contributing to this change in perception was the surprising discovery by Crovisier \\& Dickey (1983) of a power spectrum for widespread {\\it H I} emission that was comparable to the Kolmogorov power spectrum for velocity in incompressible turbulence. CO velocities were found to be correlated over a range of scales, too (Scalo 1984, Stenholm 1984). By the late 1980s, compression from interstellar turbulence was considered to be one of the main cloud-formation mechanisms (see review in Elmegreen 1991).  Here we summarize observations and theory of interstellar turbulence. This first review discusses the dense cool phases of the ISM, energy sources, turbulence theory, and simulations. {\\it Interstellar Turbulence II} considers the effects of turbulence on element mixing, chemistry, cosmic ray scattering, and radio scintillation.  There are many reviews and textbooks on turbulence. A comprehensive review of magnetohydrodynamical (MHD) turbulence is in the recent book by Biskamp (2003), and a review of laboratory turbulence is in Sreenivasan \\& Antonia (1997). A review of incompressible MHD turbulence is in Chandran (2003). For the ISM, a collection of papers covering a broad range of topics is in the book edited by Franco \\& Carraminana (1999). Recent reviews of ISM turbulence simulations are in V\\'azquez-Semadeni et al. (2000), Mac Low (2003), and Mac Low \\& Klessen (2004), and a review of observations is in Falgarone, Hily-Blant \\& Levrier (2003). A review of theory related to the ISM is given by V\\'azquez-Semadeni (1999).  Earlier work is surveyed by Scalo (1987). General discussions of incompressible turbulence can be found in Tennekes \\& Lumley (1972), Hinze (1975), Lesieur (1990), McComb (1990), Frisch (1995), Mathieu \\& Scott (2000), Pope (2000), and Tsinober (2001). The comprehensive two volumes by Monin \\& Yaglom (1975) remain extremely useful. Work on compressibility effects in turbulence at fairly low Mach numbers is reviewed by Lele (1994). Generally the literature is so large that we can reference only a few specific results on each topic; the reader should consult the most recent papers for citations of earlier work.  We have included papers that were available to us as of Dec. 2003. A complete bibliography including paper titles is available at:   http://www.as.utexas.edu/astronomy/people/scalo/research/ARAA/ ",
        "conclusions": "Interstellar turbulence has been studied using power spectra and structure functions of the distributions of radial velocity, emission, and absorption, using statistical properties of line profiles, unsharp masks, and wavelet transforms, one-point probability distribution functions of column densities and velocity centroids, fractal dimensions and multifractal spectra, and various other techniques including the Spectral Correlation Function and Principal Component Analysis, which are applied to spectral line data cubes.  The results are often ambiguous and difficult to interpret. The density of the neutral medium seems to possess spatial correlations over a wide range of scales, possibly from the sub-parsec limit of resolution to hundreds of parsecs. Such correlations probably reflect the hierarchical nature of turbulence in a medium with a very large Reynolds number. The power spectrum of the associated intensity fluctuations is a robust power law with a slope between -1.8 and -2.3 for the Milky Way, LMC, and SMC. A steeper slope has been obtained for Galactic CO using Principle Component Analysis. However, no clear power spectrum or autocorrelation function has been found yet for centroid velocities as a function of spatial lag inside individual clouds, even though they might be expected if the clouds are turbulent.  Statistical properties of the ISM are difficult to recover with only line-of-sight motions in projected cloud maps that have limited spatial extent and substantial noise. Linewidths often seem correlated with region size in an average sense, but there are large variations between different regions and surveys.  At the moment, this relation seems to be dominated by scatter. Because the linewidth-size correlation is analogous to a second order structure function, which is well-defined for incompressible turbulence, the regional variations are difficult to understand in the context of conventional turbulence theory. Certainly a case could be made that ISM turbulence is not conventional: it is not statistically homogeneous, stationary, or isotropic.  Considering the difficulty of observing three-dimensional structure and motion in space, the value of statistical descriptors lies primarily in their ability to make detailed comparisons between observations and simulations. Many such comparisons have been made, but a comprehensive simulation of a particular region including many of the descriptors listed above is not yet available.  Among the many differences between interstellar turbulence and classical incompressible turbulence is the broad spatial scale for energy input in space. In spite of the wide ranging spatial correlations in emission and absorption maps, there is no analogy with classical turbulence in which energy is injected on the largest scale and then cascades in a self-similar fashion to the very small scale of dissipation. In the ISM, energy is injected over a wide range of scales, from kiloparsec disturbances in spiral density waves, shear instabilities, and superbubbles, to parsec-sized explosions in supernovae, massive-star winds and {\\it H II} regions, to sub-parsec motions in low-mass stellar winds and stellar gravitational wakes, to AU-sized motions powered by cosmic ray streaming instabilities. Estimates of power input indicate that stellar explosions are the largest contributor numerically, but this does not mean that other sources are unimportant on other scales.  Energy is also dissipated over a wide range of scales.  Shock fronts, ambipolar and Ohmic diffusion, Landau wave damping, and viscosity in vortex tubes should all play a role. The geometrical structures of the dissipating regions are not well constrained by observations.  Energy is also dissipated in regions with little density sub-structure, as in the ionized gas that produces scintillation or along perturbed magnetic field lines that scatter some cosmic rays. Viscosity and Landau damping can produce heat from the tiny wave-like motions that occur in these regions (see Interstellar Turbulence, Part II), although the nature of this dissipation below the collision mean free path is not understood yet.  Self-gravity makes interstellar turbulence more difficult to understand than terrestrial turbulence. The contribution to ISM motions from self-gravity appears to increase with cloud mass until it dominates above $10^4$--$10^5$ M$_\\odot$. Many selection effects may contribute to this correlation, however. Self-gravity is also important in the smaller clumps that form stars. What happens between these scales remains a mystery.  We would like to know how the distribution function of the ratio of virial mass to luminous mass for ISM clouds and clumps varies with the spatial scales of these structures. What fraction of clouds or clumps are self-gravitating for each mass range? Does the apparent trend toward diminished self-gravity on small scales turn around on the scale of individual protostars?  The balance between solenoidal and compressible energy density in the ISM varies with time. Compressibility transfers energy between kinetic and thermal modes, short-circuiting the cascade of energy in wavenumber space that occurs in a self-similar fashion for incompressible turbulence. Consequently the ISM turbulent power spectra for kinetic and magnetic energy are not known from terrestrial analogies, and there is not even a theoretical or heuristic justification for expecting self-similar or power-law behavior. Numerical simulations make predictions of these quantities, but these simulations are usually idealized and they are always limited by resolution and other numerical effects.  A considerable effort has been put into modelling compressible MHD turbulence under idealized conditions, i.e., without self-gravity and without dispersed and realistic energy sources (e.g., explosions). Some models predict analytically and confirm numerically a Kolmogorov energy spectrum transverse to the mean field when the magnetic energy density is not much larger than the kinetic energy density. These models also predict intermittency properties identical to hydrodynamic turbulence in this trans-field direction. Other models find steeper power spectra, however. One has mostly solenoidal motions on large scales and sheet-like dissipation regions on small scales. Another has more realistic heating and cooling. At very strong fields, the turbulence becomes more restricted to the two transverse dimensions and then the energy spectrum seems to become flatter, as in the Iroshnikov-Kraichnan model.  Overall the situation regarding the energy spectrum of MHD ISM turbulence is unsettled.  Our understanding of the ISM has benefited greatly from numerical simulations. In the 10 or so years since the first simulations of ISM turbulence, numerical models have reproduced most of the observed correlations and scaling relations.  They have demonstrated that supersonic turbulence always decays quickly and concluded that star formation is equally fast, forcing a link between cold clouds and their energy-rich environments. They have also predicted a somewhat universal probability function for density in an isothermal gas, implying that only a small fraction of the gas mass can ever be in a dense enough state to collapse gravitationally -- thereby explaining the low efficiency of star formation.  Simulations have demonstrated that magnetism does not support clouds, prevent collapse, systematically align small clumps, or act like an effective pressure. Magnetism can slow collapsing cores and remove angular momentum on large and small scales, and it may contribute to filamentary structures that control the accretion rate of gas onto protostars. Simulations have also shown that thermal instabilities do not lead to bimodal density distributions as previously believed, although they probably enhance the compressibility of the ISM and might contribute to the turbulent motions.  The high abundance of atomic gas in what would be the thermally unstable regime of static models is easily explained in a turbulent medium.  Simulations have a long way to go before solving some of the most important problems, however. Models are needed that include realistic cooling, ionization balance, chemistry, radiative transfer, ambipolar diffusion, magnetic reconnection, and especially realistic forcing because it is not clear that any of these effects are negligible. Detailed models should be able to form star clusters in a realistic fashion, continuing the progress that has already been made with restricted resolution and input physics. Complete models need to include energy sources from young stellar winds, ionization, and heating, and they should be able to follow the orbital dynamics of binaries. Successful models should eventually explain the stellar initial mass function, mass segregation, the binary fraction and distribution of binary orbital periods, the mean magnetic flux in a star, and the overall star-formation efficiency at the time of cloud disruption.  Realistic simulations of galactic-scale processes are also challenging because they should include a large range of scales and fundamentally different physical processes in the radial and vertical directions. Background galactic dynamics and the possibility of energy and mass exchange with a three-dimensional halo are also important.  The turbulent model of the ISM also raises substantial new questions. Why is the power spectrum of ISM density structure a power law when direct observation shows the ISM to be a collection of shells, bubbles, comets, spiral wave shocks and other pressurized structures spanning a wide range of scales? Do directed pressures and turbulence combine to produce the observed power spectrum or does one dominate to other? Does the input of energy on one scale prevent the turbulent cascade of energy that is put in on a larger scale?  In this respect, is ISM turbulence more similar to turbulent combustion than incompressible turbulence because of the way in which energy is injected?  Our current embrace of turbulence theory as an explanation for ISM structures and motions may be partly based on an overly-simplification of available models and a limitation of observational techniques. This state of the field guarantees more surprises in the coming decade.   {\\it Acknowledgments:} We are grateful to J. Ballesteros-Paredes, J. Dickey, N. Evans, E. Falgarone, A. Goodman, A. Lazarian, P. Padoan, and E. V\\'azquez-Semadeni for helpful comments on Section 2;  to C. Norman for a careful reading of Section 3; to A. Brandenberg, A. Lazarian, P. Padoan, T. Passot, E. V\\'azquez-Semadeni, and E. Vishniac for helpful comments on Section 4; and to A.G. Kritsuk, M-M. Mac Low, E. Ostriker, P. Padoan, and E. V\\'azquez-Semadeni for helpful comments on Section 5. We also thank J. Kormendy for helpful comments on the manuscript."
    },
    "0404/astro-ph0404384_arXiv.txt": {
        "abstract": "{In this paper we present a modified version of the CORS method based on a new calibration of the Surface Brightness function in the Str\\\"omgren photometric system. The method has been tested by means of synthetic light and radial velocity curves derived from nonlinear pulsation models. Detailed simulations have been performed to take into account the quality of real observed curves as well as possible shifts between photometric and radial velocity data. The method has been then applied to a sample of Galactic Cepheids with Str\\\"omgren photometry and radial velocity data to derive the radii and a new PR relation. As a result we find $\\log R = (1.19 \\pm 0.09) + (0.74 \\pm 0.11) \\log P$ (r.m.s=0.07). The comparison between our result and previous estimates in the literature is satisfactory. Better results are expected from the adoption of improved model atmosphere grids. ",
        "introduction": "Classical Cepheids are the cornerstone of the extragalactic distance scale. Thanks to their characteristic Period-Luminosity (PL) and Period-Luminosity-Color (PLC) relations they are traditionally used to derive the distances to Local Group galaxies, and (with the advent of space observations) to external galaxies distant up to about 25 Mpc (targets of a Hubble Space Telescope Key Project, see Freedman et al. 1997, 2001). As primary indicators they are used to calibrate a number of secondary distance indicators (see e.g. Freedman et al. 2001) reaching the region of the so called {\\it Hubble flow} where the Hubble law can be applied and an estimate of the Hubble constant can be derived.\\\\ Moreover, the comparison between Cepheid physical parameters (stellar mass, luminosity, chemical composition) based on evolutionary and pulsation models supplies the unique opportunity to pin point the occurrence of deceptive systematic errors (Bono et al. 2001a; Moskalik 2000) on the Cepheid distance scale.\\\\ In particular radius determinations are important to constrain both the intrinsic luminosity, through the application of the Stefan-Boltzmann law, provided that an effective temperature calibration is available, and the stellar mass, by adopting a Period-Mass-Radius relation (e.g. Bono et al. 2001).\\\\ Many investigations have been devoted during the last decade to the derivation of accurate Period-Radius (PR) relations for Classical Cepheids both from the empirical (see e.g. Laney \\& Stobie 1995; Gieren, Fouqu\\'e, \\& Gomez 1998; Ripepi et al. 1997) and the theoretical (Bono, Caputo, Marconi 1998; Marconi et al. 2003) point of view.\\\\ Empirical Cepheid radii are generally derived either by means of the Baade Wesselink (BW) method (Moffet \\& Barnes 1987, Ripepi et al. 1997; Gieren et al. 1998, just to list a few examples) both in the classical form and in subsequent modified versions, or with interferometric coupled with trigonometric parallaxes techniques (Nordgren, Armstrong \\& German 2000, Lane, Creech-Eakman \\& Nordgren 2002).\\\\ The latter method is more direct and less model dependent but up to now it has been applied only to a limited number of stars. On the other hand, the different versions of the BW technique can be applied to relatively large Cepheid samples but require both accurate photometric and radial velocity data.\\\\ A powerful modification of the BW technique is the so called CORS method (Caccin et al. 1981), which has the advantage of taking into account the whole light curve rather than selecting phase points at the same color (as in the classical BW implementation), but relies on the adoption of an accurate Surface Brightness (SB) calibration.\\\\ Originally Sollazzo et al., (1981) adopted the empirical SB photometric calibration in the Walraven system provided by Pel (1978). More recently Ripepi et al. (1997) modified the method, by adopting the empirical calibration of the reduced surface brightness $F_V$  as a function of $(V-R)$ provided by Barnes, Evans $\\&$ Parson  (1976). This modified version of the CORS method was tested, for different colors selections, through the application to synthetic light and radial velocity curves based on nonlinear convective pulsation models (Ripepi et al. 2000).\\\\ The recent release of new Cepheid data in the Str\\\"omgren photometric system (Arellano-Ferro et al. 1998), and the known sensitivity of intermediate band colors to stellar physical parameters (e.g. gravity and  effective temperature)  suggested us to investigate the possibility of extending the CORS method to the Str\\\"omgren filters.\\\\ To this purpose we have derived, in this system, a SB calibration based on model atmosphere tabulations. In this paper we present a modified version of the CORS method based on this new calibration and the application to a sample of Galactic Cepheids.\\\\ The organization of the paper is the following: in Sect. 2 we summarize the assumptions and the philosophy of the traditional CORS method; in Sect. 3 we introduce the modified CORS method based on the new SB calibration; in Sect. 4 we test the new method by means of pulsation models; in Sect. 5 we apply the method to a sample of Galactic Cepheids and  present the comparison with the literature; in Sect. 6   our final results concerning  the PR relation for Galactic Cepheids are shown and the theoretical fit of observed light and radial velocity variations for Cepheid Y Oph is used as an additional check. Some final remarks close the paper. ",
        "conclusions": ""
    },
    "0404/astro-ph0404047_arXiv.txt": {
        "abstract": "Recent {\\it Chandra} and XMM--{\\it Newton} surveys have uncovered a large fraction of the obscured AGN responsible of the hard X--ray background. One of the most intriguing results of extensive programs of follow--up observations concerns the optical and near--infrared properties of the hard X--ray sources counterparts. More specifically, for a significant fraction of hard X--ray obscured sources the AGN responsible of the high X--ray luminosity remains elusive over a wide range of wavelengths from soft X--rays to near--infrared. This very observational result opens the possibility to investigate the host of bright obscured quasars in some detail. Here we briefly report on some preliminar results obtained for a small sample of elusive AGN in the {\\tt HELLAS2XMM} survey. ",
        "introduction": "Hard X-ray surveys represent an efficient probe to unveil the super-massive black hole ({\\tt SMBH}) accretion activity, which is recorded in the cosmic X-ray background ({\\tt XRB}) spectral intensity. The advent of hard X-ray (i.e., 2--10~keV) imaging instruments, from {\\tt ASCA} to {\\it BeppoSAX} and, more recently, {\\it Chandra} and \\xmm, has provided a dramatic advance in the field of X--ray surveys.  The combination of deep/ultra-deep X--ray surveys with {\\it Chandra} (Alexander et al. 2003; Giacconi et al. 2002) and the shallower surveys with \\xmm\\  (Baldi et al. 2002) has allowed to resolve \\hbox{$\\approx$~80\\%} of the {\\tt XRB} in the \\hbox{2--10~keV} band and to unveil classes of cosmic sources which were previously unknown or marginally represented by a few ambiguous and sparse cases.  Within this context, the {\\tt HELLAS2XMM} survey plays an important role. Using suitable \\xmm\\ archival observations, this project aims at covering 4 square degrees of sky down to X--ray fluxes of the order of $10^{-14}$ \\cgs, sampling the bright tail of the X--ray luminosity function. The key scientific issue is a solid estimate of the luminosity function and evolution of the obscured AGN responsible for a large fraction of the {\\tt XRB}. In order to fulfill the goals of such an ambitious objective, the spectroscopic identification of a large number of hard X--ray selected sources is mandatory. Unfortunately the identification of a sizable fraction of optically faint obscured AGN is already challenging the capabilities of the 8--10 m class telescopes calling for alternative approaches based on multiband optical photometry, detection of redshifted iron K$\\alpha$ lines, or statistical methods. Some of the results obtained by our group in these regards,  based on the multiwavelength observations of the {\\tt HELLAS2XMM} 1 degree field, will be briefly reported. ",
        "conclusions": ""
    },
    "0404/astro-ph0404271_arXiv.txt": {
        "abstract": "{The rate coefficient for radiative and dielectronic recombination of berylliumlike magnesium ions was measured with high resolution at the Heidelberg heavy-ion storage ring TSR. In the electron-ion collision energy range 0--207 eV resonances due to $2s \\to 2p$ ($\\Delta N = 0$) and $2s \\to 3l$ ($\\Delta N=1$) core excitations were detected. At low energies below 0.15 eV the recombination rate coefficient is dominated by strong $1s^2\\,(2s\\,2p\\,\\,^3P)\\,7l$ resonances with the strongest one occuring at an energy of only 21 meV. These resonances decisively influence the \\ion{Mg}{ix} recombination rate coefficient in a low temperature plasma. The experimentally derived \\ion{Mg}{ix} dielectronic recombination rate coefficient ($\\pm 15\\%$ systematical uncertainty) is compared with the recommendation by Mazzotta et al.\\ (1998, A\\&AS, 133, 403) and the recent calculations by Gu (2003, ApJ, 590, 1131) and by Colgan et al.\\ (2003, A\\&A, 412, 597). These results deviate from the experimental rate coefficient by 130\\%, 82\\% and 25\\%, respectively, at the temperature where the fractional abundance of \\ion{Mg}{ix} is expected to peak in a photoionized plasma. At this temperature a theoretical uncertainty in the $1s^2\\,(2s\\,2p\\,\\,^3P)\\,7l$ resonance positions of only 100 meV would translate into an uncertainty of the plasma rate coefficient of almost a factor 3. This finding emphasizes that an accurate theoretical calculation of the \\ion{Mg}{ix} recombination rate coefficient from first principles is challenging. ",
        "introduction": "\\label{sec:intro}  For the accurate calculation of the ionization equilibrium in astrophysical plasmas, rate coefficients for the population and depopulation of the various ion charge states have to be known precisely. To date most of the required rate coefficients are derived from theoretical calculations, and, hence, experimental benchmarks are required for testing and improving the theoretical methods. This is especially true for recombination in photoionized plasmas, which occurs at relatively low plasma temperatures of only a few electron volts. At such low temperatures the recombination rate coefficient, depending on the ion under consideration, can strongly be influenced by the existence of low-energy dielectronic recombination (DR) resonances that are difficult to theoretically predict with sufficient accuracy.  Experimentally derived plasma rate coefficients were previously published for the recombination of \\ion{C}{iv} \\citep{Schippers2001c}, \\ion{O}{vi} \\citep{Boehm2003a}, \\ion{Ti}{v} \\citep{Schippers1998}, \\ion{Ni}{xviii} \\citep{Fogle2003b}, \\ion{Ni}{xxvi} \\citep{Schippers2000b} and \\ion{Fe}{xviii} -- \\ion{Fe}{xxii} \\citep{Savin1997,Savin1999,Savin2002a,Savin2002c,Savin2003a}. Here the \\ion{Mg}{ix} recombination rate coefficient derived from experimental measurements at a heavy ion storage ring is provided.    The approximate temperature ranges where berylliumlike magnesium forms in photoionized and in collisionally ionized plasmas can be obtained from the work of \\citet{Kallman2001} who calculated the fractional abundances of ions in plasmas for a variety of physical conditions. For photoionized plasmas they find that the fractional \\ion{Mg} {ix} abundance peaks at an ionization parameter of $\\log\\zeta = 0.9$ corresponding to a temperature of about 2.8 eV. The `photoionized zone' may be defined as the range of temperatures where the fractional abundance of a given ion exceeds 10\\% of its peak value. For \\ion{Mg}{ix} this corresponds to the temperature range 2--13~eV. Using the same criterion and the results of \\citet{Kallman2001} for coronal equilibrium the \\ion{Mg}{ix} `collisionally ionized zone' is estimated to extend over the temperature range 60--170~eV. It should be kept in mind that these temperature ranges are only indicative. In particular, they depend on the accuracy of the atomic data base used by \\citet{Kallman2001} and, in case of the photoionization zone, on the assumed $1/E$ energy dependence of the ionizing radiation. Nevertheless, the above given temperature ranges will be used in the discussion below.   From the present measurements it is found that the \\ion{Mg}{ix} recombination rate coefficient in the photoionization zone is decisively influenced by the presence of a strong DR resonance at an electron-ion collision energy of 21 meV. Any theoretical calculation aiming at an accurate \\ion{Mg}{ix} rate coefficient will have to predict this resonance's position with an error of less than a few meV. Presently available atomic-structure computer codes are generally not capable of providing results with such an accuracy ab initio. In the standard theoretical approaches the situation is often improved by using spectroscopically observed target energies for the atomic structure \\citep{Colgan2003a,Gu2003b}. As will be shown below, a 100 meV uncertainty on the theoretical resonance position would translate into an uncertainty of the \\ion{Mg}{ix} DR rate coefficient of a factor 2.7 at the temperature where the \\ion{Mg}{ix} abundance is expected to peak in photoionization equilibrium. ",
        "conclusions": "\\label{sec:dis}   In Figure \\ref{fig:Mg8plasma} the present experimentally derived DR+TR rate coefficient is compared with the recommendation of \\citet{Mazzotta1998} and with the recent calculations of \\citet{Gu2003b} and \\citet{Colgan2003a}. The recommended \\ion{Mg}{ix} DR rate coefficient of \\citet{Mazzotta1998} deviates by up to 130\\% from the experimental result in the temperature range 2--13~eV where \\ion{Mg}{ix} exists in photoionized plasmas. This rather large discrepancy questions the usefulness of the recommended \\ion{Mg}{ix} DR rate coefficient for the modeling of photoionized plasmas. In the temperature range 60-170~eV (collisionally ionized zone) the recommended \\ion{Mg}{ix} DR rate coefficient of \\citet{Mazzotta1998} agrees within 11\\%, i.\\,e., within the 15\\% experimental uncertainty, with the experimentally derived one.  In this temperature range the theoretical rate coefficient of \\citet{Gu2003b} shows very good agreement with the experimentally derived rate coefficient. The difference between the two curves in this range is less than 7\\%. At lower temperatures, however, the discrepancy is up to 99\\% at 2~eV. At these low temperatures where \\ion{Mg}{ix} exists in photoionized plasmas, the theoretical result of \\citet{Colgan2003a} agrees better with the experimental rate coefficient. The deviation is less than 35\\% for $k_BT > 2$~eV, less than 30\\% for $k_BT > 15$~eV, and less than 15\\% for $k_BT > 100$~eV. At the temperature 2.8 eV  where the fractional abundance of \\ion{Mg}{ix} is expected to peak in a photoionized plasma \\citep{Kallman2001} the results of \\citet{Gu2003b} and \\citet{Colgan2003a} deviate from the experimental rate coefficient by 82\\% and 25\\%, respectively.  The rather large low-temperature deviation of the theoretical rate coefficient of \\citet{Gu2003b} from the experimentally derived curve is most probably due to an inaccurate calculation of the DR resonance energies below 0.15~eV (Figure \\ref{fig:Mg8hires}). At these very low energies the DR rate coefficient is very sensitive to slight variations of resonance positions. This is highlighted in Figure \\ref{fig:Mg8shift} that displays the effect of hypothetical shifts of the low-energy resonance positions by $\\pm 50$ meV and $\\pm 100$~meV with respect to the tabulated values (Table \\ref{tab:fit}) on the \\ion{Mg}{ix} DR rate coefficient. Obviously, a $\\pm 100$~meV uncertainty of the low-energy $1s^2\\,(2s\\,2p\\,\\,^3P)\\,7l$ resonance positions translates into an uncertainty of a factor 2.7 of the plasma rate coefficient in the temperature range where the fractional abundance of \\ion{Mg}{ix} peaks in a photoionized gas.   \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{0380fig6} \\caption{\\label{fig:Mg8shift}Impact of a shift of the low energy $1s^2\\,(2s\\,2p\\,\\,^3P)\\,7l$ resonances displayed in Figure \\ref{fig:Mg8hires} on the \\ion{Mg}{ix} DR plasma rate-coefficient. The full curve is the experimental result, the long-dashed curves correspond to resonance shifts of $\\Delta E_\\mathrm{res} = \\pm 50$~meV, and the short-dashed  curves to a shift of $\\Delta E_\\mathrm{res} = \\pm 100$~meV. Thereby, the lower curves correspond to resonance shifts to lower energies. The vertical dash-dotted line marks the temperature 2.8~eV where the fractional abundance of \\ion{Mg}{ix} peaks in a photoionized gas \\citep{Kallman2001}. The horizontal dash-dotted lines indicate the factor 2.7 uncertainty of the plasma rate coefficient at this temperature that would be introduced by a $\\pm 100$~meV uncertainty of the low energy resonance positions.} \\end{figure}   These findings demonstrate that an accurate calculation of DR rates for ions that exhibit DR resonances close to zero energy is challenging. In case of more complex ions this task is certainly beyond the capabilities of present-day atomic structure codes as is exemplified by a recent combined theoretical and experimental DR study of argon-like Sc$^{3+}$ \\citep{Schippers2002a}. Recombination experiments at heavy-ion storage rings are certainly required for guiding the future development of the theoretical methods, especially in the case of low temperature DR rate coefficients for complex ions."
    },
    "0404/astro-ph0404487_arXiv.txt": {
        "abstract": "We analyze the properties of quasar variability using repeated SDSS imaging data in five UV-to-far red photometric bands, accurate to 0.02 mag, for $\\sim$13,000 spectroscopically confirmed quasars. The observed time lags span the range from 3 hours to over 3 years, and constrain the quasar variability for rest-frame time lags of up to two years, and at rest-frame wavelengths from 1000\\AA\\ to 6000\\AA. We demonstrate that $\\sim$66,000 SDSS measurements of magnitude differences can be described within the measurement noise by a simple function of only three free parameters. The addition of POSS data constrains the long-term behavior of quasar variability and provides evidence for a turn-over in the structure function. This turn-over indicates that the characteristic time scale for optical variability of quasars is of the order 1 year. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404164_arXiv.txt": {
        "abstract": "{ Thick disks are faint and extended stellar components found around several disk galaxies including our Milky Way. The Milky Way thick disk, the only one studied in detail, contains mostly old disk stars ($\\approx\\!10$\\,Gyr), so that thick disks are likely to trace the early stages of disk evolution. Previous detections of thick disk stellar light in external galaxies have been originally made for early-type, edge-on galaxies but detailed 2D thick/thin disk decompositions have been reported for only a scant handful of mostly late-type disk galaxies. We present in this paper for the first time explicit 3D thick/thin disk decompositions characterising the presence and properties (\\eg scalelength and scaleheight) for a sample of eight lenticular galaxies by fitting 3D disk models to the data. For six out of the eight galaxies we were able to derive a consistent thin/thick disk model. The mean scaleheight of the thick disk is 3.6 times larger than that of the thin disk. The scalelength of the thick disk is about twice, and its central luminosity density between 3-10\\% of, the thin disk value. Both thin and thick disk are truncated at similar radii. This implies that thick disks extend over fewer scalelengths than thin disks, and turning a thin disk into a thick one requires therefore vertical but little radial heating. All these structural parameters are similar to thick disk parameters for later Hubble-type galaxies previously studied. We discuss our data in respect to present models for the origin of thick disks, either as pre- or post-thin-disk structures, providing new observational constraints. ",
        "introduction": "\\label{introduction} \\begin{table*} \\begin{center} {\\normalsize \\begin{tabular}{ l  c c l  r@{.}l  c  c  c r@{.}l c } \\hline \\rule[+0.4cm]{0mm}{0.0cm} Galaxy &RA &DEC &RC3 &\\multicolumn{2}{c}{T} &Diam. &$v_{\\sun}$ &$v_{\\rm vir}$ &\\multicolumn{2}{c}{D} &$M^{\\scriptscriptstyle 0}_{\\rm B}$ \\\\[+0.1cm] &\\multicolumn{2}{c}{(2000.0)} &type &\\multicolumn{2}{c}{} &[\\ \\arcmin\\ ] &[\\ \\kms] &[\\ \\kms] &\\multicolumn{2}{c}{[Mpc]} &[mag]\\\\ \\rule[-3mm]{0mm}{5mm}{\\scriptsize{\\raisebox{-0.7ex}{\\it (1)}}} &{\\scriptsize{\\raisebox{-0.7ex}{\\it (2)}}} &{\\scriptsize{\\raisebox{-0.7ex}{\\it (3)}}} &\\hspace*{0.1cm}{\\scriptsize{\\raisebox{-0.7ex}{\\it (4)}}} &\\multicolumn{2}{c}{{\\scriptsize{\\raisebox{-0.7ex}{\\it (5)}}}} &{\\scriptsize{\\raisebox{-0.7ex}{\\it (6)}}} &{\\scriptsize{\\raisebox{-0.7ex}{\\it (7)}}} &{\\scriptsize{\\raisebox{-0.7ex}{\\it (8)}}} &\\multicolumn{2}{c}{{\\scriptsize{\\raisebox{-0.7ex}{\\it (9)}}}} &{\\scriptsize{\\raisebox{-0.7ex}{\\it (10)}}} \\\\ \\hline\\hline \\\\[-0.2cm] NGC\\,1596   &04 27 38.1&$-$55 01 40 &.LA..*/ &$-$2&0 &3.9 &1510 &1229   &17&1&$-$19.3\\\\ NGC\\,2310   &06 53 54.0&$-$40 51 45 &.L..../ &$-$2&0 &4.1 &1187 & 943   &13&1&$-$18.6\\\\ ESO\\,311-012&07 47 34.1&$-$41 27 08 &.S..0?/ &   0&0 &3.7 &1131 & 893   &12&4&$-$20.0\\\\ NGC\\,3564   &11 10 36.4&$-$37 32 51 &.L...*/ &$-$2&0 &1.7 &2812 &2639   &36&7&$-$20.2\\\\ NGC\\,3957   &11 54 01.5&$-$19 34 08 &.LA.+*/ &$-$1&0 &3.1 &1687 &1607   &22&3&$-$19.0\\\\ NGC\\,4179   &12 12 52.6&$+$01 17 47 &.L..../ &$-$2&0 &4.1 &1248 &1277   &17&7&$-$19.6\\\\ NGC\\,4521   &12 32 47.7&$+$63 56 21 &.S..0.. &   0&0 &2.6 &2500 &2767   &38&4&$-$20.2\\\\ NGC\\,5047   &13 15 48.5&$-$16 31 08 &.L..../ &$-$2&0 &2.7 &6330 &6292   &87&4&$-$21.7\\\\ \\hline \\end{tabular} } \\caption[]{Global parameters of the observed lenticular galaxies: {\\scriptsize{\\it (1)}} Principal name, {\\scriptsize{\\it (2)}} right ascension, {\\scriptsize{\\it (3)}} declination,  {\\scriptsize{\\it(4)}} RC3 coded Hubble-type, and the {\\scriptsize{\\it(5)}} Hubble parameter T are taken from \\cite{rc3}. The {\\scriptsize{\\it(6)}} diameter in arcminutes, the {\\scriptsize{\\it(7)}} heliocentric radial velocities, and the B-Band absolute magnitude {\\scriptsize{\\it(10)}} are taken from  LEDA. According to the heliocentric radial velocities corrected for the Local Group infall into the Virgo cluster {\\scriptsize{\\it(8)}} from LEDA, we estimated the {\\scriptsize{\\it (9)}} distances following the Hubble relation with the Hubble constant from the HST key project of $H_{0}\\!=\\!72$ km s$^{-1}$Mpc$^{-1}$ \\cite[]{hst_h0}. \\label{sample} } \\end{center} \\end{table*} The knowledge of the detailed distribution of stars in galaxies is of fundamental importance to address the formation and evolution of those systems. To a first approximation, a disk galaxy can be described by a set of distinct stellar entities: a disk population, a bulge component, and a stellar halo. Deep surface photometry of external early-type galaxies \\cite[]{bursteinI,bursteinII,bursteinIII,tsikoudi1979,tsikoudi1980} and later elaborate measurements in our own Galaxy \\cite[]{gilmore1983} revealed the need for an additional component of stars. This was called `thick disk' \\cite[]{bursteinIII}, since it exhibiting a disk-like distribution with larger scaleheight compared to the inner, dominating `thin disk'.  There are three distinct families of hypotheses for its creation. The first group considers the thick disk as a separate entity produced in an early phase of enhanced star formation during the initial proto-galactic collapse \\cite[i.e.~the ELS scenario,][]{els,gilmore1984,burkert1992}. Another family of models regards the thick disk as an extension (by dynamical heating) of the thin disk. They assume that after the initial collapse all gas settles down into the galactic plane and starts forming stars. On this thin stellar disk a variety of constant or violent heating mechanisms could act: spiral density waves \\cite[]{barbanis1967,carlberg1985}, encounters with giant molecular clouds \\cite[e.g.][]{spitzer1953,lacey1984}, scattering by massive black holes \\cite[]{lacey1985}, energy input by accretion of satellite galaxies \\cite[]{carney1989,quinn1993,velazquez1999, aguerri2001}, or bar bending instabilities \\cite[]{raha1991}. For example, \\cite{gnedin2003} recently used N-body simulations to show that tidal heating in a cluster is sufficient to thicken stellar disks by a factor of 2-3. This kinematic heating and vertical expansion will lead to a significant morphological transformation of a normal spiral galaxy into a lenticular. The third model suggests that thick disks are mostly made of debris material from accreted satellites. Recent cosmological N-body+SPH galaxy formation models of \\cite{abadi2003} locating the thick disk formation before $z\\!\\approx\\!1$ find that more than half of the thick disk stars are actually tidal debris from disrupted satellites. Therefore the thick disk is not a former thin disk thickened by a minor merger. To decide which of these hypotheses could explain the thick disk phenomenon best we need first a more general and complete statistic of thick disk properties. Naturally, these are rather global ones for external galaxies whereas our particular position in the Milky Way makes it possible to determine much finer details.  Since the work of Tsikoudi and Burstein \\cite[]{bursteinI,bursteinII, bursteinIII,tsikoudi1979,tsikoudi1980} it appears well known that thick disks are quite common in \\s0 galaxies. However, none of the more recent detections, except for two galaxies in \\cite{degrijs1996} and a short remark in \\cite{degrijs1997}, quantifying detailed parameters such as the ratio of thick to thin disk scaleheight or scalelength, is actually made in \\s0 galaxies. All these galaxies are of later Hubble type. In addition, we have not found a detailed 2D thin/thick disk decomposition for any \\s0 galaxy in the literature. Subsequent numerical decompositions dealing with \\s0 galaxies after the pioneering work in the early 80's treated the thick disk either as an outer flattened but exponential halo \\cite[for NGC\\,4452 and NGC\\,4762:][]{hamabe1989}, or as a spheroidal bulge component \\cite[for NGC\\,1381:][]{decarvalho1987} \\cite[for NGC\\,3115:][]{capaccioli1987,silva1989}.  The detections of possible halo or thick disk stellar light in disk galaxies of later Hubble type have been made in a scant handful of mostly nearby edge-on galaxies \\cite[\\eg for ESO\\,342-017, IC\\,5249, NGC\\,891, NGC\\,4565, NGC\\,5907, and NGC\\,6504:][]{vdk81,vdk81b,shaw1989,morrison1994,vandokkum1994,sackett1994, morrison1997,naeslund1997,lequeux1998,zheng1999,neeser2002}. Quite recently, \\cite{dalcanton2002} suggest the detection of extended, ubiquitous thick disks in a large sample of late-type, edge-on galaxies by means of multi-colour imaging. However, their thick disks are solely detected by vertical colour gradients for which dust extinction complicates the interpretation. In addition, their vertical colour profiles, especially $(R-K)$, typically extend out to only very few vertical disk scaleheights. At those z-heights where they attribute the red colour to an additional component the thin disk may still be dominant over a potential thick disk and even determine the measured colour.  In this paper, we analyse a set of eight edge-on \\s0 galaxies using the classical approach for the identification of thick disks in external galaxies: the need for an additional disk component when attempting to fit single disk models to the light distribution in a deep image. Thereby we characterise the presence and properties (scalelengths, -heights, and central surface brightnesses) of thick disks by directly measuring their structure. ",
        "conclusions": "\\label{Discussion} \\subsection{Thick disks: discrete or continuous?} A key question is: {\\sl How to describe a thick disk in general? } With reference to the proposed different formation scenarios described in  Sect.~\\ref{introduction} we can assume that the disk component is either characterised as a superposition of two discrete and independent isothermal disk systems (as done here), or built from the contribution of multiple velocity-dispersion components \\cite[e.g.][]{wielen1992,dove1993}. These two approaches on how to treat a thick disk seem to be incompatible. Therefore any derived parameters for the two-disks model appear useless for the multi-component disk. In addition, the latter seems to be superior since it is consistent with the model of continuous disk heating leading to the well known observable age-velocity dispersion relation \\cite[]{wielen1977}.  However, \\cite{majewski1993} already states in his review that detailed studies of the spatial and kinematical distribution of stars in the Milky Way do not make it possible to decide if the thick disk is a discrete component or a more continuous sequence of stellar populations. Recently, \\cite{nissen2003} concluded that the latest studies \\citep[e.g.][]{bensby2003} argue for the separate entity picture. This confirms our result that 75\\% of the chosen \\s0 galaxies are well described with a distinct, two-component thin/thick disk system. Note that these results exclude all thick disk formation scenarios based solely on heating. Especially the elemental abundance trends found by \\cite{bensby2003} favour a merger scenario where a satellite galaxy either merges with the parent galaxy or sheds significant amounts of its material to form the thick disk as proposed by \\cite{abadi2003}.  However, as already noted by \\cite{jacobi1994} it does not seem possible to distinguish between our simple model and more sophisticated ones. Therefore the very good description with the applied discrete two-component system for our galaxies is only a first step towards studying the luminosity distribution of external galaxies. In addition, one has to keep in mind that our model only fits well in a restricted radial range (only out to $2.2$\\hn$-3.9$\\hn) where most of the inner part is ``hidden'' by the central bulge, bar, or ring components. One key to the nature of the thick disk may lie in studying the outer parts where the breaks in the radial profiles of both thin and thick disk should provide an additional constraint. Any vertical colour gradients could provide additional information but the only available colour map for our sample \\cite[NGC\\,3957 in][]{pohlen2001} suffers from low S/N and is not conclusive. \\subsection{Are thick disks of \\s0s exceptional ?} The fact that the range of thick disk parameters for all known \\s0 galaxies are not too different from those of late-type galaxies, even compared to a low surface brightness Scd galaxy, is especially surprising, since  at first glance one expects \\s0 galaxies to posses more prominent thick disks. However, in terms of the thick-to-thin-disk scaleheight ratio our values agree well with those derived for all other galaxies (\\cf Table \\ref{comparison}). Even taking into account the central surface brightness or central luminosity density the comparison yields similar values. Does this point to a general formation process for thick disks independent of bulge-to-disk ratio and Hubble type ? At this stage we are not able to answer this question. Although the numbers of galaxies used for the comparison are the same, all literature values come from different sources, using sometimes very different methods and models to derive their parameters. In the case of NGC\\,4565 it is obvious that depending on the model, one can find an even larger variety of $\\zkdzn$-values than the full range for our six galaxies. In addition, as pointed out by \\cite{knapen1991}, and again discussed in detail by \\cite{pohlen2001}, fitting disk-model components to surface brightness data of edge-on galaxies is a delicate business, and comparing the results of different authors can be misleading. To overcome this problem one has to extend this survey of thick disk parameters also to late-type galaxies. To reduce the inevitable influence of their prominent dust lanes this has to be done in the near-infrared. After applying the same fitting method one is able to address the questions if thick disks are really common around all Hubble-types, and if their parameters are really similar, as suggested here. This would entail a common formation scenario for thick disks independent of their normal evolution along the Hubble sequence.  Of our galaxies with successful thick disk models there are only two for which there are rotational velocity measurements in the literature. Therefore any correlation between mass and thick disk parameters could not be derived. Consequently it is still unclear if the mass, as a galaxy characteristic, is related to the thick disk parameters. \\subsection{Thick disk scalelengths} We find significantly larger scalelength ratios ($\\hkdhn\\!=\\!1.9$) for the thick disks of our \\s0 galaxies than in the literature ($\\hkdhn\\!\\ltsim\\!1.4$). Is this a possible distinction in thick disk parameters between late-type and our early-type galaxies ? Again, differences in fitting methods, especially for the scalelength in edge-on galaxies, could be responsible for this apparent disagreement. In particular, fitting profiles with clear breaks using a single disk (with an infinitely sharp truncation or infinite exponential as done here) will entail systematic errors (\\cf \\sec\\ref{breaks}), therefore this question could only be addressed unambiguously applying exactly the same method to late-type galaxies.  However, although the exact number is uncertain the scalelength of the thick disk is without doubt systematically larger than that of the thin disk. Except in one case this is also true for all the literature values. Does this larger scalelength imply a different formation process or even restrict the available formation scenarios ? At a first glance the different scalelengths contradict a dynamical heating scenario, since this should only alter the vertical distribution. However, assuming two distinct disk components with radially constant thickness the velocity dispersion scales with radius \\cite[]{vdk86}. It depends therefore on the exact way (radial distribution) the proposed mechanisms (\\cf\\sec\\ref{introduction}) dispenses the energy in the vertical motions of the stars. In the case of heating the disk by satellite accretion, the N-body simulations of \\cite{quinn1993} show that the scalelength of the disk is nearly unchanged (even slightly smaller than the initial scalelength) for the inner parts (out to $3h_{\\rm ini}$). Due to the migration of material outwards in radius, at larger radius (out to $\\approx\\!6h_{\\rm ini}$) the final disk shows a second shallower component (\\cf their Fig.~4). However, \\cite{aguerri2001} find in their N-body simulation a global outward transport of disk material leading to a general increase of the disk scalelength of 10\\% to  60\\%. Although this increase is below our measured scalelength ratios, the latter seems to coincide in general with observations of larger thick disk scalelengths. One has to keep in mind that in the satellite accretion scenario a pre-existing thin disk gets heated. The thin disk to be observed today must be rebuilt out of gas remaining after the merger process. It is not guaranteed that this thin disk will have the same scalelength as its predecessor since we do not have a unique explanation of how the disk scalelength is determined out of an initial (or new) gas distribution \\cite[cf.][]{pohlen2000}.  Looking from a slightly different angle, one can also try to apply the larger thick disk scalelength as a general argument against an internal heating scenario. Assuming infinitely exponential disks, the larger scalelength implies a larger angular momentum content of the thick disk. Any valid heating mechanism must therefore add angular momentum, which is not the case for the internal heating scenarios. However, note that the disks are not at all infinitely exponential but show clear breaks in their profiles (\\cf\\sec\\ref{breaks}). The disks exhibit the break at roughly similar radial radius (\\eg\\fig\\ref{n5047}). In this sense the thick disk is truncated ``earlier'' (in respect to its scalelength) than the thin disk. In addition to a probably lower rotational velocity this could add up again to a similar angular momentum content for thick and thin disk, ruling out the whole argument against the heating scenarios. The key point here is again the origin of these breaks in the radial profiles. \\subsection{Residuals} In addition to comparing thick disks across a large range of characteristic parameters of galactic disks, fitting and subtracting a combined disk yields valuable information on the structure and size of galactic bulges. As discussed in Appendix \\ref{individual} the residual images highlight the deviation of the galaxy from the disk model. Structures only faintly indicated in the profiles become obvious. One important result related to this is that for our sample of \\s0 galaxies none of the present bulges could be described as a traditional early-type $R^{1/4}$ bulge in agreement with \\cite{balcells2003}. However, a detailed structural analysis of the remaining central structures has to be done with great care. As discussed in \\sec\\ref{method}, choosing the vertical model is not a unique process and therefore any variation of $f(z)$ for the thick and thin disk could alter the shape of the disk profile in the inner bulge parts. Especially for galaxies where the bulge component along the minor axis extends vertically above the disk (ESO\\,311-012 and NGC\\,3957) this could be even more unconstrained."
    },
    "0404/astro-ph0404352_arXiv.txt": {
        "abstract": "If ultra-high energy cosmic rays (UHECRs) are accelerated at astrophysical point sources, the identification of such sources can be achieved if there is some kind of radiation at observable wavelengths that may be associated with the acceleration and/or propagation processes. No radiation of this type has so far been detected or at least no such connection has been claimed. The process of photopion production during the propagation of UHECRs from the sources to the Earth results in the generation of charged and neutral pions. The neutral (charged) pions in turn decay to gamma quanta and electrons that initiate an electromagnetic cascade in the universal photon background. We calculate the flux of this gamma radiation in the GeV-TeV energy range and find that for source luminosities compatible with those expected from small scale anisotropies in the directions of arrival of UHECRs, the fluxes can be detectable by future Cerenkov gamma ray telescopes, such as VERITAS and HESS, provided the intergalactic magnetic field is not larger than $\\sim 10^{-10}$ Gauss and for source distances comparable with the loss length for photopion production. ",
        "introduction": "The discovery of the cosmic microwave background radiation (CMBR) \\cite{penzias} has been a major breakthrough not only for cosmology but for cosmic ray physics as well. Soon after its discovery, it was shown that cosmic rays would suffer severe energy losses due to the inelastic process of photopion production, provided the energy of cosmic rays being higher than the kinematic threshold for this process. This effect was predicted to result in a flux suppression at energies around $10^{20}$~eV \\cite{greisen66,zk66} for a homogeneous distribution of the sources, something which is now known as the GZK feature. Despite the efforts dedicated to the experimental detection of this feature, to date we do not have a definite answer to whether the GZK suppression is there or not. The two largest detectors currently operating, AGASA and HiRes, give results which appear discrepant, although the statistics of events does not allow us to infer a definitive conclusion on this issue \\cite{demarco1}. The detection of the GZK feature in the spectrum of UHECRs would be the clearest proof that UHECRs are generated at extragalactic sources.  Despite the efforts in building several large experiments for the detection of particles with the highest energies in the cosmic ray spectrum, the nature of the sources of these particles and the acceleration processes at work are still unknown. No trivial counterpart has been identified by any of the current experiments: there is no significant association with any large scale local structure nor with single sources. The lack of counterparts is particularly puzzling at the highest energies, where the loss length of the detected particles is small enough that only a few sources can be expected to be located within the error box of current experiments.  It has been recently proposed that a statistically significant correlation may exist between the arrival directions of UHECRs with energy above $2\\times 10^{19}$ eV and the spatial location of BL Lac objects, with redshifts larger than 0.1 \\cite{tkachev1,tkachev2} (but see also \\cite{efs} for a criticism and \\cite{ttreply} for the reply). Since at these energies the loss length is comparable with the size of the universe, no New Physics would be needed to explain the possible correlation. On the other hand, since no BL Lac is known to be located close to the Earth, the spectrum of UHECRs expected from these sources would have a very pronounced GZK cutoff.  The potential for discovery of the sources has recently improved after the identification of a few doublets and triplets of events clustered on angular scales comparable with the experimental angular resolution of AGASA. A recent analysis of the combined results of most UHECR experiments \\cite{uchihori} revealed 8 doublets and two triplets on a total of 92 events above $4\\times 10^{19}$ eV (47 of which are from AGASA). The recent HiRes data do not show evidence for significant small angle clustering, but this might well be the consequence of the smaller statistics of events and of the energy dependent acceptance of this experiment: one should remember that in order to reconstruct the spectrum of cosmic rays, namely to account for the unobserved events, a substantial correction for this energy dependence need to be carried out in HiRes (the acceptance is instead a flat function of energy for AGASA at energies above $10^{19}$ eV). This correction however does not give any information on the distribution of arrival directions of the {\\it missed} events.  The statistical significance of these small scale multiplets of events have recently been questioned in \\cite{finley2003}. If the appearance of these multiplets in the data will be confirmed by future experiments as not just the result of a statistical fluctuation or focusing in the galactic magnetic field \\cite{harari}, then the only way to explain their appearance is by assuming that the sources of UHECRs are in fact point sources. This would represent the first true indication in favor of astrophysical sources of UHECRs, since the clustering of events in most top-down scenarios for UHECRs seems unlikely. A clear identification of the GZK feature in the spectrum of UHECRs would make the evidence in favor of astrophysical sources even stronger.  In some recent work \\cite{blasi2003,will,tom,dubovsky,fodor} it has been shown that the small scale anisotropies can be used as a powerful tool to measure the number density of sources of UHECRs. In particular, in \\cite{blasi2003} the number density of sources that best fits the observed number of doublets and triplets was found to be of order $10^{-6}-10^{-5}$ $\\rm Mpc^{-3}$, with an average luminosity per source of $2\\times 10^{42}-2\\times 10^{43}\\rm erg~s^{-1}$ at energies in excess of $10^{19}$ eV. The extrapolation of these numbers down to GeV energies are strongly dependent upon the injection spectrum, assumed here in such a way to fit the observed spectrum of cosmic rays above $10^{19}$ eV.  Alternative estimates of the density of sources of UHECRs recently appeared in the literature: the number obtained in \\cite{yoshi} is in fair agreement with the result of \\cite{blasi2003}, while more complex is the comparison with the case investigated in \\cite{sigl1,ensslin}. In \\cite{sigl1}, a source density of $\\sim 10^{-4}-10^{-3} \\rm Mpc^{-3}$ was derived, while the large scale isotropy of UHECRs was used to infer a local field strength of $\\sim 0.1 \\mu G$ in the local supercluster. This conclusion was mainly due to the fact that the authors neglected the contribution of distant sources of cosmic rays with energy above $4\\times 10^{19}$ eV. When this effect is accounted for, as in a more recent paper \\cite{ensslin} by the same authors, the revised estimate of the source density appears to be in rough agreement with that of \\cite{blasi2003} for the cases of weak magnetization. The authors of \\cite{ensslin} also investigate the cosmic ray propagation in a magnetized universe, as obtained from numerical simulations of large scale structure formation with passive magnetic fields. In these cases the estimate of the source density becomes more model dependent. It is probably worth stressing however, that the magnetic fields obtained in \\cite{ensslin} do not appear to be compatible with those derived in the numerical MHD simulations of \\cite{grasso}: while the magnetic field obtained inside clusters of galaxies in both numerical approaches appears to have roughly the same strength, the fields outside clusters differ by 2-4 orders of magnitude. Not surprisingly the conclusions that the two groups draw on the propagation of UHECRs are incompatible with each other. It would be auspicable that the two pieces of work can achieve some kind of agreement, so to clarify the situation.  The interactions of cosmic ray protons at the highest energies with the photons in the CMBR, responsible for the appearance of the GZK feature in the diffuse spectrum of UHECRs, also generate gamma radiation through the decay of neutral pions and electrons through the decay of charged pions (actually mostly positrons are produced, but hereafter we will refer to both electrons and positrons as {\\it electrons}). The very high energy gamma radiation produced in this way initiates an electromagnetic cascade that results in the appearance of a gamma ray flux in the energy region accessible to Cerenkov detectors. We investigate in detail this process and assess the detectability of this signal from the directions where the sources of UHECRs are located.  The paper is organized as follows: in section \\ref{sec:cascade} we present an analytical estimate of the development of the electromagnetic cascade initiated by UHECRs; in section \\ref{sec:background} we detail the background radiation field we used for our calculations; in section \\ref{sec:detailedcascade} we present the detailed calculation of the electromagnetic cascade developement; in section \\ref{sec:results} we present our results for different luminosities of the sources and different values and topologies of the extragalactic magnetic field; in section \\ref{sec:conclusions} we draw our conclusions. ",
        "conclusions": "\\label{sec:conclusions}  The sources of UHECRs are still unknown mainly because we do not know yet whether they are generated in a few very powerful sources or rather in numerous low power sources. This degeneracy has probably been broken after the discovery of small scale anisotropies in the arrival directions of UHECRs in the AGASA experiment, although the statistical significance of this discovery appears to be still controversial \\cite{finley2003}. In \\cite{blasi2003}, the AGASA data were used to infer that the most likely number density of sources of UHECRs is around $10^{-6}-10^{-5}~\\rm Mpc^{-3}$, which corresponds to a source luminosity of $2\\times 10^{43}-2\\times 10^{42}~\\rm erg~s^{-1}$ at energies above $10^{19}$ eV. These estimates, obtained in the absence of intergalactic magnetic fields, remain valid as long as the magnetic field does not exceed $\\sim 10^{-10}$ G over cosmological distances.  During their propagation from the sources to the Earth UHECRs suffer photopion production if their energy is above threshold for this process, therefore generating charged and neutral pions. While the former decay mainly into electrons and neutrinos, the latter generate gamma rays. Both gamma rays and electrons initiate an electromagnetic cascade that transfers energy from high energies to lower energies, more easily accessible to gamma ray telescopes. We calculated the development of this electromagnetic cascade with and without the presence of an intergalactic magnetic field, and investigated the detectability of the gamma ray signal with ground based telescopes such as MAGIC, VERITAS and HESS, as well as with space-borne telescopes such as GLAST \\footnote{The fluxes derived in the present paper could be higher by up to a factor of $\\sim 2$ due to the contribution from proton pair production, that was neglected here.}.  The cascade radiation appears to be detectable by VERITAS and HESS for a source luminosity around $2\\times 10^{43}~\\rm erg~s^{-1}$ in the form of cosmic rays with energy larger than $10^{19}$ eV, if the magnetic field in the IGM is constant over all the propagation volume and smaller than $10^{-9}$ G (This luminosity corresponds to the lower source density compatible with the small scale anisotropies observed by AGASA). If the magnetic field is patched, stronger fields are allowed close to the Earth without appreciably changing the previous conclusions. If the region with stronger fields is close to the source location, the initial stages of the cascade may be affected if the magnetized region is large enough, so that the energy content of the cascade gets suppressed. A large field close to the Earth does not affect the conclusions in any appreciable way, since the cascade is already fully developed at that point. A magnetic field of 1nG near the source implies slightly lower gamma ray fluxes if it is extended in a region of about 10 Mpc. A stronger field, of say 10 nG, in a 1 Mpc region either around the source or the observer does not change the results in any appreciable way.  The main uncertainty involved in our calculations is the luminosity of the sources in the form of UHECRs. The numbers we adopted are derived in \\cite{blasi2003}, on the basis of the appearance of small scale anisotropies in the AGASA data and in the assumption of sources that accelerate UHECRs continuously. For bursting sources, the luminosity per source may be many orders of magnitude larger than that adopted here, but it would be concentrated within a short time interval. The diffuse and time-spread appearance of the observed UHECRs is achieved, within this class of models, by assuming the presence of a small magnetic field in the intergalactic medium, responsible for sizeable time delay that make the burst look like a continuous signal when observed through charged particles. This is the case for the gamma ray burst models \\cite{vietri95,waxman95,daniel3} of UHECRs.  A similar time delay is introduced in the development of the electromagnetic cascade. Moreover, an additional time delay is introduced due to the presence of electrons and positrons in the cascade itself. These effects cannot be accounted for within the framework proposed here, and will be considered in a forthcoming publication."
    },
    "0404/astro-ph0404508_arXiv.txt": {
        "abstract": "We report the detection of atomic chlorine emissions in the atmosphere of Io using Hubble Space Telescope observations with the Goddard High Resolution Spectrograph (GHRS).  The \\ion{Cl}{1} $\\lambda$1349 dipole allowed and \\ion{Cl}{1}] $\\lambda$1386 forbidden transition multiplets are detected at a signal to noise ratio (SNR) of 6 and 10, respectively, in a combined GHRS spectrum acquired from 1994 through 1996.  Oxygen and sulfur emissions are simultaneously detected with the chlorine which allows for self-consistent abundance ratios of chlorine to these other atmospheric species.  The disk averaged ratios are:  Cl/O $=$ 0.017 $\\pm$ 0.008, Cl/S $=$ 0.10 $\\pm$ 0.05, and S/O $=$ 0.18 $\\pm$ 0.08.  We also derive a geometric albedo of 1.0$\\,\\pm\\,$0.4\\,$\\%$ for Io at 1335\\,\\AA\\ assuming an SO$_2$ atmospheric column density of 1\\,$\\times$10$^{16}$\\,cm$^{-2}$. ",
        "introduction": "Io is the most volcanically active body in the solar system.  First viewed by Voyager \\citep{mor79}, observable volcanic activity revived the search for an atmosphere around Io.  After several years of multi-wavelength studies \\citep{mcg04,lea96,spe96}, it is known that a tenuous atmosphere maintained by sublimation and volcanic outgassing and composed primarily of SO$_2$ \\citep{pea79} and SO \\citep{leb96} exists at Io.  This atmosphere is denser in the equatorial regions of the satellite and above active volcanic plumes \\citep{jes04,fel00,mcg00}.  Minor atmospheric species including sulfur and oxygen \\citep{bal87}, the atomic by-products of SO$_2$, and sodium \\citep{sch87} have been observed at Io.  Recently, emissions from chlorine \\citep{re00a}, and NaCl \\citep{lel03} have also been detected in the atmosphere with inferred abundances lower than those of sulfur and oxygen. Extended clouds of potassium have been reported \\citep{tra75} suggesting potassium local to Io, and \\citet{gei04} identified potassium as the most probable source of IR emission detected in the equatorial glows, but it has not yet been confirmed in the bound atmosphere.  Since Io's orbit around Jupiter takes it into and out of the densest regions of the Io plasma torus, electrons in the torus, co-rotating with Jupiter's magnetic field, continually bombard Io's atmosphere, collisionally exciting and ionizing atmospheric species. The excited species de-excite by emitting photons and produce a rich ultraviolet spectrum.  The ionized material can be captured by the magnetic field and populates the torus.  This and other loss mechanisms account for a supply of $\\sim\\,$10$^{30}$ amu per second from Io to the jovian magnetosphere.  As a result, elements observed in the atmosphere are expected to be components of the torus.  Likewise, species detected in the torus should be found in some form in Io's atmosphere.  \\citet{kup00} reported the discovery of Cl$^+$ in the Io plasma torus in the near infrared.  From their measurements, the abundance of chlorine ions was estimated to be 4\\% relative to sulfur ions, indicating a chlorine to total ion mixing ratio of 2\\%.  For their abundance calculations, they assumed that S$^+$, S$^{2+}$, O$^+$, and Cl$^+$ ions constitute 85\\% of the torus ion density.  In a follow up study by Schneider, Park, \\& K\\\"uppers (2000), Cl$^{2+}$ was detected at comparable levels to Cl$^+$, implying a slightly larger chlorine to total ion mixing ratio for the torus than first estimated. In a more recent study, \\citet{fel01} detected both Cl$^+$ and Cl$^{2+}$ in a far-UV spectrum of the torus and derived a Cl$^{2+}$ abundance of 3\\% relative to S$^{2+}$, implying an abundance ratio of 2\\% for all chlorine ions to all sulfur ions, and a ratio of chlorine to all torus ions of at most 1\\%.  In 2001, \\citet{fel04} acquired better signal to noise ratio spectra and confirmed the presence of chlorine ions in the torus, but with a 30\\% lower chlorine to sulfur abundance than previously inferred.  This reduction in chlorine abundance observed in data acquired one year later is highly suggestive of a variable source of chlorine.  Ionized sodium has also been tentatively detected in the torus with a similar estimated abundance of 1\\% \\citep{hal94}.  The presence of chlorine ions in the torus has motivated the search for chlorine in the atmosphere of Io.  If the atmospheric abundance is similar to the $\\leq\\,$1\\% torus abundance, then an atmospheric chlorine column density of $\\sim\\,$1\\,$\\times$10$^{14}$\\,cm$^{-2}$ is implied assuming a nominal SO$_{2}$ column density of 1\\,$\\times$10$^{16}$\\,cm$^{-2}$ \\citep{fel00,mcg00,str01,lel03}.  This implies that neutral chlorine, like oxygen and sulfur, should be a detectable component of Io's atmosphere. Chlorine is known to be outgassed in terrestrial volcanos and thermochemical models show NaCl as the dominant sodium and chlorine bearing gas in Io's volcanic gases (Fegley \\& Zolotov 2000; Moses, Zolotov, \\& Fegley 2002b). This prompted a search for NaCl, which was successfully detected by \\citet{lel03} on both the orbital leading and trailing hemispheres of Io. They inferred a disk averaged NaCl abundance of 0.3\\% relative to SO$_2$. Some of the observed NaCl emission lines are broader than the SO$_2$ emission lines, implying locally enhanced NaCl or higher NaCl temperatures.  Since NaCl is seen on both hemispheres, may be locally enhanced, and has a short lifetime to photolysis, \\citet{lel03} inferred that volcanic activity is the most likely source of the NaCl.  A variable volcanic source of chlorine could also explain the change in chlorine torus abundance reported by \\citet{fel04}.  As a result of rapid NaCl photolysis in Io's upper atmosphere, atomic chlorine and sodium are more abundant than NaCl \\citep{mo02b}.  In an independent study of data acquired on two dates by the Hubble Space Telescope (HST) Space Telescope Imaging Spectrograph (STIS), \\citet{ret02} detected two chlorine emission lines in the equatorial spots of Io and calculated a chlorine to oxygen abundance of 0.07--1\\%.  Here we present unambiguous detections of two ultraviolet multiplets of atomic chlorine present in three years of archived HST data which provide constraints on the atmospheric chlorine abundance.  Although the data show temporal variations over the three years of acquisition, this data set gives the most complete disk and time averaged abundances for the minor species in Io's atmosphere and aids in understanding the nominal atmospheric composition.  We first describe the observations and data reduction, then present the analysis of the chlorine detection, giving estimates of the relative chlorine abundances at Io, and finally, discuss the implications of our findings. ",
        "conclusions": "Utilizing the definition of optical depth, $\\tau\\,=\\,{\\cal N}\\sigma$, where ${\\cal N}$ is the column density and $\\sigma$ is the absorption cross section, and the determination of the \\ion{Cl}{1} $\\lambda$1349 multiplet being optically thick in our data ($\\tau>\\,$1), a lower bound to the disk averaged chlorine column density of 2.3$\\times$10$^{12}$\\,cm$^{-2}$ is implied.  The chlorine absorption cross section is calculated with the oscillator strength from \\citet{bie94} at a temperature of 1000 K.  In the same manner, the \\ion{S}{1} $\\lambda$1405 multiplet appears to be optically thick in our data giving a lower bound to the disk averaged sulfur column density of 2.4$\\times$10$^{13}$\\,cm$^{-2}$ using an oscillator strength from \\citet{tay98}.  This result is consistent with the disk averaged sulfur column density lower bound estimated by \\citet{fea02} for the \\ion{S}{1} $\\lambda$1814 multiplet in IUE and FOS data.  This independent estimate of the sulfur column density lends credence to the range of sulfur column densities found by \\citet{fea02}.  The S/O results presented here are in agreement with the \\citet{wol01} analysis of HST/STIS data.  They find a disk averaged S/O brightness ratio of 1 : 0.81 when comparing the combined \\ion{S}{1} $\\lambda$1479 multiplets to the \\ion{O}{1}] $\\lambda$1356 doublet.  Applying the 40\\% : 60\\% contribution of the forbidden and allowed multiplets \\citep{fea02}, respectively, to the \\citet{wol01} analysis and utilizing the rate coefficients from Table 3, their brightness ratio translates to an S/O abundance ratio of 0.10 $\\pm$ 0.06.  Combining the S/SO$_2$ ratio of 3--7$\\times$10$^{-3}$ measured at three spatially resolved locations by \\citet{mcg00} with our results, an estimate of the Cl/SO$_2$ ratio can be made, which leads to an approximate Cl/SO$_2$ abundance ratio of 3.0--7.0$\\times$10$^{-4}$ (Table 4).  In addition, combining the O/SO$_2$ ratio of 0.05 inferred with Cassini data \\citep{gei04} with our results gives a Cl/SO$_2$ value of 8.5$\\times$10$^{-4}$.  Recently, \\citet{lel03} detected a comparable abundance of NaCl with a disk averaged NaCl/SO$_2$ ratio of 7$\\times$10$^{-4}$--0.02 if NaCl is more localized than SO$_2$, with a preferred value of 3.5$\\times$10$^{-3}$ (Table 4).  In assessing escape rates of sodium, chlorine, and NaCl from Io, \\citet{mo02b} establish that chlorine and sodium have escape fluxes 10--20 times larger than NaCl.  Our estimated atmospheric chlorine abundance and the NaCl abundance of \\citet{lel03} are both 1--2 orders of magnitude lower than the 1\\% torus chlorine mixing ratio \\citep{fel01} indicating that there exists an intricate relationship between the transport of volcanic gases to the upper atmosphere and escape to the torus.  There may also be a higher concentration of chlorine compared to sulfur and oxygen near volcanic outflows which are at lower altitudes than we probe and would boost the overall atmospheric chlorine abundance measured in our data. In the Pele-type volcanic atmospheres modeled by \\citet{mo02b} which include chlorine, sodium, and potassium, the initial chlorine abundance is assumed from the high chlorine torus abundance of \\citet{sch00}.  They do not check the sensitivity of their models to lower chlorine abundances, so we do not compare our results for chlorine.  We do however compare the S/O abundance ratios.  \\citet{mo02a} determine an equilibrium ratio of 10 for S/O, $\\sim\\,$50 times larger than our data shows.  This suggests that either the initial oxygen abundances in the models are too low or the disk averaged atmosphere we detect at Io is dominated by Prometheus-type (SO$_2$ dominated) outflows rather than Pele-type (sulfur dominated).  \\citet{mo02a} do not check the sensitivity of their models to the initial S/O abundance.  Lastly, our Cl/S ratio adds plausibility to the identification of solid Cl$_2$SO$_2$ or ClSO$_2$ in NIMS/Galileo spectra near the volcanic center of Marduk as suggested by \\citet{sch03}.  Unidentified absorption features are present in the infrared spectra and based on their spectroscopic arguments, Cl$_2$SO$_2$ is the favorite candidate with ClSO$_2$ a potential alternative. \\citet{sch03} estimate that for abundant Cl$_2$SO$_2$ formation, the gaseous [Cl-(Na+K)]/S ratio in a volcanic plume needs to be larger than 0.015, and for ClSO$_2$, the ratio must be even larger.  When there is an excess of halogens as compared to alkalis, Cl/(Na+K) $>$ 1, sulfur chlorides and oxychlorides are formed with the surplus chlorine.  Since volcanic outgassing is assumed to be the only source of chlorine in Io's atmosphere, plumes should have at least a comparable, most likely higher, abundance of chlorine relative to the disk averaged atmosphere.  For comparable Cl and NaCl disk averaged abundances, consistent with our data and that of \\citet{lel03}, an estimated Cl/Na ratio of $\\sim\\,$10 is shown in Figure 3d of \\citet{feg00}, and for our Cl/S ratio of 0.1, a Cl/Na ratio $\\geq\\,$2 can be estimated from Figure 5 of \\citet{feg00}.  These ratios imply that gaseous chlorine is at least twice as high in abundance as sodium.  Therefore, our disk averaged Cl/S ratio is consistent with the requirements of \\citet{sch03} for the presence of Cl$_2$SO$_2$ or ClSO$_2$."
    },
    "0404/astro-ph0404022_arXiv.txt": {
        "abstract": " ",
        "introduction": "This paper is conceived as a pedagogical introduction to the motivations for inflation, to its mathematical settings, and to its implications for the Cosmic Microwave Background (CMB) anisotropies.\\footnote{ For the other aspects of CMB physics, we refer to the other contributions of the special issue of {\\it Comptes Rendus Physique} (2003) {\\bf 4}, 8. The full text is available on line at {\\bf http://www.sciencedirect.com}. Notice also that the present text is a slightly expanded version of that published in the above review.} We shall put a special emphasis on justifying the surprising result that quantum mechanical rules must be used to compute the primordial spectrum.  For the non-specialist, two remarks should be made from the outset. First, the simplest inflationary model has successfully passed the multiple checks based on recent observational data. Secondly, inflation is a rather conservative hypothesis since it rests, on one hand, on the set of cosmological observations conventionally interpreted, and on the other hand, on Einstein's equations, i.e., on the hypothesis that the action of General Relativity (GR), or a slight generalization thereof, governs the evolution of the space-time properties from cosmological scales, of the order of a thousand megaparsecs ($1$Mpc $ \\simeq 10^{22}$ meters), down to some microscopic scale which is not far from the Planck length ($l_P \\simeq 10^{-35}$ meters), before the threshold of Quantum Gravity.  Following these preliminary remarks, one should explain why the search for a causal explanation of the large scale properties of the universe calls for a long period of accelerated expansion (where long is defined with respect to the expansion rate during inflation). This need stems from the fact that the `standard model' of cosmology\\cite{Weinberg,Peebles}, the Hot Big Bang scenario, is {\\it incomplete}: the causal structure from the big bang is such that no processes could have taken place to explain the homogeneity and the isotropy at large scales. (What ``large scales'' means depends on the quantity under consideration. For the density fluctuations on the last scattering surface, at the moment of decoupling, large scales concern distances larger than hundred kiloparsecs. These scales are seen today as solid angles larger than 1 degree). The importance of these considerations can only be appreciated in the light of today's understanding of cosmology. ",
        "conclusions": ""
    },
    "0404/astro-ph0404214_arXiv.txt": {
        "abstract": "When a white dwarf (WD) is weakly magnetized and its accretion disk is thin, accreted material first reaches the WD's surface at its equator. This matter slows its orbit as it comes into co-rotation with the WD, dissipating kinetic energy into thermal energy and creating a hot band of freshly accreted material around the equator. Radiating in the extreme ultraviolet and soft X-rays, this material moves toward the pole as new material piles behind it, eventually becoming part of the WD once it has a temperature and rotational velocity comparable with the surface. We present a set of solutions which describe the properties of this ``spreading layer'' in the steady state limit based on the conservation equations derived by Inogamov \\& Sunyaev (1999) for accreting neutron stars. Our analysis and subsequent solutions show that the case of WDs is qualitatively different. We investigate example solutions of the spreading layer for a WD of mass $M=0.6M_\\odot$ and radius $R=9\\times10^{8}\\textrm{ cm}$. These solutions show that the spreading layer typically extends to an angle of $\\theta_{\\rm SL}\\approx0.01-0.1$ (with respect to the equator), depending on accretion rate and the magnitude of the viscosity. At low accretion rates, $\\dot{M}\\lesssim10^{18}\\textrm{ g s}^{-1}$, the amount of spreading is negligible and most of the dissipated energy is radiated back into the accretion disk. When the accretion rate is high, such as in dwarf novae, the material may spread to latitudes high enough to be visible above the accretion disk. The effective temperature of the spreading layer is $\\sim(2-5)\\times10^5\\textrm{ K}$ with approximately $T_{\\rm eff}\\propto\\dot{M}^{9/80}$. This power-law dependence on $\\dot{M}$ is weaker than for a fixed radiating area and may help explain extreme ultraviolet observations during dwarf novae. We speculate about other high accretion rate systems ($\\dot{M}\\gtrsim10^{18}\\textrm{ g s}^{-1}$) which may show evidence for a spreading layer, including symbiotic binaries and supersoft sources. ",
        "introduction": "In close binaries containing a white dwarf (WD), the companion star can fill its Roche lobe either through stellar evolution or angular momentum loss and begin accreting onto the WD to form a cataclysmic variable (CV; Warner 1995). Before material can reach the WD it must transport angular momentum outward via an accretion disk. Observations and theoretical studies show that when the disk can radiate the internally dissipated energy (Pringle 1981) it will be thin (i.e. the vertical height of the disk is much smaller than the WD radius; Shakura \\& Sunyaev 1973; Meyer \\& Meyer-Hofmeister 1982) and rotating at nearly Keplerian velocity. If the WD is weakly magnetized, the accreted material reaches the surface at the equator and then must pass through a transition region to settle into the WD and become part of the star. When the WD spin is much less breakup, nearly half of the accretion luminosity is released in this region, making it as luminous as the accretion disk and likely crucial to understanding the luminosity from accreting WDs.  The transition from the accretion disk to the WD surface is typically treated in a boundary layer (BL) model, which is simply an extension of the accretion disk with additional torques provided by the WD to decelerate disk material to the WD rotational velocity. In BL studies the important coordinate is assumed to be radial and the vertical structure is solved by assuming a vertical pressure scale height (e.g. Popham \\& Narayan 1995). The geometry of the BL is of a fattened disk near the WD surface. Other BL studies solve the two-dimensional problem simultaneously in both radial and latitudinal directions using numerical techniques (e.g. Kley \\& Hensler 1987; Kley 1989a, 1989b). Kippenhanh \\& Thomas (1978) investigated the properties of ``accretion belts'' on WDs by assuming that the accreted material is marginally stable to the Richardson criterion. We now want to understand the state of recently accreted material on much shorter timescales ($\\ll10^{4}\\textrm{ yr}$), and assess the effects of a turbulent viscous stress on the spreading properties. In Inogamov \\& Sunyaev (1999; hereafter IS99), this problem was approached from a new angle to study accreting neutron stars (NSs). Their method follows the latitudinal flow of matter on the compact object which provides information about the spreading area of hot, freshly accreted material which is not captured in BL models. We now apply these same methods to the case of WDs and call this model the spreading layer (SL), to differentiate it from BL studies. We solve the conservation equations derived by IS99 for fluid confined to the WD surface and find that the steady state limit naturally attracts toward solutions where local viscous dissipation is nearly balanced by local radiative cooling. This is qualitatively different than the results of IS99 who found that the SLs of NSs show a great deal of advection of dissipated energy up to higher latitudes where it is radiated. This is mainly due to their high accretion rates near the Eddington limit.  For a fiducial WD of mass $M=0.6M_\\odot$ and radius $R=9\\times10^{8}\\textrm{ cm}$, our solutions yield a thin hot band of extent $\\theta_{\\rm SL}\\approx0.01-0.1$ with an effective temperature of $\\sim(2-5)\\times10^5\\textrm{ K}$ over a range of accretion rates from $10^{17}-10^{19}\\textrm{ g s}^{-1}$, implying that the SL should contribute to accreting WD spectra in the extreme ultraviolet (EUV) or soft X-rays. We also find the SL's dependence on accretion rate and viscosity. These solutions show that at low accretion rates ($\\dot{M}\\lesssim10^{18}\\textrm{ g s}^{-1}$) the SL does not appear above the accretion disk, so that radiation from the SL may be obscured by the accretion disk or its winds. Only when $\\dot{M}\\gtrsim10^{18}\\textrm{ g s}^{-1}$, such as in dwarf novae, symbiotic binaries, or supersoft sources, will the SL reach high enough latitudes to be directly seen. Perhaps the best chance for testing the SL model are dwarf novae in outburst and superoutburst. Some observations show a fairly constant effective temperature over a large change in luminosity (for example SS Cyg; Mauche, Raymond \\& Mattei 1995) which is consistent with a SL. Direct comparisons between our simple model and observations can be complicated if the SL emission is absorbed and re-radiated by accretion disk winds, as is likely for systems which show that the EUV is not eclipsed by the secondary whereas the WD is (such as for the edge-on OY Car in superoutburst; Mauche \\& Raymond 2000). Nevertheless, the SL may be important for calculating the underlying continuum scattered by these winds. We predict a quick timescale for spreading, $t_{\\rm SL}\\sim50\\textrm{ s}$, which means that the typical mass of the spreading region ($\\lesssim10^{20}\\textrm{ g}$) is much less than the ``accretion belt'' of Kippenhahn \\& Thomas (1978). Most likely, as angular momentum is tranferred between the SL and the underlying WD surface, the underlying surface layers will spin up similar to what Kippenhahn \\& Thomas (1978) calculate. The short spreading timescale also means that following a dwarf nova outburst, when accretion onto the WD has decreased dramatically, the SL will spread over the WD surface too quickly to be observed as a separate hot component.  In \\S 2 we solve for the radial structure of the SL using a one-zone model. We derive the differential equations which describe the SL in the steady state limit in \\S 3, using IS99 as a guide. In \\S 4 we discuss the solutions of the spreading equations. We study physically motivated analytic estimates which capture the essential features of the numerical integrations. These provide power law relations which show how the SL changes with accretion rate, viscosity, and the properties of the accreting WD. We then present the results of integrating the equations numerically to undestand the SL in more detail. In \\S 5 we discuss observational tests for the SL and consider the current data in relation to our calculations. We conclude in \\S 6 with a summary of our findings along with a discussion about further work which can be done with this model. ",
        "conclusions": "The SL model, first conceived by IS99, provides a new way of understanding the properties of newly accreted material and its interaction with the stellar surface. In this model the latitudinal direction is used as the independent variable, as opposed to standard BL models which follow how the disk changes near the star along the radial direction. This allows an investigation of the radiating area of the hot belt which forms near the equator, and it describes the properties of the accreting material when it first comes into hydrostatic balance on the star. We have investigated the solutions of the SL model when applied to WDs, accreting in the range of $\\dot{M}\\sim10^{17}-10^{19}\\textrm{ g s}^{-1}$. We find that the integrations for the spreading flow naturally attract toward solutions in which initially local radiative cooling balances local viscous dissipation. As the accretion rate gets higher, we find increasingly more advection in the flow, similar to the solutions for NSs investigated by IS99. The spreading angle on WDs is found to typically be $\\theta_{\\rm SL}\\approx0.01-0.1$, depending on both the accretion rate and the viscosity. The SL has an effective temperatures in the range of $(2-5)\\times10^5\\textrm{ K}$ with a pressure scale height of $10^7-10^8\\textrm{ cm}$.  To be clearly observed in actively accreting systems, the SL must extend to a large enough angle to be seen above the accretion disk. This requires the accretion rate to be high, $\\dot{M}\\gtrsim10^{18}\\textrm{ g s}^{-1}$, assuming that the accretion disk is similar to a thin Shakura-Sunyaev disk. The best candidates to show the SL are dwarf novae in outburst, supersoft sources, and symbiotic binaries. Current observations of the dwarf novae may be the best opportunity to see if spreading is actually occurring. The scaling of effective temperature with accretion rate is much weaker in the case of a SL because of the change in radiating area. At lower accretion rates, $\\dot{M}\\lesssim10^{18}\\textrm{ g s}^{-1}$, the SL will not spread far from the equator and most of the dissipated energy will radiate back into the accretion disk. Although this may make comparison between theory and observations difficult, because the SL cannot be seen directly, it may still have important consequences for the accretion disk structure and spectra.  The short timescale for spreading, $t_{\\rm SL}\\sim50\\textrm{ s}$, suggests that a hot belt around the WD equator will be difficult to see after an outburst because the hot material should spread over the WD surface quickly. This also means that mixing between freshly accreted material and much deeper layers appears difficult on accreting WDs. Rosner et al. (2001) and Alexakis et al. (2003) have argued that the mixing of recently accreted H/He with the underlying layer of C/O due to rotational shearing explains observations of CNO nuclei in the ejecta of nova. To have $t_{\\rm SL}$ of order the accretion time ($\\sim10^3\\textrm{ yr}$) requires $\\alpha\\lesssim10^{-11}$, so that the viscosity must be even less than what is estimated from microphysics (see Appendix A). This makes it difficult to imagine mixing via this mechanism all the way down to the C/O layers.  The initial calculations by IS99 and in this paper describe some of the general features of the SL model, a new area of investigation which will lead to further studies of how accreted material settles onto stars. This may include studying how differential rotation and angular momentum transfer affect the underlying stellar surface or the possibility of nonradial oscillations present in the SL. Spectral modeling of the SL, along comparisons with observations, may also help in identifying if and when spreading is present.  We thank Rashid Sunyaev for helpful and enthusiastic discussions about spreading, and Christopher Mauche for answering our questions about cataclysmic variable observations. We have also benifitted from conversations with Phil Arras, Philip Chang, Chris Deloye, Aristotle Socrates, and Dean Townsley. We thank the referee for a careful reading of this paper along with many helpful suggestions. 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}"
    },
    "0404/astro-ph0404091_arXiv.txt": {
        "abstract": "We apply an eclipse mapping technique using `genetically modified fire-flies' to the eclipse light curves of HU~Aqr and EP~Dra. The technique makes as few assumptions as possible about the location of accretion stream material, allowing the emission to be located anywhere within the Roche lobe of the white dwarf. We model two consecutive eclipses in the $UBVR_c$-band for HU~Aqr, and four consecutive `white'-light eclipses for EP~Dra, to determine the changing brightness distribution of stream material. We find fire-fly distributions which are consistent with accretion through a curtain of material in both HU Aqr and EP Dra, and show that the previously assumed two part ballistic and magnetic trajectory is a good approximation for polars. Model fits to the colour band data of HU~Aqr indicate that the material confined to the magnetic field lines is brightest, and most of the emission originates from close to the white dwarf. There is evidence for emission from close to a calculated ballistic stream in both HU~Aqr and EP~Dra.We propose that a change in the stream density causes a change in the location of the bright material in the accretion stream in EP~Dra. ",
        "introduction": "AM Hers, or polars, are a sub-class of the magnetic cataclysmic variable (CV) interacting binaries. A main sequence secondary transfers material to a white dwarf primary through an accretion stream, from mass overflow at the inner Lagrangian point ($L_1$). The primary has a magnetic field of order 10$\\--$200~MG, which prevents the formation of an accretion disk (as found in other types of cataclysmic variables), and instead, at some threading radius from the primary, the material is confined to follow the magnetic field lines to accrete directly onto the white dwarf surface (see Cropper 1990 for a review). A number of these systems have inclinations to the line of sight such that the system is eclipsed by the secondary, and this can be used as a means of isolating the emission from discrete parts of the system. In particular the brightness distribution of material along the accretion stream can be inferred as successive sections of the accretion stream material are eclipsed.  Attempts to determine the brightness distribution of the accretion stream material has evolved in complexity from initially one-dimensional streams confined to the orbital plane (Hakala 1995), to three dimensional tubes carrying material far out of the orbital plane (Kube, G\\\"ansicke \\& Beuermann 2000; Vrielmann \\& Schwope 2001). These methods have been applied to both emission line (Kube, G\\\"ansicke \\& Beuermann 2000; Vrielmann \\& Schwope 2001) and continuum observations (Hakala 1995; Harrop-Allin et al. 1999b, 2001; Bridge et al. 2002). The assumption common to all previous model techniques is that of a stream trajectory determined prior to the modelling process, and fixed for the duration of the model. In an attempt to remove as many of the assumptions about the stream location as possible, a technique has been developed by Hakala, Cropper \\& Ramsay (2002), which makes fewer assumptions about the location of bright stream material. In principle, stream material can be located anywhere within the Roche lobe of the primary.  The application of the technique to synthetic data sets was demonstrated in Hakala et al. (2002), and here we apply the method to optical light curves of two eclipsing polars: EP~Dra and HU~Aqr. The two systems show variations in the brightness and trajectory of the accretion stream over the timescale of the orbital period (Bridge et al. 2002, 2003), with EP~Dra also showing a variation in brightness over a longer phase range. This variation appears to be related to the brightness of the accretion stream, and is attributed to a combination of cyclotron beaming and absorption in an extended accretion curtain (Bridge et al. 2002, 2003).  The eclipse mapping technique based on the model of Harrop-Allin et al. (1999a) was found (by Bridge et al. 2002) to be particularly sensitive to the input parameters and the signal-to-noise ratio of the data used, when applied to the selected light curves of HU~Aqr, and hence restricted the interpretation of the results. We therefore apply this new model technique to the EP~Dra and HU~Aqr light curves in an attempt to circumvent these limitations. ",
        "conclusions": "We have applied the technique of eclipse mapping using genetically modified fire-flies to the eclipse lights curves of EP~Dra and HU~Aqr. The modelling shows that the technique is applicable to relatively good signal-to-noise ratio light curves of adequate phase range. These distributions of fire-flies show that the previously assumed ballistic free-fall plus magnetically confined streams are a good approximation for the accretion streams in polars.  We applied the model to $UBVR_c$-band light curves of HU~Aqr cycles 29994 and 29995, and `white'-light curves of EP~Dra cycle 56962, 56976, 56977 and 56978. We have demonstrated that the technique will distinguish regions of brightness in the different colour bands. This may indicate different temperatures or densities, and hence heating and/or cooling processes in the accretion stream or the location of photoionised material.  The fire-flies distributions in both HU~Aqr and EP~Dra show the accretion stream to be brightest near to the white dwarf. A possible threading region is present at which the fire-fly distributions appear to deviate from a calculated purely ballistic trajectory, and this region is seen to be broad, implying accretion along many field lines.  Differences in the location and concentration of fire-flies between the three cycles, 56976, 56977 and 56978, of EP~Dra indicate a changing brightness distribution between the cycles. This could be the result of a change in the temperature of the accretion stream material, or a change in the amount of material stripped from the ballistic trajectory and coupled to the magnetic field lines of the primary."
    },
    "0404/astro-ph0404572_arXiv.txt": {
        "abstract": "{ New empirical calibrations of the Red Giant Branch Tip in the I,J,H and K bands based on two fundamental pillars, namely $\\omega$~Centauri and 47 Tucanae, have been obtained by using a large optical and near infrared photometric database.  Our best estimates give $\\mi = -4.05 \\pm 0.12$, $\\mj = -5.20 \\pm 0.16$, $\\mh = -5.94 \\pm 0.18$ and $\\mk = -6.04 \\pm 0.16$ at $[M/H]\\simeq-1.5$ ($\\omega$~Cen) and $\\mi = -3.91 \\pm 0.13$, $\\mj = -5.47 \\pm 0.25$, $\\mh = -6.35 \\pm 0.30$ and $\\mk = -6.56 \\pm 0.20$ at $[M/H]\\simeq-0.6$ (47~Tuc).  With these new empirical calibrations we also provide robust relations of the TRGB magnitude in I, J, H and K bands as a function of the global metallicity. It has also been shown that our calibrations self-consistently provide a distance modulus of the Large Magellanic Cloud in good agreement with the standard value ($(m-M)_0\\simeq 18.50$)). ",
        "introduction": "The use of Tip of the Red Giant Branch (TRGB) as a standard candle is now a mature and widely used technique to estimate the distance to galaxies of any morphological type \\cite[see][for a detailed description of the method, recent reviews and applications]{lfm93,mf95,mf98,walk}. The underlying physics is well understood \\citep{mf98,scw} and the observational procedure is operationally well defined \\citep{mf95}. The key observable is the sharp cut-off occurring at the bright end of the RGB Luminosity Function that can be easily detected with the application of an edge-detector filter \\citep[Sobel filter,][]{mf95,smf96} or by fitting the LF with a proper modeling function \\citep{mendez}. The necessary condition for a safe application of the technique is that the observed RGB Luminosity Function should be well populated, with more than $\\sim 100$ stars within 1 mag from the TRGB \\citep{mf95,draco}.  The RGB develops to its full extension in stellar populations having age \\gtsima 1-2 Gyr, therefore the natural local systems with which the method may be calibrated are Galactic globular clusters \\citep{da90,f00}. However, any calibration based on globulars must rely on the RR Lyrae (or Horizontal Branch) distance scale, whose zero-point is still quite uncertain \\citep{carla,walk}. Moreover the estimate of the TRGB in globular clusters is plagued by the low number of RGB stars available in each cluster, implying a considerable statistical uncertainty and forcing the adoption of a different operational definition of the observable \\cite[e.g., the brightest non-variable RGB star instead of the LF cut-off, see][and reference therein, for details and discussion]{da90,f00,tip}.  To overcome these drawbacks \\citet[][hereafter Pap-I]{tip} used a large photometric database of the globular cluster $\\omega$~Centauri (the most luminous in the whole Galaxy, therefore the one having the largest number of RGB stars) and the distance estimate to this cluster by \\citet[][based on the double-lined detached eclipsing binary OGLE 17]{ogle}, to provide an accurate zero point to the TRGB calibration in the Cousins' I passband, independent of the RR Lyrae distance scale. The detection of the cut-off in the LF of the upper RGB was clean and supported by a large sample comprising more that 180 stars in the brightest 1 mag of the RGB, e.g. fully satisfying all the requirements for a safe estimate of the TRGB magnitude as defined by \\citet{mf95} and \\citet{draco} by means of numerical experiments. Furthermore the use of an essentially geometrical distance estimate to the considered cluster to calibrate the TRGB makes it, in practice, a {\\em primary} standard candle. The calibration has been already applied by several authors in different contexts \\citep{tosi,apella,draco,alves,alves2}.  In the present paper we shortly re-discuss the analysis of Pap-I at the light of some new results appeared in the literature and we extend the calibration to the Near Infrared (NIR) passbands with a statistically robust estimate of the absolute magnitude of the TRGB of $\\omega$~Cen in the J,H and K bands based on the large NIR photometric database assembled by \\citet{antonio}, complemented with 2MASS \\citep{cutri} data. While the metallicity dependence of the magnitude of the TRGB is minimized in the Cousins' I band, the most suitable for distance estimates, the calibration of the standard candle in the NIR may prove very useful since (a) it may allow a safe application also in the cases of very high interstellar extinction, and (b) it will provide a natural tool for future studies performed with powerful telescopes optimized for infrared observations, as, for example, the James Webb Telescope that will push the application limit of the TRGB technique to much larger distances. Nevertheless, it has to be recalled that a safe application of the NIR calibrations requires a good estimate of the metallicity of the considered stellar system, because of the strong dependence of the TRGB brightness in these passbands (K and H, in particular).  To check our calibration also at different metallicity regimes (with respect to $\\omega$~Cen) we obtained B,V,I photometry of more that 100000 stars in a $\\sim 0.5\\times 0.5$ deg$^2$ field centered on the metal rich globular cluster 47~Tucanae. With this new large sample, complemented with NIR photometry from \\citet{paolo} and 2MASS, we estimate the TRGB magnitude in I,J,H and K bands also in this cluster.  The plan of the paper is the following: in Sect.~2 we shortly re-discuss the calibration of Pap-I and we consider the possibility of extension of the calibration to other clusters (i.e. metallicity); in Sect.~3 and 4 we present our new estimates of the absolute magnitude of the TRGB in the various considered passbands for $\\omega$~Cen and 47~Tuc, respectively; in Sect.~5 we present the new calibrations of the TRGB in I, J, H and K bands as a function of the global metallicity ([M/H], see \\citet{scs93}) and we compare our derived distance scale with the ``standard'' distance modulus of the Large Magellanic Cloud \\cite[$(m-M)_0=18.50$, see][]{alves2}. Finally, the main results of the paper are summarized in Sect.~6. ",
        "conclusions": "We have refined the zero-point of the calibration of the TRGB as a standard candle provided by \\citet[][Pap-I]{tip} adopting a more accurate estimate of the reddening to $\\omega$ Centauri, the fundamental pillar of our distance scale. The calibration of Pap-I has been accordingly revised and extended to NIR (J,H,K) passbands using a new large photometric dataset of $\\omega$ Cen giants taken from S03 and 2MASS.  A large optical and NIR photometric database has been assembled for the metal rich globular cluster 47~Tucanae. From this database we obtained the best estimate of the TRGB presently available for this cluster in I, J, H and K, providing a second observational point for our TRGB calibrations, at $[M/H]\\simeq -0.6$.  With this new observational material we provide new robust calibrations of the TRGB magnitude in I, J, H and K bands as a function of the global metallicity. The obtained NIR calibrations are in excellent agreement with the theoretical predictions by \\citet{sc98}, while the marginal discrepancy in the I band already noted in Pap-I persists and it is probably due to the color transformation applied to theoretical models. It has also been shown that our calibrations self-consistently provide a distance modulus of the Large Magellanic Cloud in good agreement with the ``standard value'' \\cite[see][]{alves2,walk,gisella}.  With the present contribution, we think we have provided the ``state of the art'' instrument of stellar astrophysics needed to afford the study of extragalactic distances with the Population II distance scale \\citep{walk,alves2}, e.g. a fundamental complement and consistency test for Cepheid and SNIa distance scales."
    },
    "0404/astro-ph0404058_arXiv.txt": {
        "abstract": "We present results of a binary population synthesis study on the orbital period distribution of wide binary millisecond pulsars forming through four evolutionary channels. In three of the channels, the progenitor of the millisecond pulsar undergoes a common envelope phase prior to the supernova explosion which gives birth to the neutron star. In the fourth channel, the primary avoids the common-envelope phase and forms a neutron star when it ascends the asymptotic giant branch. The four formation channels yield an orbital period distribution which typically shows a short-period peak below 10 days, a long-period peak around 100 days, and a cut-off near 200 days. The agreement with the orbital period distribution of observed binary millisecond pulsars in the Galactic disk is best when the common-envelope ejection is efficient, the mass-transfer phase responsible for spinning up the pulsar is highly non-conservative, and no or moderate supernova kicks are imparted to neutron stars at birth. ",
        "introduction": "In close binaries containing a neutron star (NS) and a non-compact star, the evolution of the latter and/or the orbit may drive the system into a semi-detached state where the non-compact star transfers mass to the NS. If the NS is able to accrete some of the transferred mass, the associated transfer of angular momentum will spin the NS up until a millisecond pulsar is born. At the end of the mass-transfer phase, the donor star's core is exposed as a white dwarf (WD) and the binary appears as a binary millisecond pulsar (BMSP). If the mass-transfer phase takes place when the donor is on the giant branch, the relation between the core mass and radius of the donor star, and the relation between the Roche-lobe radius and the orbital separation, lead to a correlation between the orbital period of the BMSP and the WD mass (e.g. Joss, Rappaport \\& Lewis 1987).  In this paper, we compare the orbital period distribution of simulated wide BMSPs from the BiSEPS ({\\it Bi}nary {\\it S}tellar {\\it E}volution and {\\it P}opulation {\\it S}ynthesis) code with the orbital period distribution of BMSPs observed in the Galactic disk. We particularly investigate whether a set of binary evolution parameters can be found which is able to reproduce the observed distribution without including observational selection effects or detailed pulsar physics. ",
        "conclusions": ""
    },
    "0404/astro-ph0404602_arXiv.txt": {
        "abstract": "We present an 82 ksec {\\sl Chandra} ACIS-I observation of a large-scale hierarchical complex, which consists of various clusters/groups of galaxies and low-surface brightness X-ray emission at $z = 0.247$. This high-resolution {\\sl Chandra} observation allows us for the first time to separate unambiguously the X-ray contributions from discrete sources and large-scale diffuse hot gas. We detect 99 X-ray sources in a $17^\\prime \\times 17^\\prime$ field. Ten of these sources are identified as members of the complex and are mostly radio-bright. Whereas unresolved X-ray sources tend to be associated with galaxies in intermediate density environments, extended X-ray emission peak at bright radio galaxies in the central cluster. In particular, a distinct X-ray trail appears on one side of the fast-moving galaxy C153, clearly due to ram-pressure stripping. The diffuse X-ray emission from the central cluster can be characterized by a thermal plasma with a characteristic temperature of $3.2_{-0.4}^{+0.5}$ keV and a heavy element abundance of $0.24_{-0.12}^{+0.15}$ solar (90\\% confidence uncertainties). In comparison, a patch of low-surface brightness X-ray emission apparently originates in relatively low density intergalactic gas with a characteristic temperature of $0.98_{-0.27}^{+0.22}$ keV and an abundance of $\\lesssim 0.09$ solar. The {\\sl Chandra} observation, together with extensive multi-wavelength data, indicates that the complex represents a projection of several galaxy sub-structures, which may be undergoing major mergers. We discuss the dynamic states of the complex and its sub-structures as well as properties of X-ray-emitting galaxies and the relationship to their environments. ",
        "introduction": "The structure of the universe is believed to have formed by clustering hierarchically from small to large scales. The outcome of this hierarchical formation process depends largely on the interplay between galaxies and their environments. But how and where such galaxy-environment interactions primarily occur remain greatly uncertain (e.g., David et al. 1996; Wang \\& Ulmer 1997; Ponman et al. 1999; Fujita 2001; Balogh et al. 2002; Bekki et al. 2002).  We have identified a large-scale hierarchical complex (Fig.\\ 1) that is well-suited for investigating the structure formation process and the environmental impact on galaxy properties. Revealed in a survey of 10 Butcher \\& Oemler clusters observed with the \\rosat\\ PSPC (Wang \\& Ulmer 1997), this complex contains various X-ray-emitting features, which are associated with concentrations of optical and radio galaxies (Fig.\\ 1; Wang, Connolly, \\& Brunner 1997a; Owen et al. 1999; Dwarakanath \\& Owen 1999). The overall optical galaxy concentration of the region was originally classified as a cluster Abell 2125 (richness 4). The \\rosat\\ image and follow-up optical observations, however,  have shown that the complex contains three well-defined X-ray bright clusters (Wang et al. 1997a). In addition, substantial amounts of unresolved low-surface brightness X-ray emission (LSBXE) are also present. The overall angular size of the entire X-ray-emitting complex seems to extend more than $\\sim 30^\\prime$. But the three main concentrations of galaxies (LSBXE, the central A2125 cluster, and Cluster B) are identified within a smaller projected region of dimension $\\sim 12^\\prime$ ($1^\\prime = 0.23$ Mpc; the cosmological parameters, $H_0 = 71 {\\rm~km~s^{-1}~Mpc^{-1}}$, $\\Omega_{total} = 1$, and $\\Omega_\\Lambda = 0.73$ are adopted throughout the paper). The complex thus represents an X-ray-bright hierarchical filamentary superstructure, as predicted by numerical simulations of the structure formation (e.g.,  Cen \\&  Ostriker 1996).   \\begin{figure}[htb!] \\caption{\\protect\\footnotesize \\rosat\\  PSPC X-ray image of the Abell 2125 complex and its vicinity in the 0.5-2~keV band (Wang et al. 1997a). Point-like X-ray sources detected in the image have been excised. Each contour is 50\\% (2$\\sigma$) above its lower level. Spectroscopically confirmed, radio detected members of the complex (Owen et al. 2004a) are marked by {\\sl pluses}. The box outlines the field covered by our \\chandra\\ ACIS-I observation (Fig. 2). The large cross and circle show the center position and radius used in the Abell catalog.} \\label{fig1} \\end{figure}   \\begin{figure*} \\unitlength1.0cm \\centerline{\\vbox{ \\vspace{-2cm} \\vspace{-2.2cm} }} \\caption{\\protect\\footnotesize \\chandra\\ ACIS-I images of the Abell 2125 field in the 0.5-2 keV (upper panel) and 2-8 keV (lower panel) bands. The images are smoothed with a Gaussian of FWHM equal to 3\\as. The small circles represent the regions removed for discrete sources, which are labeled as in Table 1. The two large circles represent the cluster and the LSBXE, while the large ellipse outlines the diffuse emission to the southwest. The two squares outline the off-source background regions, from which spectra are extracted. } \\label{fig2} \\end{figure*}  The nomenclature adopted here  (e.g., in Fig. 1) follows that used in Wang et al. (1997), based on X-ray identifications.   We call the entire elongated diffuse X-ray enhancement as seen in Fig. 1 as the Abell 2125 complex. Historically, however, this congregation of optical galaxies was loosely called a cluster. The centroid listed in the Abell catalog (Abell 1958) is R.A., Dec. (J2000) = $15^h40^m55^s$, 66$^\\circ 19^\\prime15^{\\prime\\prime}$, as marked in Fig. 1.  The Abell radius, 9\\farcm2 (2 Mpc), encloses the bulk of the galaxy concentrations in the field of Fig. 1, including Cluster B and much of the LSBXE. The work by Butcher et al. (1983) and Butcher \\& Oemler (1984), however, assumed a cluster center that is close to what is inferred from the X-ray for the central Abell 2125 cluster (Fig. 1). Furthermore, their imaging field of view (FoV $= $ 55 arcmin$^2$) is comparable to the size of the X-ray cluster, but is much smaller than the area enclosed in the Abell radius. The central cluster itself, which is only part of the complex, is not particularly rich. This can, at least partly, explain its relatively large blue galaxy fraction ($f_b \\sim 20\\%$; Butcher \\& Oemler 1984). But, the Abell 2125 complex does seem to contain a distinct population of active galaxies (Owen et al. 2004a). There is an over-abundance of radio galaxies in Abell 2125, $\\sim 9\\%$ compared to a typical 2\\% for rich clusters at similar redshifts (Morrison \\& Owen 2003).  We have obtained a deep \\chandra\\ observation that covers part of the Abell 2125 complex (Fig. 1) to characterize its detailed X-ray properties and to study the interplay between galaxies with their environments. The high spatial resolution of \\chandra\\ enables us to examine diffuse X-ray structures down to a scale of $\\sim 3.8$~kpc ($\\sim 1^{\\prime\\prime}$). In this paper, we concentrate on presenting the observation and the results on the detection of discrete sources and on the characterization of large-scale diffuse X-ray emission. Detailed analysis of the X-ray data in conjunction with observations in other wavelength bands will be discussed elsewhere (e.g., Owen et al. 2004a,b), although a few galaxy counterparts will be mentioned in the present work. X-ray sources are labeled with a two-digit number (with a prefix ``X-'' in the text), whereas optical IDs are in five-digits (e.g., 00047). ",
        "conclusions": "The above results provide a detailed characterization of the diffuse X-ray emission and discrete sources in the Abell 2125 complex. In the following, we try to address several key issues about the complex by incorporating the results with information learned from observations in other wavelength bands.  \\subsection{Abell 2125 as a Large-scale Hierarchical Complex}  This distinct complex of galaxies and hot gas shows several remarkable characteristics:  \\begin{itemize} \\item a large velocity dispersion which appears to be due to two or three major components  with a total projected extent of several Mpc (Miller et al 2004);  \\item an exceptionally large fraction of radio galaxies, which are located primarily outside the central Abell 2125 cluster (Owen et al. 2004a);  \\item the presence of multiple X-ray-emitting clusters and LSBXE features, each with substantially lower luminosity and temperature than expected from the overall galaxy richness and large velocity dispersion of the complex (see also Wang et al. 1997). \\end{itemize}  Comparisons with various numerical simulations indicate that the complex represents a projection of multiple components, which might be in a process of merging with each other (Miller et al. 2004; Owen et al. 2004a). Projection effects, together with enhanced activities during this process, may explain the observed characteristics.  \\subsection{Nature of the LSBXE}  Globally speaking, the prominent southwest patch of the LSBXE characterized here probably represents a local temperature and density enhancement of a large-scale diffuse intergalactic medium (IGM) structure, as indicated first in the \\rosat\\ PSPC image (Fig. 1), which is more sensitive to very soft X-rays ($\\lesssim 0.5$ keV) than the \\chandra\\ observation.  The LSBXE is distinctly different from the central cluster. Even the relatively X-ray-bright southwest patch does not seem to be centrally peaked, as may be expected from a more-or-less virialized intracluster medium (ICM). The kinematics (Miller et al. 2004) indicates that the galaxies in this patch are probably loosely bound, although their exact line-of-sight distribution is not clear. Figs. 17c,d  show a sub-field of the LSBXE, as sampled by an {\\sl HST} WFPC-2 image. X-14 and X-22 may be associated with two bright ellipticals in the field. But the centroid of X-14 is closer to a faint point-like optical object about 4$^{\\prime\\prime}$ offset from the corresponding elliptical. Therefore, the X-ray source may represent a background AGN. Most of other galaxies in the field appear to be spirals. X-ray emission from such galaxies are typically dominated by X-ray binaries and/or AGNs and is therefore expected to have a relatively hard spectrum. But the diffuse X-ray enhancement associated with the galaxy concentration is soft, indicating an origin in diffuse hot gas around individual galaxies, in groups of galaxies, and/or in the intergroup IGM.  The heavy element abundance in the LSBXE appears to be substantially lower than that in the ICM of the central cluster. The abundance may be underestimated in the LSBXE, if it contains multiple temperature components. Although the single temperature thermal plasma gives a satisfactory fit, we also tried a model with two temperature components. This model gives a marginally improved fit to the LSBXE spectrum ($\\chi^2/d.o.f = 18.2/23$, compared with that that in Table 2), but provides no useful constraints on the abundance, because of the limited counting statistics of the data. If the low metallicity obtained from the single temperature plasma fit is real, there may be two possible explanations: 1) Metals (mainly irons) are largely locked in dust grains within or around galaxy groups (e.g., Nollenberg, Williams, \\& Maddox, 2003); 2) The metals are still bounded by individual galaxies (e.g., Tsai \\& Mathews 1996). Dust grains may survive, and may even not mix well with, the X-ray-emitting gas in a quiescent, low density environment. But the mixing should likely to occur during or after the merger of a galaxy or group with a cluster. The dust grains would then be destroyed rapidly by sputtering in the ICM of high density and temperature. This may explain the metal abundance difference in the X-ray-emitting gas between the LSBXE and the central cluster.  \\subsection{Dynamical State of the Central  Cluster}  This central cluster (Abell 2125; Fig. 1) represents the strongest enhancement of the diffuse X-ray emission in the complex. The relative values of our measured temperature and luminosity of the  ICM (Table 2) are consistent with those for typical clusters (e.g., Wu et al. 1999). From the mass-temperature relationship (e.g., Shimizu et al. 2003), we may then expect the virial mass of the cluster as $2-5 \\times 10^{14} M_\\odot$. For a virialized system, the temperature would also predict a galaxy velocity dispersion of $\\sim 6.5 \\times 10^{2} {\\rm~km~s^{-1}}$. This velocity dispersion is considerably smaller than $\\sim 8.3 \\times 10^{2} {\\rm~km~s^{-1}}$ (assuming to be isotropic), inferred from the radial velocity dispersion of $\\sim 4.8 \\times 10^{2} {\\rm~km~s^{-1}}$ estimated from the modeling of the velocity field (Miller et al. 2004). Interestingly, the two brightest compact radio galaxies, 00047 and 00057, are moving at high velocities ($\\sim 1.5-1.8 \\times 10^3 {\\rm~km~s^{-1}}$) relative the mean of the cluster. It is conceivable that these two galaxies might simply represent a background cluster/group projected in the same sky by chance. The projected closeness of the galaxies to the center of the cluster and their enhanced diffuse X-ray emission, however, suggest that the association is physical, consistent with that the cluster is experiencing a major merger probably in a direction close to the line of sight. Indeed, 00047 is a strongly disturbed disk-like galaxy and the associated diffuse X-ray emission forms a trail, as expected from a ram-pressure stripping. 00057 and other two radio galaxies in the core of the cluster are all cD-like ellipticals (Figs. 7c and 15).  Furthermore, the 2-D morphology of the cluster in the sky, as shown in \\S 3.2, also suggests that it is far from a relaxed system. The X-ray morphology is strongly elongated and shows the centroid shift with scales, strongly indicating an ongoing merger (e.g., Roettiger et al. 1993). The measured X-ray ellipticity of the cluster is the largest in the Butcher \\& Oemler sample (Wang \\& Ulmer 1997). The X-ray centroid is further offset from the cD galaxies in the cluster core (Fig. 15). A northwest-southeast elongation is also seen in the galaxy distribution, similar to the X-ray morphology (Owen et al 2004b). One separate sub-cluster, at least, can be found in the NW extension of the cluster and is centered on a radio-loud cD-like galaxy (00039) and coincides with the bright soft X-ray source X-65 (Fig. 17a). This galaxy has a similar optical magnitude as those three cDs in the core. Furthermore, the velocity of the galaxy is comparable to those of the north and northeast cDs in the core, suggesting that the merger axis is nearly perpendicular to the line of sight. This northwest-southeast sub-cluster merger provides a natural explanation for the wide-angle tailed radio source attached to one (00106) of the two low-velocity ellipticals in the cluster core (Fig. 15), which suggests a pressure gradient or relative motion of the local gas in the plane of the sky. In short, the central cluster is undergoing a merger, possibly in multiple directions. As clusters are expected to be found at intersections of cosmic webs, such a complicated merger process is probably common for clusters still in early formation stages.  What might be the relationship between the LSBXE and the central cluster? The velocity centroids of the associated galaxy concentrations appear to be comparable (Miller et al. 2004). Spatially, the two concentrations are separated by $\\sim 7^\\prime$, or a projected distance $d_p \\sim $ 1.6 Mpc. One possibility is that the cluster and the LSBXE are merging in the plane of the sky. There are galaxies occupying the area between the LSBXE and the central cluster with the similar velocity (as well as a separate higher velocity component; Miller et al. 2004). There is also faint diffuse X-ray emission in the area, although the quality of the data does not allow for a distinction between a projection effect and a true dynamical interaction between these two concentrations. Alternatively, the two concentrations may be physically separated by a turn-around distance $d_{t}$ so that their relative velocity is near zero. Following Sarazin (2003), we estimated $d_{t} \\approx 2(GM)^{1/3} (t_t/\\pi)^{2/3} \\approx (7 {\\rm~Mpc}) (M/10^{15} M_\\odot)^{1/3} (t_t/10^{10} {\\rm yr})^{2/3}$, where $M$ is the total gravitational mass of the two concentrations, and $t_t$ is the age of the Universe at the turn-around time. If the system starts to collapse for the first time at $z = 0.247$, the axis is then oriented relative to the line of sight by an angle $\\sim {\\rm sin}^{-1}(d_p/d_t) \\sim 16^\\circ$ for $M \\sim 5 \\times 10^{14} M_\\odot$ (see above). If the two concentrations have already passed across each other (i.e., the first collapse occurred much earlier) and is about to re-collapse, the angle would then be larger.  \\subsection {Galaxy-Environment Interaction}  As marked in Table 1 and Fig. 7b, 10 X-ray sources are identified as member galaxies of the Abell 2125 complex. Eight of them are radio galaxies, preferentially in the LSBXE region. Most of these X-ray-loud galaxies are not resolved by {\\sl Chandra} and are probably dominated by AGN activities. Two of the X-ray sources are apparently resolved and are associated with the giant elliptical galaxy 00039 and the disturbed disk-like galaxy 00047 (C153) in the central Abell 2125 cluster (Figs. 15 and 17a,b). As our source detection algorithms are optimized to detect point-like sources, extended sources such as the one associated with 00057 (Fig. 15) is not listed in Table 1. These extended soft X-ray enhancements probably represent hot gas associated with individual massive galaxies or groups of galaxies, which may have entered the cluster for the first time. Outside the central cluster, the ambient density and relative velocity are typically low and the ram-pressure stripping is probably not important, at least for massive galaxies. The intergalactic gas around galaxies may even cool fast enough to replenish the gas consumed for star formation (Bekki et al. 2002). As they are plugging into a cluster, the surrounding gas may then be compressed by the high ram-pressure of the ICM, resulting in enhanced soft X-ray emission. The eventual stripping of the gas and dust may be important in both enriching the ICM and transforming the galaxies (e.g., Bekki et al. 2002.   \\begin{figure}[htb!] \\centerline{ } \\caption{\\protect\\footnotesize {\\sl HST} WFPC-2 V-band images of a field near the core of the Abell 2125 cluster (a and b) and a field in the LSBXE (c and d). The overlaid X-ray intensity contours are at 2.50,  2.75,  3.25,  4.00,  5.00,  6.25,  7.75,  9.5, 14.25, 26.75, and 51.75 for the 0.5-2 keV band (a); 2.80,  3.20,  4.00,  5.20,  6.80,  8.80,  11.2,  14 21.6,  41.6,  and 81.6  for the 2-8 keV band (b); 1.55,  1.6,      1.7,      1.85,      2.05,      2.3,      2.65,      3.9, 6.4,   and   11.4  for the 0.5-2 keV band (c); 2.1,      2.2,      2.4,      2.7,      3.1,      3.6,      4.3,      6.8, 11.8,   and   21.8  for the 2-8 keV band (d); all in units of $10^{-3} {\\rm~counts~s^{-1}~arcmin^{-2}}$. X-ray source numbers (Table 1) are labeled. } \\label{fig17} \\end{figure}   C153 represents an extreme case of the ram-pressure stripping. This galaxy probably went through the cluster central region quite recently (Fig. 15). The absence of the trail above 1.5 keV (Fig. 15) suggests that the gas in the trail is substantially cooler ($kT \\lesssim 1.5$ keV) than the ambient ICM, assuming collisional ionization equilibrium. (The the gas, if heated from a cool phase, may well be out of the equilibrium.) The emissivity of the hot gas in the ACIS-I 0.5-1.5 keV band  (Fig. 15) peaks at $kT \\sim 0.7$ keV, but is within a factor of 2 as long as $kT \\gtrsim 0.35$ keV. We estimate the luminosity of the trail as $\\sim 5 \\times 10^{41} {\\rm~ergs~s^{-1}}$ in the 0.5-2 keV band, which only weakly depends on the assumed gas temperature. To proceed further, we assume that the entire trail can be approximated as a uniform cylinder of $\\sim 88$ kpc long and  16 kpc diameter (Figs. 15-16; \\S 3.3). The mean electron density can then be estimated as $\\sim 1.0 \\times 10^{-2} {\\rm~cm^{-3}}\\xi^{-0.5} $, where $\\xi$ is the gas metallicity in the solar units. The dependence on $\\xi$ assumes that metal lines dominate the X-ray emission. If the X-ray-emitting gas in the trail is roughly in a pressure balance with the ICM ($p/k \\sim 7.7 \\times 10^4 {\\rm~cm^{-3}~K}$; \\S 3.2), the required gas temperature is then $kT \\sim 0.6 {\\rm~keV} \\xi^{0.5}$. We further estimate the total mass and  radiative cooling time scale of the gas in the trail as $\\sim 5 \\times 10^9 \\xi^{-0.5} M_\\odot$ and $ \\sim 1$ Gyr. If the gas is far from uniform in the trail, however, the cooling time could then be much shorter and may be compared to the crossing time scale of the galaxy through the cluster (a few times $10^8$ years).  A more thorough multiwavelength investigation of galaxy properties and their relationship to the local and global environment will be presented by Owen et al. (2004a,b). In particular, our X-ray source detection reported here is optimized for point-like sources and the detection threshold is quite conservative, which minimizes the global probability of including spurious sources in our list. With the priori positions of individual galaxies, one can lower the threshold for a positive detection of their X-ray counterparts. Using a threshold of $P = 10^{-3}$, for example, we have tentatively identified another 15 faint X-ray counterparts of the radio members of the Abell 2125 complex  (Owen et al. 2004a). These studies suggest that the Abell 2125 complex is in a special phase of the cluster formation, probably triggered by mergers among major components of a large-scale structure (e.g., Owen et al. 2004a). The presence of an unusually large number of active galaxies may be especially important for the heating of the IGM, affecting the subsequent evolution of the complex."
    },
    "0404/astro-ph0404328_arXiv.txt": {
        "abstract": "{We present high resolution spherically symmetric relativistic magnetohydrodynamical simulations of the evolution of a pulsar wind nebula inside the free expanding ejecta of the supernova progenitor. The evolution is followed starting from a few years after the supernova explosion and up to an age of the remnant of 1500 years. We consider different values of the pulsar wind magnetization parameter and also different braking indices for the spin-down process. We compare the numerical results with those derived through an approximate semi-analytical approach that allows us to trace the time evolution of the positions of both the pulsar wind termination shock and the contact discontinuity between the nebula and the supernova ejecta. We also discuss, whenever a comparison is possible, to what extent our numerical results agree with former self-similar models, and how these models could be adapted to take into account the temporal evolution of the system. The inferred magnetization of the pulsar wind could be an order of magnitude lower than that derived from time independent analytic models. ",
        "introduction": "\\label{sec:intro}  Pulsars are rapidly spinning magnetized neutron stars that usually form as the result of the core collapse of massive stars ($8-16\\ M_\\odot$) in supernova events (SN). The typical energy released in a supernova explosion is of order $\\sim10^{53}$ erg. Most of this is carried away by neutrinos, while only a small fraction (about 1\\%) goes into a blast wave that sweeps up the outer layers of the star and produces a strong shock propagating in the surrounding medium. The ejected material is initially heated by the blast wave and set into motion. Then, while the heat is converted into kinetic energy, the ejecta accelerate until the pressure becomes so low as to be dynamically unimportant. When this happens the material finally sets into homologous expansion (\\cite{chevalier89}; \\cite{matzner99}). This phase is usually referred to as free expansion of the ejecta.  As a consequence of the electromagnetic torques acting on it, the pulsar supplies a late energy input to the remnant in the form of a relativistic magnetized wind, mainly made of electron-positron pairs and a toroidal magnetic field (\\cite{goldreich69}; \\cite{michel99}). Most of the pulsar rotational energy is carried away by this wind, whose propagation velocity is ultra-relativistic, with typical Lorentz factors that, far enough from the light cylinder, are estimated to be in the range $~10^{4}$--$~10^{7}$. The interaction of the wind with the ejecta expanding at non-relativistic speed produces a reverse shock that propagates toward the pulsar (\\cite{rees74}). In the region bound by the wind termination shock on the inner side, and by the ejecta on the outer side, a bubble of relativistically hot magnetized plasma is created. This shines through synchrotron and Inverse Compton emission in a very broad range of frequencies, from radio wavelengths up to gamma rays: this is what we call a pulsar wind nebula (PWN) or plerion.  The evolution of a PWN inside the free expanding ejecta depends on many different parameters such as the pulsar luminosity, the density and velocity distribution in the SN ejecta (\\cite{dwarkadas89}; \\cite{featherstone01}; \\cite{blondin96}), the presence of large and/or small scale inhomogeneities (\\cite{chevalier89}; \\cite{campbel03}). In the case of constant luminosity, and if spherical symmetry is assumed, it is possible to derive a simple evolutionary equation for the radius of the PWN as a function of time, that results in a power law (\\cite{chevalier92}; \\cite{swaluw01}). In the case of SN ejecta with a constant density profile the PWN contact discontinuity evolves as $t^{6/5}$. For a more detailed description of the various phases of the PWN-SNR evolution see, for example, van der Swaluw et al.\\ (2001) and references therein.  While many analytic and numerical models exist in the literature for the evolution of SNRs, until recently only two classes of analytic models in the proper relativistic magnetohydrodynamical regime have been presented for PWNe: the steady state solution by \\cite{kennel84} (KC hereafter), and the self-similar solution by \\cite{emmering87} (EC hereafter), which allows for a non zero velocity of the termination shock.  Only lately the evolution of PWN-SNR systems has began being investigated through numerical simulations. These have been performed mainly in the classical hydrodynamical (HD) (\\cite{blondin01}; \\cite{swaluw01}), or classical MHD (\\cite{swaluw03}) regime. However the recent development of codes for relativistic magnetohydrodynamics (RMHD) allows one to investigate such systems in a proper regime and to quantify the accuracy of approximate analytic solutions (\\cite{bucciantini03}).  Both KC and EC models rely on two strong assumptions: a constant pulsar spin-down luminosity, and a constant velocity at the outer boundary of the nebula, neither of which applies to a real case nor is consistent with the PWN evolution inside free expanding ejecta. While both assumptions are known to be unrealistic, the most crucial one, as far as the long-term evolution of the system is concerned, is probably that of constant pulsar energy input. As we have already mentioned, the PWN is powered by the rotational energy lost by the star due to electromagnetic braking, and this loss translates into an increase with time of the pulsar rotation period (\\cite{lyne98}).  In the case of a dipolar magnetic field the torque exerted on the star results in the following relation between the spin-down rate and the pulsar frequency $\\Omega$ (e.g. \\cite{michel99}): \\be \\dot \\Omega \\propto -\\Omega^{3}; \\label{eq:omegadot} \\ee while the power supplied to the wind changes with time $t$ as: \\be -I \\Omega \\dot \\Omega = L(t) = \\frac{L_{o}}{(1+t/\\tau)^{2}}, \\label{eq:ltomega} \\ee where $I$ is the momentum of inertia of the pulsar, $\\tau$ is the characteristic spin-down time, and $L_o$ is the initial pulsar luminosity. More generally, if the field is not exactly dipolar one can write the pulsar (or wind) luminosity as: \\be L(t) = \\frac{L_{o}}{(1+t/\\tau)^{n}}, \\label{eq:lt} \\ee where $n=(\\beta+1)/(\\beta-1)$, with $\\beta$ the braking index.  Estimated values of $L_{o}$ may be up to $10^{38}$--$10^{40}$~erg/s. Determining the braking index from observations is extremely complicated, as it requires detailed and precise pulsar timing over long time-spans. A measure of $n$ is presently available for four pulsars only (\\cite{camilo00}). Among these, one, the Vela pulsar, has $n=6$, while all the others have $2<n<3$. The value of $\\tau$ can be derived from the pulsar period and spin-down rate once $n$ is known.  In the following we present 1D high resolution numerical simulations of the structure and evolution of a PWN inside free expanding SN ejecta. We extend, for this initial stage, that lasts for about 2000-3000 years (until the reverse shock propagating in the SNR collapses on the PWN), previous work on the subject by Bucciantini et al. (2003).  The plan of the paper is as follows. In Section~\\ref{sec:theor} we present a semi-analytic model for the evolution of the PWN radius and the pulsar wind termination shock, valid for any power-law profile of the ejecta and for any value of $n$. After briefly describing, in Section~\\ref{sec:numsim}, the numerical code employed and the initial conditions of our simulations, we present, in Section~\\ref{sec:numres}, the results obtained for different values of the braking index, $n=0,2,3$, and different magnetizations of the wind. We discuss these results and compare them with the expectations of the analytic EC model. A formula is derived for the spin-down factorization, and the new results are applied to the case of Crab Nebula. In Section~\\ref{sec:fin} we summarize our conclusions. ",
        "conclusions": "\\label{sec:fin} Our results show that, in the spherically symmetric approximation, the EC model gives a good description of the internal profiles of PWNe when the effects of spin-down can be neglected (i.e. when the central object is young and not very powerful). Noticeable discrepancies arise only at the outer boundary of the nebula, where the initial conditions and the evolutionary history of the system both play a major role. However, the region just behind the termination shock and further out, up to the position of the equipartition point, is well reproduced. This is the region of the nebula from which the high energy emission originates.  Even if the radial profiles of the various quantities are in good agreement, the ratio $R_{cd}/R_{ts}$ inferred from the EC model by matching the boundary velocity for a given $\\sigma$ is generally underestimated, due to the discrepancies at the boundary. Hence, if the ratio $R_{cd}/R_{ts}$ estimated in this way is compared with the observed value to evaluate the wind magnetization, one generally overestimates the value of $\\sigma$. A better agreement is obtained if $R_{cd}$ is taken equal to the radius at which the singularity of the EC model is found.  There is a good agreement in the HD case while a small difference, less than $10\\%$, is found in the magnetized cases, probably because in these cases $V_{cd} \\sim V_{asy}$. The EC model can be adapted to the evolution of $V_{cd}$ since, with respect to the sound speed crossing time in the PWN, variations at the border are very slow: changes of $V_{cd}$ are compensated by analogous variations of $V_{ts}$.  More generally, we expect that deviations from the spherically symmetric approximation might occur in the pulsar wind region, leading to the formation of a turbulent flow in the nebula that only multidimensional simulations can handle. In spite of being a simplification of the problem, 1D models can help understanding the importance of various processes like spin-down, mass loading or synchrotron losses, and to clarify how to keep them into account when comparing models with observations.  Once the spin-down is included, the ratio $R_{cd}/R_{ts}$ increases with respect to the constant luminosity case. Again, if this effect is not taken into account, the pulsar wind magnetization inferred from the PWN size might well be overestimated. With a simplified analytic model we are able to reasonably reproduce the ratio $R_{cd}/R_{ts}$ given by the simulations. We also found a fit to the results of our simulations in the form of a correction coefficient that should multiply the standard EC expectation. However, the radial profiles of the various quantities given by the EC model fail to reproduce the structure inside the PWN beyond the point where magnetic pressure reaches its maximum.  Given our fit, if the wind magnetization is known from the analysis of the non-thermal emission of the PWN, the size of the nebula can be used, in principle, to estimate $n$ and $\\tau$. Viceversa, if the pulsar spin-down properties are known, from the same expression it is possible to estimate $\\sigma$.  As shown in the case of the Crab Nebula, inclusion of the spin-down effect reduces the estimated $\\sigma$ by about one order of magnitude with respect to the value found by KC. We want to stress that this is not just a dynamical problem related to the size of the nebula alone. A less magnetized wind creates a less magnetized nebula. Synchrotron losses are less efficient and high energy particles have a longer life and can move farther from the termination shock. This suggests that the interpretation of synchrotron maps should take into account the new results we have presented."
    },
    "0404/astro-ph0404420_arXiv.txt": {
        "abstract": "We report on results from a 2003 February {\\em INTEGRAL} observation of the black hole candidate (BHC) 4U 1630--47 and the high energy sources in the Norma Region of the Galaxy and on {\\em RXTE} and radio observations of 4U 1630--47.  To date, there have been 38 new or previously poorly studied sources found by {\\em INTEGRAL} (the ``IGR'' sources), and 15 of the IGR sources are in the field of view of our observation.  We detect 12 of the 15 sources, and we speculate that a subset of the IGR sources may be part of a new population of persistent, hard, intermediate luminosity X-ray sources.  During our observations, 4U 1630--47 was in an intermediate state with strong soft and hard components, a power-law index of $\\Gamma = 2.11$-2.18, a 4\\% rms level of timing noise (0.1-10 Hz), and the source was not detected in the radio band.  We discuss the evolution of the energy spectrum during the 293 ks {\\em INTEGRAL} observation, and we report on rapid, high amplitude flaring behavior that has been seen for much of the 2002-2004 outburst. ",
        "introduction": "The X-ray flux observed from transient X-ray binaries changes by several orders of magnitude between quiescence and outburst, providing an opportunity to learn about black hole candidates (BHCs), accretion onto compact objects, and the processes responsible for high energy emission.  Although these objects are known to emit hard X-rays and gamma-rays, the origin of this emission is not well-understood. The high energy emission provides a probe of the regions of space close to the compact object, allowing us to study strong gravity and jet formation processes.  4U 1630--47 is among the most active of the BHC X-ray transients and has produced strong hard X-ray emission during its 17 detected outbursts \\citep{tk00,oosterbroek98,kuulkers97,parmar97} as well as highly polarized radio emission \\citep{hjellming99}, indicating the presence of jets.  The current outburst from 4U 1630--47, which began in 2002 September, is one of the brightest and longest recorded outbursts from this system, and we had the opportunity to observe 4U 1630--47 with {\\em INTEGRAL} in 2003 February.  Due to the large {\\em INTEGRAL} field of view, we also detect many other hard X-ray sources in the vicinity of 4U 1630--47. Several of the sources were unknown or poorly studied prior to {\\em INTEGRAL} and have ``IGR'' designations.  Although we are just beginning to understand the nature of the IGR sources, many of these appear to be highly absorbed ($N_{\\rm H} = 10^{23-24}$ cm$^{-2}$), and some of them are bright IR sources \\citep[e.g.,][and papers by Foschini, Walter, and others in these proceedings]{mg03,walter04,patel04}.  The concentration of IGR sources near 4U 1630--47 is almost certainly due to the location of the source in the Galactic plane ($l = 336.91^{\\circ}$, $b = +0.25^{\\circ}$).  This is in the Norma region of the Galaxy, and the line of sight to 4U 1630--47 is tangent to the Norma-Scutum Galactic arm.  Below, we present our results on:  The hard X-ray sources detected by {\\em INTEGRAL} in the Norma region; and X-ray and radio observations of 4U 1630--47. ",
        "conclusions": "\\noindent $\\bullet$~Fifteen of the 38 IGR sources are in the 4U 1630--47 field. Only three of the sources are not detected, indicating that as many as 12 of the sources are persistent.  It is possible that a subset of the IGR sources are part of a new population of persistent, hard, intermediate luminosity X-ray sources.\\\\ $\\bullet$~For much of its current outburst, 4U 1630--47 has shown dramatic flaring that is reminiscent of the flaring seen from GRS 1915+105.  For 4U 1630--47, the hardness and intensity are correlated during the flares, indicating that these are hard flares that are related accretion (and possibly ejection) mechanisms rather than being due to, e.g., absoption from the outer edge of the disk.\\\\ $\\bullet$~The ISGRI light curve (Figure~\\ref{fig:lc3}a) and the spectra shown in Figure~\\ref{fig:evolution} show that 4U 1630--47 spectrum gradually changed, and the main change is in the normalization of the power-law component.\\\\ $\\bullet$~The {\\em RXTE} spectra indicate that the high energy cutoff also changes during the observations, and this sort of evolution must be kept in mind when combining {\\em INTEGRAL} data from long observations."
    },
    "0404/astro-ph0404566_arXiv.txt": {
        "abstract": "Cosmological gamma ray bursts (GRBs) are thought to occur from violent hypercritical accretion onto stellar mass black holes, either following core collapse in massive stars or compact binary mergers. This dichotomy may be reflected in the two classes of bursts having different durations.  Dynamical calculations of the evolution of these systems are essential if one is to establish characteristic, relevant timescales. We show here for the first time the result of dynamical simulations, lasting approximately one second, of post--merger accretion disks around black holes, using a realistic equation of state and considering neutrino emission processes. We find that the inclusion of neutrino optical depth effects produces important qualitative temporal and spatial transitions in the evolution and structure of the disk, which may directly reflect upon the duration and variability of short GRBs. ",
        "introduction": "The coalescence of two compact objects due to the emission of gravitational radiation is a possible origin of cosmological gamma ray bursts of the short variety \\citep[see e.g.,][for a review]{m02}, lasting a few tenths of a second \\citep{ls76,bp86,elps89,npp92}. Previous multidimensional studies have shown that the outcome of such a merger is in all likelihood a dense torus surrounding a supramassive neutron star, likely to collapse to a black hole on a short timescale, or a black hole, if one was already present in the system \\citep[e.g][]{rs94,kl98,rj99,rrd03}. This class of systems has been studied observationally since the 1970s, when the Hulse--Taylor system was discovered \\citep{ht75}, and the recent detection of PSR J0737-3039 \\citep{burgay03}, the tightest binary yet in this class, has only heightened interest in these events. The final instants before the actual merger will produce a powerful burst of gravitational waves, and the merger waveform itself is expected to reveal details about the equation of state of matter at high densities \\citep[see e.g.,][]{thorne95}.  The merger process is intrinsically a three--dimensional event, with no particular symmetry which can be exploited to reduce the complexity of the problem, and must be addressed numerically. No such calculation is currently evolved for more than a few tens of milliseconds, because of computational limitations. Information about the state of the system after the violent, dynamical merger is essentially over can be obtained, but the final state of the system, on a timescale comparable to the duration of a short GRB, cannot be addressed.  The steady state structure and composition of post--merger accretion disks has been the focus of several recent studies \\citep{pwf99,npk01,km02,dpn02,pwh03,b03}, which have included an increasing amount of microphysical detail. However, the dynamics of the problem have never been followed for an extended amount of time, although progress is being made in this direction \\citep{srj04}.  In a previous paper \\citep{lrr02}, we studied the dynamical evolution of the accretion structures formed after a black hole--neutron star merger under a set of simplifying assumptions. Initial conditions were taken directly from the late stages of the coalescence, modeled in three dimensions using a simple equation of state \\citep{l01}. The disk was then mapped onto two dimensions assuming azimuthal symmetry, and its evolution was followed for 0.2 seconds, comparable to the duration of a short GRB. The equation of state was simple (ideal gas with adiabatic index $\\gamma=4/3$), and we assumed that all the energy dissipated by viscosity (modeled with the full expression of the stress tensor, using an $\\alpha$--prescription for the magnitude of the viscosity) was radiated away.  In this {\\em Letter} we present for the first time dynamical calculations of post--merger accretion disks in two dimensions, $(r,z)$, that: (1) use a realistic and self--consistent equation of state for the fluid in the disk; (2) consider the cooling of the disk through neutrino emission, using appropriate rates; (3) account for the effects of neutrino opacities and the corresponding optical depths through a simplified treatment and (4) last for a minimum of 0.4 seconds, comparable to the duration of a short GRB. ",
        "conclusions": "\\label{disc}  The evolution of accretion disks resulting from dynamical three dimensional binary coalescence calculations, where a neutron star is tidally disrupted before being swallowed by its black hole companion, is studied numerically. Angular momentum transport and the associated energy dissipation are modeled using an $\\alpha$ prescription. By assuming azimuthal symmetry we are able to follow the time dependence of the disk structure for a fraction of a second, a time comparable with the duration of short GRBs. During this time, the rate of mass supply to the central black hole is of the order of a fraction of a solar mass per second; i.e. much greater than the Eddington rate. Although the gas photon opacities are large, the disk becomes sufficiently dense and hot to cool via neutrino emission. There is in principle no difficulty in dissipating the disk internal energy, but the problem is in allowing these neutrinos to escape from the Thomson thick inflowing gas.  At sufficiently low accretion rates, $\\alpha \\lesssim 0.01$, we find that the energy released by viscous dissipation is almost completely radiated away on a timescale given by $t_{\\rm cool} \\approx E_{\\rm int}/L_{\\nu} \\sim 0.1$~s. In contrast, for a higher mass supply, $\\alpha \\gtrsim 0.1$, energy advection remains important until the entire disk becomes optically thin, at $t[\\tau_\\nu=1]\\sim 30$ ms. The restriction on the cooling rate imposed by high optical depths is key because it allows the energy loss to be spread over an extended period of time (i.e.  $t_{\\rm cool}$ or $t[\\tau_\\nu=1]$) during which the neutrino luminosity stays roughly constant.  In principle, neutrinos could give rise to a relativistic pair dominated wind if they converted into $e^\\pm$ pairs in a region of low baryon density. This gives a characteristic timescale for energy extraction, and may be essential for determining the duration of neutrino-driven, short GRBs.  An alternative way to tap the torus energy is via magnetic fields threading the disk: the energy liberated by accretion is converted efficiently into magnetic form and emitted as a magnetically dominated outflow. The launching of a jet probably requires the existence of a poloidal magnetic field of magnitude $B_p$ over a scale of radius $R$, where $B_p$ is smaller than the field strength $B_{\\rm disk}$ associated with the dynamo-driven magnetic disk viscosity \\citep{lop99}. This results directly from estimating the accretion and jet luminosities as $L_{\\rm acc}\\sim\\dot{M}_{\\rm acc}v_{\\phi}^{2}$, and $L_{\\rm jet}\\sim B_{\\rm p}^{2}R^{2}v_{\\phi}$ respectively, which gives $(B_{\\rm p}/B_{\\rm disk})^{2}\\sim (L_{\\rm jet}/L_{\\rm acc})(H/R)$. Assuming the scaling $B_{\\rm p}\\sim (H/R)B_{\\rm disk}$, derived by \\citet{tp96}, one obtains $L_{\\rm jet}/L_{\\rm acc}\\sim H/R$. In order to achieve $L_{\\rm jet} \\sim L_{\\rm acc}$, \\citet{lpk03} argue that a number $R/H$ of neighboring annuli in the disk need to provide a locally generated net poloidal field in the same direction. They notice that since the local dynamos vary on timescales of $t_{\\rm dyn}$, the timescale for establishing a change in the magnetic field in the disk should be of the order $t_{\\rm jet} \\sim t_{\\rm dyn}2^{R/H}$. In our case $R/H\\sim 3$, which gives $t_{\\rm jet}\\sim 10-100$ ms. This corresponds to the typical variability timescale of short duration GRBs \\citep{rf00,np02}.  Two further physical effects, noted previously in a more general context, are relevant for the present study. First, the mechanism by which a jet is launched is not scale--free \\citep{pringle93}, and thus the poloidal field must be established in the inner regions of the disk. Second, in these same regions, the inflow speed caused by angular momentum loss to the outflow is increased, up to a factor $R/H$ over the value expected from magnetic viscosity alone \\citep{lpk03}. Thus the poloidal flux generated by the dynamo can be effectively trapped. This probably implies that once the dynamo process generates a global poloidal field, it is able to maintain it during the lifetime of the thick portion of the disk, $t[\\tau_\\nu=1]$. For the typical parameters studied here, this timescale could be of the order of seconds when the rate of mass supply is low (i.e. $\\alpha\\sim 10^{-3}$). There is thus no problem in principle in accounting for sporadic large-amplitude variability on timescales as short as one millisecond, even in the most long-lived short GRBs."
    },
    "0404/astro-ph0404085_arXiv.txt": {
        "abstract": "We measure the angular power spectrum $C_\\ell$ of radio galaxies in the NRAO VLA Sky Survey (NVSS) using two independent methods: direct spherical harmonic analysis and maximum likelihood estimation.  The results are consistent and can be understood using models for the spatial matter power spectrum and for the redshift distribution of radio galaxies at mJy flux-density levels.  A good fit to the angular power spectrum can only be achieved if radio galaxies possess high bias with respect to mass fluctuations; by marginalizing over the other parameters of the model we derive a $68\\%$ confidence interval $1.53 < b_0 \\times \\sigma_8 < 1.87$, where $b_0$ is the linear bias factor for radio galaxies and $\\sigma_8$ describes the normalization of the matter power spectrum.  Our models indicate that the majority of the signal in the NVSS $C_\\ell$ spectrum is generated at low redshifts $z \\la 0.1$.  Individual redshifts for the NVSS sources are thus required to alleviate projection effects and probe directly the matter power spectrum on large scales. ",
        "introduction": "\\renewcommand{\\thefootnote}{\\fnsymbol{footnote}} \\setcounter{footnote}{1} \\footnotetext{E-mail: chrisb@phys.unsw.edu.au}  \\begin{figure*} \\center \\epsfig{file=nvsssurv.ps,height=12cm,angle=-90} \\caption{NVSS sources with $S_{\\rm 1.4 \\, GHz} > 200$ mJy in an equal-area projection.  The Galactic plane and Galactic latitudes $\\pm 5^\\circ$ are also plotted; sources within this region are masked from our large-scale structure analysis as many are Galactic in origin.} \\label{fignvss} \\end{figure*}  Active Galactic Nuclei (AGN) mapped in radio waves are an interesting probe of large-scale structure.  They can be routinely detected out to very large redshift ($z \\sim 4$) over wide areas of the sky and hence delineate the largest structures and their evolution over cosmic epoch.  Radio emission is insensitive to dust obscuration and radio AGN are effective tracers of mass: they are uniformly hosted by massive elliptical galaxies and have been shown to trace both clusters (Hill \\& Lilly 1991) and superclusters (Brand et al. 2003).  The current generation of wide-area radio surveys such as Faint Images of the Radio Sky at Twenty centimetres (FIRST; Becker, White \\& Helfand 1995) and the NRAO VLA Sky Survey (NVSS; Condon et al. 1998) contain radio galaxies in very large numbers ($\\sim 10^6$) and have allowed accurate measurements of the imprint of radio galaxy angular clustering.  These patterns are considerably harder to detect in radio waves than in optical light due to the huge redshift range that is probed.  Whilst this provides access to clustering on the largest scales, it also washes out much of the angular clustering signal through the superposition of unrelated redshift slices.  The angular correlation function was measured for FIRST by Cress et al. (1996) and Magliocchetti et al. (1998) and for NVSS by Blake \\& Wall (2002a) and Overzier et al. (2003).  The NVSS radio survey, covering $\\sim 80$ per cent of the sky, permits the measurement of fluctuations over very large angles.  Blake \\& Wall (2002b) detected the imprint of the cosmological velocity dipole in the NVSS surface density, in a direction consistent with the Cosmic Microwave Background (CMB) dipole.  In this study we measure the angular power spectrum, $C_\\ell$, of the radio galaxy distribution (Baleisis et al. 1998).  This statistic represents the source surface-density field as a sum of sinusoidal angular density fluctuations of different wavelengths, using the spherical harmonic functions.  The angular power spectrum is sensitive to large-angle fluctuations and hence complements the measurement of the angular correlation function, $w(\\theta)$, at small angles.  Measurement of the $C_\\ell$ spectrum has some advantages in comparison with $w(\\theta)$.  Firstly, the error matrix describing correlations between multipoles $\\ell$ has a very simple structure, becoming diagonal for a complete sky.  This is not the case for the separation bins in a measurement of $w(\\theta)$: even for a full sky, an individual galaxy appears in many separation bins, automatically inducing correlations between those bins.  Secondly, there is a natural relation between the angular power spectrum and the spatial power spectrum of density fluctuations, $P(k)$.  This latter quantity provides a very convenient means of describing structure in the Universe for a number of reasons.  Firstly its primordial form is produced by models of inflation, which prescribe the initial pattern of density fluctuations $\\delta\\rho/\\rho$.  Furthermore, in linear theory for the growth of perturbations, fluctuations described by different wavenumbers $k$ evolve independently, enabling the model power spectrum to be easily scaled with redshift.  The physics of linear perturbations are hence more naturally described in Fourier space.  In contrast, the angular correlation function is more easily related to the spatial correlation function $\\xi(r)$, the Fourier transform of $P(k)$.  Correlation functions more naturally serve to describe the real-space profile of collapsing structures evolving out of the linear regime.  We emphasize that although the two functions $C_\\ell$ and $w(\\theta)$ are {\\it theoretically} equivalent -- linked by a Legendre transform -- this is not true in an {\\it observational} sense.  For example, $w(\\theta)$ can only be successfully measured for angles up to a few degrees, but $C_\\ell$ depends on $w(\\theta)$ at {\\it all} angles.  We derive the angular power spectrum using two independent methods. Firstly we apply a direct spherical harmonic estimator following Peebles (1973).  Secondly, we use maximum likelihood estimation, commonly employed for deriving the angular power spectra of the CMB temperature and polarization maps.  These two methods are described in Section \\ref{secmeth}.  We find that these two approaches yield very similar results (Section \\ref{secres}), which is unsurprising given the wide sky coverage of the NVSS.  In Section \\ref{secpk} we interpret the NVSS angular power spectrum in terms of the underlying spatial power spectrum of mass fluctuations and the radial distribution of radio sources.  Finally in Section \\ref{secbias} we employ these models to derive the linear bias factor of NVSS radio galaxies by marginalizing over the other model parameters. ",
        "conclusions": "This investigation has measured the angular power spectrum of radio galaxies for the first time, yielding consistent results through the application of two independent methods: direct spherical harmonic analysis and maximum likelihood estimation.  The NVSS covers a sufficient fraction of sky ($\\sim 80$ per cent) that spherical harmonic analysis is very effective, with minimal correlations amongst different multipoles.  The form of the $C_\\ell$ spectrum can be reproduced by standard models for the present-day spatial power spectrum and for the radial distribution of NVSS sources -- provided that this latter is modified at low redshift through comparison with optical galaxy redshift surveys.  The results strongly indicate that radio galaxies possess high bias with respect to matter fluctuations. A constant linear bias $b_0 \\approx 1.7$ permits a good fit, and by marginalizing over the other parameters of the model we deduce a $68\\%$ confidence interval $b_0 \\, \\sigma_8 = 1.53 \\rightarrow 1.87$ where $\\sigma_8$ describes the normalization of the matter power spectrum.  We find that the majority of the angular power spectrum signal is generated at low redshifts, $z \\la 0.1$.  Therefore, in order to exploit the potential of radio galaxies to probe spatial fluctuations on the largest scales, we require individual redshifts for the NVSS sources."
    },
    "0404/astro-ph0404200_arXiv.txt": {
        "abstract": "We present a method to correct the chromatic and airmass dependent extinction for N-band spectra taken with the TIMMI2 instrument at the ESO\\,/\\,La~Silla observatory. Usually, the target and calibrator star have to be observed at similar airmass in order to obtain reliable spectrophotometric fluxes. Our method allows to correct the atmospheric extinction and substantially improves the spectrophotometric flux calibration, when the standard stars were observed at a very different airmass than the targets. Hundreds of standard star measurements in several passbands (N1, N8.9, N10.4, N11.9) were used to derive mid-IR extinction coefficients. We demonstrate that applying our correction of the differential extinction to test data results in a spectrophotometric accuracy up to 2\\% within the literature flux. ",
        "introduction": "For spectroscopy the target stars and calibrators should generally be observed at a comparable airmass to avoid a varying extinction between the two objects. Different to near-IR observations, in the mid-IR a good calibrator at the same airmass is often difficult to find and the theoretical behaviour of the extinction with airmass and wavelength is not really known up to now. Some authors claim there would be no clear dependence and thus for photometry no airmass correction needs to be applied. However these authors used dozens of standard star observations, while we apply several hundred measurements.  In mid-IR spectroscopy the influence of extinction is more evident than for photometry. No good spectral flux calibrations are possible unless the target and calibrator star are observed close in time and very close in airmass ($\\leq$~0.1 airmass distance). Therefore, some observers often use additional photometric measurements to correct the spectral fluxes. Since the extinction as a function of wavelength varies significantly within the N-band, also the spectral slope needs to be corrected, if the target and calibrator had been taken at different airmass.  Our goal is to explain the extinction in the mid-IR as a function of airmass. We further demonstrate that the extinction has a non-linear wavelength dependence within the N-band. ",
        "conclusions": ""
    },
    "0404/astro-ph0404346_arXiv.txt": {
        "abstract": "We have imaged the circumstellar envelope around the binary protostar L1551 IRS 5 in CS ($J$=7--6) and 343 GHz continuum emission at $\\sim$ 3$\\arcsec$ resolution using the Submillimeter Array. The continuum emission shows an elongated structure ($\\sim220 \\times 100$ AU) around the binary perpendicular to the axis of the associated radio jet. The CS emission extends over $\\sim400$~AU, appears approximately circularly symmetric, and shows a velocity gradient from southeast (blueshifted) to northwest (redshifted). The direction of the velocity gradient is different from that observed in C$^{18}$O ($J$=1--0). This may be because rotation is more dominant in the CS envelope than the C$^{18}$O envelope, in which both infall and rotation exist. The CS emission may be divided into two velocity components: (1) a ``high\" velocity disk-like structure surrounding the protostar, $\\pm1.0-1.5$~km~s$^{-1}$ from the systemic velocity, and (2) a ``low\" velocity structure, located southwest of the protostar, $<1.0$~km~s$^{-1}$ from the systemic velocity. The high-velocity component traces warm and dense gas with kinematics consistent with rotation around the protostar. The low-velocity component may arise from dense gas entrained in the outflow. Alternatively, this component may trace infalling and rotating gas in an envelope with a vertical structure. ",
        "introduction": "It is widely accepted that low-mass stars are formed through mass accretion in protostellar envelopes \\cite{and00,mye00}.  The dense and warm innermost part of low-mass protostellar envelopes (radius $<$ 200 AU) are likely sites of formation of protoplanetary disks \\cite{spa95,bec96}.  However, we have very little knowledge about this region because previous mm-wave observations could not separate the warm and dense regions from the overlying low-density ($\\sim10^{4-5}$ cm$^{-3}$) and cold ($\\sim$10 K) gas along the line of sight. On the other hand, submillimeter molecular lines such as CS ($J$=7--6) trace higher densities ($>$ 10$^{\\rm 7}$ cm$^{\\rm -3}$) and temperatures ($>$ 60 K) \\cite{gms95,spa95}, and submillimeter dust continuum emission is an excellent tracer of disks around protostars \\cite{oso03}.  In this $Letter$, we describe results of CS ($J$=7--6) and 343 GHz continuum observations of L1551 IRS 5 with the Submillimeter Array (SMA)\\footnote{The Submillimeter Array (SMA) 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.}.  L1551 IRS 5 is the brightest protostar ($L_{bol}$ $\\sim30 L_{\\odot}$; Keene \\& Masson 1990) in Taurus ($D \\sim140$ pc; Elias 1978). After the first discovery of the bipolar molecular outflow by Snell, Loren \\& Plambeck (1980), Kaifu et~al.\\ (1984) discovered a 0.1 pc scale rotating gas structure around L1551 IRS 5 with the Nobeyama 45 m telescope.  The elongated rotating protostellar envelope around L1551 IRS 5 was first imaged by Sargent et~al.\\ (1988) in C$^{18}$O ($J$=1--0) with the OVRO mm-array.  Ohashi et~al.\\ (1996), Saito et~al.\\ (1996), and Momose et~al.\\ (1998) have found evidence of infalling motion in the protostellar envelope with the Nobeyama Millimeter Array (NMA). Millimeter interferometric observations of dust continuum emission from L1551 IRS 5 revealed that L1551 IRS 5 is a close ($\\sim0\\farcs3$) binary \\cite{loo97,rod98}.  Lay et~al.\\ (1994) made the first submillimeter interferometric observations of dust continuum emission from L1551 IRS 5 with the JCMT-CSO interferometer. Single-dish observations of L1551 IRS 5 in submillimeter molecular lines have revealed warm ($\\geq$ 40 K) and dense ($\\geq$ 10$^{6}$ cm$^{-3}$) gas in the inner part of the envelope \\cite{ful95,gms95,hog97,hog98}.  Our SMA observations provide us with new information of the innermost part of the protostellar envelope. ",
        "conclusions": "\\subsection{Physical Conditions of the CS Emitting Region}  Our CS ($J$=7--6) observations have revealed a rather compact ($\\sim400$ AU) structure in the inner part of the elongated ($\\sim2000$ AU) protostellar envelope detected by mm-wave observations \\cite{oh96b,sai96,mom98} around L1551 IRS 5.  The critical gas density traced by the CS line is expected to be $\\geq$ 10$^{\\rm 7}$ cm$^{\\rm -3}$, and the equivalent temperature of the $J$=7 energy level is 66 K.  This gas density and temperature are much higher than those traced by the millimeter-wave tracers, such as C$^{\\rm 18}$O ($J$=1--0) ($\\sim10^{4}$ cm$^{\\rm -3}$; 5 K) \\cite{sar88,mom98} or H$^{\\rm 13}$CO$^{\\rm +}$ ($J$=1--0) ($\\sim10^{5}$ cm$^{-3}$; 4 K) \\cite{sai96}.  Indeed, Moriarty-Schieven et~al.\\ (1995) estimated gas density and temperature at the center of the envelope in L1551 IRS 5 to be $\\sim10^{7}$ cm$^{-3}$ and $\\sim40$ K from their multi-transitional CS ($J$=3--2, 5--4, 7--6) observations with JCMT and CSO. Part of the difference in the distribution between the CS ($J$=7--6) and C$^{\\rm 18}$O ($J$=1--0) emission shown in Figure 1 is likely to be due to the different physical conditions traced by these two molecular lines.  The CS emission line better traces the inner, denser and warmer part of the protostellar envelope.  The compact shape of the CS ($J$=7--6) emission without elongation could be due to the missing flux (only $\\sim$ 11$\\%$ recovered).  However, our simple simulation indicated that, if the distribution of the CS ($J$=7--6) emission were the same as that of the C$^{18}$O ($J$=1--0) emission, then the present SMA observations would not detect any of the CS emission. This suggests that there should be a compact structure of the CS emission as well as an extended CS component with $\\geq$ 11$\\arcsec$ = 1500 AU in size. We note that it is still a puzzle why the CS ($J$=7--6) emission which traces denser and warmer gas can be so extended.  It seems to be difficult to make gas temperature high enough only by the heating from the central stars (e.g., Lay et~al.\\ 1994).  There might be some external heating sources such as shocks due to infalling and/or outflowing gas \\cite{nak00,ave96}.  \\placefigure{f4}   \\subsection{Origin of the High- and Low-Velocity Components}  As described in $\\S$3, the CS ($J$=7--6) emission in L1551 IRS 5 consists of a high-velocity disk-like component and a low-velocity component whose peak locates at the southwest of the protostar.  In order to investigate the origin of these velocity components in more detail, in Figure 4 we show the position-velocity (P-V) diagram of the CS emission along the cut perpendicular to the axis of the radio jet (see Figure 3).  The velocity structure appears to be symmetric with respect to the velocity of 6.8 km s$^{-1}$, and we adopt this velocity as a systemic velocity in this paper (Note that this is significantly different from the systemic velocity of C$^{\\rm 18}$O ($J$=1--0) of 6.2 km s$^{-1}$; Momose et~al.\\ 1998). Solid color curves in Figure 4 indicate the Keplerian rotation velocity ($\\propto$ $r^{-0.5}$) for a disk with an inclination angle of 65$^{\\circ}$ \\cite{mom98} and with a central stellar mass of 0.15 M$_{\\odot}$ (red), 0.5 M$_{\\odot}$ (green), and 1.2 M$_{\\odot}$ (blue). The central stellar mass of 0.15 M$_{\\odot}$ was estimated by Momose et~al.\\ (1998) from their detailed kinematic analyses of the C$^{\\rm 18}$O infalling envelope, and that of 1.2 M$_{\\odot}$ by Rodr\\'{\\i}guez et~al.\\ (2003a) from the proper motion analyses of the binary.  Dashed curves in Figure 4 show rotation with angular momentum conserved, that is, $V_{rot} = 0.24(r/700{\\rm AU})^{-1}~~{\\rm km\\ s^{-1}}$. Momose et~al.\\ (1998) argued that this rotation law can explain the observed P-V diagram of the C$^{\\rm 18}$O infalling envelope.  The high-velocity component, with relative velocities of $\\pm$ 1.0--1.5 km s$^{-1}$ in Figure 4, seems to be consistent with the Keplerian rotation with a central stellar mass of 0.15 M$_{\\odot}$ or the rotation with the constant angular momentum, and is likely to be the rotating circumbinary disk at the innermost of the envelope.  Unfortunately, it is difficult for us to distinguish these two rotation curves to explain the kinematics of the high-velocity component with the current sensitivity.  Keplerian rotation curves with higher ($>$ 1.0 M$_{\\odot}$) stellar masses show too high velocities as compared to the CS P-V diagram.  In contrast, the low-velocity component ($\\pm$0.5 km s$^{-1}$ from the systemic velocity) seems not to follow these rotation curves.  The spatial offset of the low-velocity component from the dust disk and the possible coincidence with the radio jet (Figure 3) may suggest that the origin of the low-velocity component is outflow-related.  Interestingly, at the south-western tip of the radio jet where the peaks of the low-velocity component locate, Bally, Feigelson, \\& Reipurth (2003) have detected X-ray emission with $Chandra$, which they attribute to shocks caused by the jet. One interpretation therefore is that the low-velocity component is the shocked warm and dense molecular material pushed by the radio jet. However, one difficulty of this interpretation may be its ``low velocity''. The projected separation between the peaks of the low-velocity component and the continuum peak is $\\sim$ 100 AU, and the line of sight velocities of the low-velocity component is only $\\sim$ a few $\\times$ 0.1 km s$^{-1}$. In order for the low-velocity component to escape from the gravitational potential of the central star with a mass of 0.15 M$_{\\odot}$ \\cite{mom98}, the angle between the flow and the plane of sky must be less than 10$^{\\circ}$, while Uchida et~al.\\ (1987) estimated the inclination angle of 15$^{\\circ}$ from their large-scale CO outflow map.  It is important to compare our P-V diagram with those calculated using a model envelope with a vertical structure, because models of flared envelopes having infall as well as rotation yield peaks at lower velocity near the central star in their P-V diagrams cut along the major axis \\cite{mom98,hog01}.  The reason why these peaks at lower velocity appear near the central star is because infalling motions in the near- and far-side of the envelope are observed together with rotation in the same line of sight. These peaks cannot be explained with rotation curves calculated using spatially thin disk models, which is very similar to our case shown in Figure 4. It is therefore possible that a model envelope having vertical structures with infall and rotation motions may explain the lower velocity component.  The effect of the missing flux prevents us from arriving at a clear interpretation.  The lower-velocity component could be part of spatially extended component, which can often be seen at lower velocities.  Further observations with shorter spacings using the SMA should provide us with the more obvious view and interpretation of these velocity structures of warm and dense gas in the innermost part of the protostellar envelope."
    },
    "0404/astro-ph0404170_arXiv.txt": {
        "abstract": "We point out novel consequences of general relativity involving tidal dynamics of ultrarelativistic relative motion. Specifically, we use the generalized Jacobi equation and its extension to study the force-free dynamics of relativistic flows near a massive rotating source. We show that along the rotation axis of the gravitational source, relativistic tidal effects strongly \\emph{decelerate} an initially ultrarelativistic flow with respect to the ambient medium, contrary to Newtonian expectations. Moreover, an initially ultrarelativistic flow perpendicular to the axis of rotation is strongly \\emph{accelerated} by the relativistic tidal forces. The astrophysical implications of these results for jets and ultrahigh energy cosmic rays are briefly mentioned. ",
        "introduction": "\\label{s1} This paper is about the motion of free test particles in the background gravitational field of a Kerr black hole. The geodesics of Kerr spacetime are well known and have been extensively studied in standard systems of coordinates such as the Boyer-Lindquist coordinate system. The Boyer-Lindquist coordinates are adapted to the stationary and axisymmetric nature of the source and are Minkowskian (expressed in spherical polar coordinates) infinitely far from the black hole. We are interested however in the \\emph{relative} motion of free nearby particles; as explained in detail in~\\cite{1}, from the standpoint of relativity theory the study of relative motion is certainly preferred, since it is in keeping with the spirit of the theory. Moreover, relative motion has direct observational significance. Thus theoretical results obtained in the study of relative motion can be directly compared with suitable experimental data~\\cite{1,2}.  To study the relative motion of free particles, we need a reference particle whose motion is well known. For the purposes of the present work, we select the free motion of particles on radial escape trajectories along the rotation axis of the Kerr black hole. Thus at a given initial position along a radial direction in Schwarzschild spacetime or the rotation axis in Kerr spacetime, the reference particle has an escape velocity such that the particle eventually reaches infinity with zero kinetic energy. We then study the relative motion of the relativistic particles that start from the same location as the reference particle and travel in different directions. The origin of these relativistic particles on the axis of rotation will not be discussed in this paper. We will simply assume that such particles exist near the poles of the Kerr black hole. The nature of the central engine and the details of the accretion phenomena that could generate such relativistic particles are beyond the scope of this work. Instead, we concentrate on the gravitational dynamics of these particles relative to the reference particle. Moreover, it turns out that the existence of event horizons has no bearing on our main results; therefore, our treatment applies equally well to test particles in the exterior of a gravitational source whose field is adequately described by the exterior Kerr spacetime. However, our results are in general observationally significant only in the case of highly collapsed gravitational sources such as neutron stars and black holes.  To characterize the relative motion, we establish a quasi-inertial coordinate system along the worldline of the reference particle. This Fermi normal coordinate system~\\cite{3} is valid in a cylindrical region of radius $\\mathcal{R}$ along the worldline of the reference particle. Here $\\mathcal{R}$ is the radius of curvature of the spacetime manifold. The geodesic equations of motion for the free particles are then integrated in this Fermi coordinate system subject to the limitation that the relative distance must remain well within the admissible range of Fermi coordinates. The results of such an investigation of the relative motion of free particles would naturally depend on the choice of the reference particle. For instance, a separate investigation would be necessary if the reference trajectory is an orbit in the equatorial plane of the Kerr black hole. In fact, a judicious choice of the reference particle would be essential to explain a particular observational result. Moreover the relative distance in our approach should be less than $\\mathcal{R}$.  We find that ultrarelativistic flows in the gravitational field of collapsed  configurations can exhibit purely non-Newtonian behavior due to  general relativistic tidal effects. These include the phenomenon of the tidal  deceleration along the axis of rotation and the phenomenon of tidal acceleration perpendicular to the axis of rotation, which follow from the integration of the generalized Jacobi equation \\cite{1,2}. We study this equation and its extension in this paper and compare our results with the corresponding nonrelativistic tidal effects. In the  following we use units such that $c=1$. The signature of the spacetime metric is assumed to be $+2$.  Imagine a geodesic worldline in the gravitational field of an external source. Let $u^\\mu=dx^\\mu /d\\tau$ be the four-velocity vector of this {\\it reference} geodesic and $\\tau$ be its proper time. A neighboring geodesic with an arbitrary velocity relative to the reference geodesic is connected to it by the deviation vector $\\xi ^\\mu (\\tau )$ such that $\\xi^\\mu u_\\mu=0$. The generalized Jacobi equation \\cite{1} is given by the geodesic deviation equation to first order in $\\xi^\\mu$, namely, \\begin{eqnarray}\\label{eq1} \\nonumber &&\\frac{D^2\\xi^\\mu}{D\\tau^2}+R^\\mu _{\\;\\; \\rho \\nu \\sigma} u^\\rho \\xi^\\nu u^\\sigma \\\\ \\nonumber &&{} +\\Big(u^\\mu +\\frac{D\\xi^\\mu}{D\\tau }\\Big) \\Big( 2R_{\\zeta \\rho \\nu \\sigma}u^\\zeta \\frac{D\\xi^\\rho}{D\\tau }\\xi ^\\nu u^\\sigma +\\frac{2}{3} R_{\\zeta \\rho \\nu \\sigma}u^\\zeta \\frac{D\\xi ^\\rho}{D\\tau}\\xi^\\nu \\frac{D\\xi^\\sigma}{D\\tau} \\Big)\\\\ &&{}+2R^\\mu _{\\;\\; \\rho \\nu \\sigma} \\frac{D\\xi^\\rho}{D\\tau }\\xi^\\nu  u^\\sigma +\\frac{2}{3}R^\\mu_{\\;\\; \\rho \\nu \\sigma}\\frac{D\\xi^\\rho}{D\\tau}\\xi^\\nu \\frac{D\\xi^\\sigma}{D\\tau}=0. \\end{eqnarray} Here $D\\xi^\\mu /D\\tau =\\xi^\\mu_{\\;\\; ;\\nu}u^\\nu$ is the covariant derivative of $\\xi^\\mu$ along the reference worldline. Let $\\lambda^\\mu _{\\;\\;(\\alpha)}$ be an orthonormal tetrad frame that is parallel propagated along the reference trajectory such that $\\lambda^\\mu_{\\;\\;(0)}=u^\\mu$. Then, $\\xi^\\mu =X^i\\lambda^\\mu_{\\;\\; (i)}$ and equation \\eqref{eq1} may be written in terms of Fermi coordinates $(T,{\\bf X})$, where $T=\\tau$ along the reference geodesic, as \\cite{1,2} \\begin{eqnarray}\\label{eq2} \\nonumber &&\\frac{d^2X^i}{dT^2}+\\fr_{0i0j}X^j+2\\,\\fr_{ikj0}V^kX^j\\\\ && +(2\\,\\fr_{0kj0}V^iV^k+\\frac{2}{3}\\,\\fr_{ikj\\ell}V^kV^\\ell +\\frac{2}{3}\\,\\fr_{0kj\\ell}V^iV^kV^\\ell )X^j=0. \\end{eqnarray} Here $V^i=dX^i/dT$, $|\\mathbf{V}|< 1$ along the reference geodesic and \\begin{equation}\\label{eq3} \\fr_{\\alpha \\beta \\gamma \\delta }(T)=R_{\\mu \\nu \\rho \\sigma }\\lambda^\\mu_{\\;\\; (\\alpha )}\\lambda^\\nu _{\\;\\;(\\beta)} \\lambda^\\rho _{\\;\\; (\\gamma )}\\lambda^\\sigma _{\\;\\;(\\delta )} \\end{equation} is the projection of the Riemann curvature tensor on the orthonormal tetrad of the reference observer along its worldline. If the relative velocity is much smaller than the speed of light, $|\\mathbf{ V}|\\ll 1$, the velocity-dependent terms in \\eqref{eq1} and \\eqref{eq2} may be neglected and the generalized Jacobi equation reduces to the standard Jacobi equation that expresses the linear evolution of the deviation between two neighboring geodesics that have negligible relative motion. This relative motion generally evolves slowly on the scale of the radius of curvature $\\mathcal{R}$, which is defined such that $\\mathcal{R}^{-2}$ is the supremum of $|\\fr_{\\alpha \\beta \\gamma \\delta} |$; therefore, $|{\\bf X}|\\ll \\mathcal{R}$ in \\eqref{eq2},  since higher-order terms in the deviation equation have been neglected. In the generalized Jacobi equation, the relative velocity could approach the speed of light; therefore, the relative motion could evolve rapidly thus further restricting the temporal domain of validity of equation \\eqref{eq2}.  Equation~\\eqref{eq2} can be derived in a straightforward manner using the reduced geodesic equation in Fermi coordinates. That is, starting from the geodesic equation for a free particle in arbitrary coordinates $\\hat x^\\mu=(\\hat t, \\hat x^i)$, \\begin{equation}\\label{mu:4} \\frac{d^2\\hat{x}^\\mu}{ds^2} +\\hat{\\Gamma}^\\mu_{\\alpha\\beta} \\frac{d\\hat{x}^\\alpha}{ds}\\frac{d\\hat{x}^\\beta}{ds}=0, \\end{equation} where $ds^2=\\hat{g}_{\\mu\\nu}d\\hat x^\\mu d\\hat x^\\nu$, one can simply derive the reduced geodesic equation~\\cite{1} \\begin{equation} \\frac{d^2\\hat{x}^i}{d\\hat{t}^2} -(\\hat\\Gamma^0_{\\alpha\\beta}\\frac{d\\hat{x}^\\alpha}{d\\hat{t}}\\frac{d\\hat{x}^\\beta}{d\\hat{t}})\\frac{d\\hat{x}^i}{d\\hat{t}} +\\hat\\Gamma^i_{\\alpha\\beta}\\frac{d\\hat{x}^\\alpha}{d\\hat{t}}\\frac{d\\hat{x}^\\beta}{d\\hat{t}}=0. \\end{equation} Specializing this equation now to the Fermi coordinates established along the reference trajectory, i.e. $(\\hat t, \\hat x^i)\\mapsto (T,X^i)$, and using \\begin{eqnarray} \\label{metrict} \\nonumber g_{00}&=&-1-\\fr_{0i0j}(T)\\,X^iX^j+\\cdots,\\\\ \\nonumber g_{0i}&=&-\\frac{2}{3}\\,\\fr_{0jik}(T)\\, X^jX^k+\\cdots,\\\\ g_{ij}&=&\\delta_{ij}-\\frac{1}{3}\\,\\fr_{ikj\\ell}(T)\\, X^kX^\\ell+\\cdots, \\end{eqnarray} one arrives at the tidal equation for the \\emph{relative} motion, since in these coordinates the reference particle remains fixed at $\\mathbf{X}=0$. Restricting our attention to the terms given explicitly in~\\eqref{metrict} and thus neglecting higher-order terms of the metric in the spatial Fermi coordinates, the tidal equation reduces to the generalized Jacobi equation~\\eqref{eq2}. We study the main physical consequences of this equation in the field of a Kerr black hole (see also~\\ref{appen:A}). Higher-order tidal accelerations are discussed in~\\ref{appen:B}.  Tidal deceleration along the rotation axis is discussed in the next section. The non-Newtonian character of our results is emphasized in section~\\ref{s3}. The analysis of the generalized Jacobi equation is extended to the three-dimensional case in section~\\ref{s5}. Tidal acceleration normal to the rotation axis is discussed in section~\\ref{s6}. Finally, section~\\ref{s7} contains a brief discussion of our results. ",
        "conclusions": "\\label{s7} In this paper, we are interested in some of the observable consequences of general relativity involving ultrarelativistic flows near a gravitationally collapsed configuration. The measurement problem in general relativity is rather subtle. Suppose, for instance, that in the background of Kerr spacetime one succeeds in solving certain covariant equations of magnetohydrodynamics (MHD) in Boyer-Lindquist coordinates. Though possibly interesting, this result per se may have little to do with understanding the \\emph{observable} consequences of general relativistic MHD in this case. As a first step, we therefore concentrate on an \\emph{invariant} characterization of the dynamics of free particles in the exterior Kerr spacetime; in particular, we study \\emph{relative} motion in the simplest possible situation, namely, the motion of free particles relative to a certain fiducial class of particles moving freely on escape trajectories along the rotation axis of the Kerr source. The  relative motion of two nearby free particles is then determined by the generalized Jacobi equation.  The generalized Jacobi equation has been known for three decades~\\cite{5}, but its physical consequences for relativistic relative motion  have only recently received attention~\\cite{1,2}. An important issue regarding the generalized Jacobi equation is that for ultrarelativistic relative motion, the duration of validity of this equation is rather short. However, we have shown that even within this limit, for ultrarelativistic outward motion along the rotation axis of a Kerr black hole, there is a remarkably strong initial tidal deceleration relative to the ambient medium regardless of the value of $\\Gamma _0\\gg 1$. To go beyond this initial deceleration, let us note that a significant feature  of our basic equations~\\eqref{eq4}--\\eqref{eq6} is that they do not  explicitly depend on the reference trajectory corresponding to a  ``fixed\" marker in the ambient medium. One can therefore imagine a  number of such markers along the path of the jet such that from one marker to the next, the integration of equation \\eqref{eq8} describes the force-free jet deceleration to a good approximation. In this way,  one may approach the terminal speed of the jet by referring the jet motion to different markers. So long as the jet speed relative to any nearby ambient marker is significantly greater than $1/\\sqrt{2}$, it tends to decrease toward this terminal speed.  It is possible to view the terminal speed in another way. Along a radial direction far from a massive source, its tidal influence on two neighboring test particles is a tidal {\\it acceleration} if the relative speed of the particles is significantly below $1/\\sqrt{2}$. However, this influence turns to a tidal {\\it deceleration} if the relative speed is significantly above $1/\\sqrt{2}$.  We have neglected plasma effects in this paper; in fact, our treatment is valid for a force-free plasma. If the plasma is not force-free, the inclusion of its effects would enormously complicate matters \\cite{2,8}. In any case, tidal deceleration is most effective very close to the central source, where the jet originates. Moreover, it is expected to be more significant for microquasars, since tidal effects near a black hole generally decrease as the inverse square of the mass of the source. Thus the results of this paper may well play a significant role in the dynamics of jets, especially in microquasars~\\cite{4}. More generally, this work may be of interest in connection with the astrophysics of energetic particles such as the ultrahigh energy cosmic rays.  Let us now turn to the motion of a particle normal to the rotation axis of the black hole. Relative to the reference particle, the particle accelerates if its initial speed exceeds the critical speed $1/\\sqrt{2}$. For ultrarelativistic relative motion, this tidal acceleration can lead to ultrahigh energy particles according to the generalized Jacobi equation; more generally, particle acceleration occurs outside a cone of angle $\\theta$, where the polar angle $\\theta$ is measured from the $Z$-axis and $\\tan \\theta \\approx \\sqrt{2}$. As the particle accelerates beyond the Fermi system, one can in principle use other overlapping coordinate systems to describe the motion over an extended period of time. On the other hand, the tidal effects of the black hole diminish rapidly with increasing distance and hence the influence of the black hole can be neglected once the particle has moved sufficiently far away.  The rotation of the source leads to a preferred radial direction, i.e. the axis of rotation, but is otherwise of little importance for the tidal dynamics presented in this paper. This remark applies equally well to the exact nature of the source, except that our results become physically significant near gravitationally collapsed objects, where tidal accelerations can be substantial.  The nature of the accretion mechanisms that create ultrarelativistic particles near black holes is beyond the scope of our investigation. The analysis presented in this paper based on the generalized Jacobi equation indicates that once such particles are created near the poles of the black hole, those propagating mainly along the rotation axis of the black hole decelerate with respect to the ambient medium, while those propagating mainly normal to this axis can accelerate to almost the speed of light.  The loss of kinetic energy for the motion of an initially ultrarelativistic particle along the rotation axis and the gain in kinetic energy for motion normal to this axis may be attributed, in analogy with Newtonian gravitation, to the change in the gravitational potential energy, which is the tidal energy in the case under consideration in this paper. The underlying assumption here is that the perturbation of the background geometry caused by the motion of the test particle may be neglected. This breaks down, however, when the particle is accelerated to almost the speed of light; therefore, the back-reaction on the motion is expected to moderate this singularity leading to an ultrahigh energy particle. More generally, the gravitational and electromagnetic (in case of charged particles) radiations emitted by the decelerating/accelerating particles should be taken into account as well. The tidal acceleration mechanism may provide the key to the explanation of recent \\emph{Chandra} X-ray observations of the Crab Nebula.  It is well known that primary ultrahigh energy cosmic ray protons with energies above $10^{20}$ eV from extragalactic sources are expected to interact with the cosmic microwave background photons resulting in photopion production and pair creation~\\cite{13,14}. Therefore, it may be reasonable to assume that the most energetic particles that reach the solar system originate within our galaxy (see~\\cite{15} for a recent review of cosmic ray astrophysics). In view of the foregoing results, it would be interesting to determine whether the directional distribution of the ultrahigh energy cosmic rays is correlated with the known microquasars, taking due account of the directionality of the acceleration zones normal to the jet directions.  \\appendix \\newcounter{saveeqn}% \\newcommand{\\alpheqn}{\\setcounter{saveeqn}{\\value{equation}}% \\stepcounter{saveeqn}\\setcounter{equation}{0}% \\renewcommand{\\theequation} {\\mbox{\\Alph{section}\\arabic{equation}}}}% \\alpheqn \\renewcommand{\\thesection}{Appendix \\Alph{section}}"
    },
    "0404/astro-ph0404493_arXiv.txt": {
        "abstract": "{ We have studied an ultra-luminous X-ray source (ULX) in the dwarf galaxy NGC\\,5408 with a series of {\\it XMM-Newton} observations, between 2001 July and 2003 January. We find that its X-ray spectrum is best fitted with a power law of photon index $\\Gamma \\approx 2.6$--$2.9$ and a thermal component with blackbody temperature $kT_{\\rm bb} \\approx 0.12$--$0.14$ keV. These spectral features, and the inferred luminosity $\\approx 10^{40}$ erg s$^{-1}$ in the $0.3$--$12$ keV band, are typical of bright ULXs in nearby dwarf galaxies. The blackbody plus power-law model is a significantly better fit than either a simple power law or a broken power law (although the latter model is also acceptable at some epochs). Doppler-boosted emission from a relativistic jet is not required, although we cannot rule out this scenario. Our preliminary timing analysis shows flaring behaviour which we interpret as variability in the power-law component, on timescales of $\\sim 10^2$ s. The hard component is suppressed during the dips, while the soft thermal component is consistent with being constant. The power density spectrum is flat at low frequencies, has a break at $\\nu_{\\rm b} \\approx 2.5$ mHz, and has a slope $\\approx -1$ at higher frequencies. A comparison with the power spectra of Cyg X-1 and of a sample of other BH candidates and AGN suggests a mass of $\\sim 10^2 M_{\\odot}$. It is also possible that the BH is at the upper end of the stellar-mass class ($M \\sim 50 M_{\\odot}$), in a phase of moderately super-Eddington accretion. The formation of such a massive BH via normal stellar evolution may have been favoured by the very metal-poor environment of NGC\\,5408. ",
        "introduction": "Ultra-luminous X-ray sources (ULXs) are point-like sources, not including galactic nuclei and young supernova remnants, with apparent luminositites higher than the Eddington limit for a stellar-mass accreting black hole (BH), ie, with $L_{X}\\ga 3 \\times 10^{39}$\\,erg s$^{-1}$ (Fabbiano 1992; Colbert \\& Mushotzky 1999; Roberts \\& Warwick 2000). The true nature of these objects is still open to debate. The fundamental issue is whether the emission is isotropic or beamed along our line-of-sight; in the latter case, the total luminosity would of course be lower than the isotropic value. A possible scenario for geometrical beaming is a super-Eddington mass inflow during phases of thermal-timescale mass transfer (King 2002). Relativistic beaming, associated with Doppler boosting, would instead be the effect of our looking into the jet of a microquasar (K\\\"{o}rding et al.~2002). Alternatively, if the emission is isotropic and the Eddington limit is not violated, ULXs must be fuelled by accretion onto intermediate-mass BHs, with masses $\\sim 10^{2}\\,M_{\\odot}$. Possible mechanisms of formation of such massive remnants include the collapse of massive population {\\rm III} stars (Madau \\& Rees 2001), or runaway stellar mergers in young clusters (Portegies Zwart et al.~2004).   A ULX in the starburst dwarf irregular galaxy NGC\\,5408 could help discriminate between these two alternatives. Radio observations of this galaxy (Stevens et al.~2002) showed four main emission regions, mostly coincident with massive young star clusters and H$\\alpha$ emission. The X-ray emission is dominated by a single point source, with a luminosity $>10^{40}$ erg s$^{-1}$ (Fabian \\& Ward 1993; Kaaret et al.~2003), located outside the H{\\footnotesize II} regions, but apparently coincident with a weak, steep-spectrum radio source. On the basis of the X-ray, radio, and optical fluxes, Kaaret et al.~(2003) concluded that the most likely explanation for the ULX was beamed relativistic jet emission from a stellar-mass BH. However, they could not rule out an intermediate-mass BH model. In this paper we present the first results of our {\\it XMM-Newton} study of the ULX in NGC\\,5408, providing new insights into the nature of this source. ",
        "conclusions": "Our spectral analysis, based on five {\\it XMM-Newton}/EPIC observations between 2001 July and 2003 January, shows a thermal component with $kT_{\\rm bb} \\approx 0.12$--$0.14$ keV, significantly detected at all epochs, in addition to a soft power-law ($\\Gamma \\approx 2.6$--$2.9$). A broken power-law model, as suggested in Kaaret et al.~(2003), is acceptable at some epochs, but is not consistent with the full series of observations, and is generally worse than a blackbody plus power-law model or a Comptonised blackbody model.  Ruling out a simple broken power-law model does not rule out the possibility that the X-ray emission is due to inverse-Compton emission from a jet, enhanced by relativistic Doppler boosting (Kaaret et al.~2003). It was suggested in Georganopoulos (2002), and Georganopoulos \\& Kazanas (2003), that a more realistic jet spectrum would have a curvature around the break energy, thus mimicking a soft thermal component. Spectral analysis over a larger energy range is required to test this possibility.  In any case, the X-ray spectrum is consistent with a more massive accreting BH in a high/soft or very high state. Given its inferred isotropic luminosity of $\\approx 10^{40}$ erg s$^{-1}$ in the $0.3$--$12$ keV band, a mass $\\ga 100 M_{\\odot}$ is required to satisfy the Eddington limit. Thus, the ULX in NGC\\,5408 appears very similar to a group of ``canonical'' ULXs in nearby dwarf and spiral galaxies, with a low-temperature thermal component at $kT_{\\rm bb} \\sim 0.1$ keV and emitted luminosity $\\sim (1$--$2) \\times 10^{40}$ erg s$^{-1}$. In all these systems, if the soft thermal component is interpreted as the emission from the inner part of a standard accretion disk (Shakura \\& Sunyaev 1973), its temperature would imply masses $> 10^5 M_{\\odot}$, unphysically high for a non-nuclear BH. However, it was also suggested (King 2003) that these ULXs may be accreting BHs during a super-Eddington accretion phase: they could have a luminosity close or even a factor of a few above the Eddington limit, accompanied by a radiatively-driven, Compton-thick outflow from the accretion disk. In this scenario, the soft thermal component would come from the photosphere of the outflow. Hence, the BH mass would not need to be higher than $\\sim 20$--$50 M_{\\odot}$, still consistent with a stellar origin.  We note that this ULX is located in a very metal-poor dwarf galaxy ($Z \\approx 0.07 Z_{\\odot}$). An association between ULXs and metal-poor environments was suggested by Pakull \\& Mirioni (2002). Low metal abundance implies a reduced mass-loss rates in the radiatively-driven wind from the O-star progenitor ($\\dot{M}_{\\rm w} \\sim Z^{0.85}$: Vink et al.~2001; see also Bouret et al.~2003); this leads to a more massive stellar core, which may then collapse into a more massive BH, via normal stellar evolution. Hence, low abundances may explain the formation of isolated BHs with masses up to $\\approx 50 M_{\\odot}$ (ie, up to $\\approx$ half of the mass of the progenitor star).   We found a moderate long-term spectral variability between the various epochs of the {\\it XMM-Newton} observations, with a ratio between the maximum and minimum fluxes of $\\approx 1.4$. This is consistent with the long-term behaviour of the source from previous {\\it Einstein}, {\\it ROSAT}, {\\it ASCA} and {\\it Chandra} observations (Kaaret et al.~2003). Our preliminary timing analysis shows rapid variability and flaring-like behaviour, particularly in the 2003 Jan observation. The variability is larger in the hard band ($E > 1.5$ keV), where the flux is consistent with zero during the dips.  The {\\tt {bb+po}} spectral model offers a clean, simple interpretation of the flaring behaviour. In this scenario, we have interpreted the lightcurves as the combination of a rapidly variable power-law component, plus an additional, non-variable thermal component below 1.5 keV, consistent with our spectral analysis. The power-law flux appears to be suppressed during the dips, when the X-ray spectrum becomes softer; this is supported by flux-dependent spectral analysis. The power-law X-ray emission in accreting BH binaries is generally explained as inverse-Compton scattering of soft disk photons in a hot corona or Compton cloud, located either above the inner part of the accretion disk, or as a quasi-spherical region inside the truncation radius of the disk.  A flaring behaviour with soft dips or transitions has also been observed in Galactic microquasars such as GRS 1915$+$105 (Naik et al.~2001) and, in one case, XTE J1550$-$564 (Rodriguez et al.~2003). It was interpreted (Vadawale et al.~2003) as evidence that the matter responsible for the Comptonised component is recurrently ejected from the inner regions; the ejected matter may be responsible for the optically-thin synchrotron radio emission in those systems. A more detailed discussion of this hypothesis, and of whether it may explain the steep-spectrum radio emission detected by the Australia Telescope Compact Array at the position of the ULX (Kaaret et al.~2003; Stevens et al.~2002), will be addressed in further work.  The power density spectrum for this ULX is flat below $\\approx 2.5$ mHz, and has a slope of $\\approx -1$ at higher frequencies. This is similar to what is found in ``canonical'' Galactic BH candidates and AGN: it is believed that the break frequency is inversely proportional to the mass of the accreting BH (eg, Belloni \\& Hasinger 1990; Nowak et al.~1999; Uttley et al.~2002; Markowitz et al.~2003; Cropper et al.~2004). For example, the break in the power density spectrum of Cyg X-1 occurs at $\\nu \\sim 0.4$--$0.04$ Hz in different accretion states (Nowak et al.~1999). This suggests that the mass of the BH powering the ULX in NGC\\,5408 is $\\sim 10^2 M_{\\odot}$. This is also in good agreement with the inferred luminosity, if we require that it does not exceed the Eddington limit. Such a high mass would rule out (cf.~Fig.~3 of Kaaret et al.~2004) an association of this ULX with the young star clusters in the starburst region of NGC\\,5408, located $\\approx 300$ pc to the north-west. However, we are aware that the relation between break frequency and BH mass is very uncertain, and the power-density spectra are often more complicated than a simple broken power law (see, eg, the case of GRS 1915$+$105: Morgan et al.~1997).  In conclusion, we argue that both the spectral and timing results from our {\\it XMM-Newton} study are consistent with the possibility that this ULX is at the upper end of the mass range for BHs of stellar origin, a few times more massive than ``standard'' BH candidates such as Cyg X-1. In that case, its luminosity would be very close or, more likely, slightly higher than its Eddington limit, and a strong outflow would be expected from its accretion disk. Doppler boosting in a relativistic jet is not required to explain the observed properties, although we cannot rule it out."
    },
    "0404/astro-ph0404391.txt": {
        "abstract": "We present new measurements of the Sunyaev-Zel'dovich (SZ) effect from clusters of galaxies using the Sunyaev-Zel'dovich Infrared Experiment (SuZIE~II). We combine these new measurements with previous cluster observations with the SuZIE instrument to form a sample of 15 clusters of galaxies. For this sample we calculate the central Comptonization, $y_0$, and the integrated SZ flux decrement, $S$, for each of our clusters. We find that the integrated SZ flux is a more robust observable derived from our measurements than the central Comptonization due to inadequacies in the spatial modelling of the intra-cluster gas with a standard Beta model. This is highlighted by comparing our central Comptonization results with values calculated from measurements using the BIMA and OVRO interferometers. On average, the SuZIE calculated central Comptonizations are $\\sim60\\%$ higher in the cooling flow clusters than the interferometric values, compared to only $\\sim12\\%$ higher in the non-cooling flow clusters. We believe this discrepancy to be in large part due to the spatial modelling of the intra-cluster gas. From our cluster sample we construct $y_0$--$T$ and $S$--$T$ scaling relations. The $y_0$--$T$ scaling relation is inconsistent with what we would expect for self-similar clusters; however this result is questionable because of the large systematic uncertainty in $y_0$. The $S$--$T$ scaling relation has a slope and redshift evolution consistent with what we expect for self-similar clusters with a characteristic density that scales with the mean density of the universe.  We rule out zero redshift evolution of the $S$--$T$ relation at $\\sim90$\\% confidence. ",
        "introduction": "\\label{intro}  Clusters of galaxies are the largest gravitationally bound objects in the Universe, and formed at relatively early times over a critical redshift range ($0<z \\lesssim 3$) during which the dark energy came to dominate the total energy density of the Universe. A measurement of the evolution of the cluster number density with redshift is sensitive to various cosmological parameters, including $\\sigma_8$, $\\Omega_M$, $\\Omega_{\\Lambda}$, and the dark energy equation of state \\citep{wang98,holder01}.  A direct measurement of the cluster number density can be made through a survey utilizing the Sunyaev-Zel'dovich (SZ) effect.  The SZ effect is particularly well-suited for cluster surveys because an SZ survey will detect every cluster above a mass limit that is independent of redshift \\citep[see][for example]{carl02}. Active and planned SZ surveys should result in the discovery of tens of thousands of clusters over the next several years, \\citep[see][for a review]{white03}.  The cosmological constraints from any SZ survey may ultimately be limited by how closely clusters behave as standard candles. \\citet{haiman01} showed that future SZ surveys are likely to be limited by systematic uncertainties due to the assumption that clusters are virialized objects whose density scales with the mean background density. Observations of relatively nearby clusters ($z \\lesssim 0.1$) in X-rays have shown that clusters are at least remarkably regular objects whose observable properties seem to obey well-behaved scaling relations. These scaling relations include the mass-temperature \\citep[e.g.,][]{fin01}, size-temperature \\citep[e.g.,][]{mohr00}, and luminosity-temperature \\citep[e.g.,][]{mark98} scaling relations. \\citet{verde02} have argued that an integrated SZ flux versus X-ray temperature scaling relation should have an exceptionally small scatter, compared to other cluster scaling relations, and should be especially useful in testing possible deviations from virialization.  Investigations of SZ scaling relations have been limited so far due to a scarcity of measurements. \\citet{cooray99} and \\citet{mccarthyb} have compiled SZ measurements drawn from the literature and constructed SZ scaling relations; however these studies suffered from several draw-backs. Firstly, both papers drew upon measurements from several different instruments. While no systematic differences across instruments was known of at the time, we believe this paper offers the first evidence that a significant systematic discrepancy does exist. Secondly, both papers concentrated on scaling relations that use the central decrement of the cluster. The calculated central decrement often relies on an assumed spatial distribution of the intra-cluster (IC) gas, whose density is still best constrained by higher resolution X-ray data. The traditional parameterization of the IC gas distribution with the single Beta model \\citep{betaref1,betaref2} causes a large systematic uncertainty in the central decrement calculated from SZ measurements. In this paper, we address both these issues by constructing scaling relations using SZ measurements from only one instrument, and instead focusing on the integrated SZ flux scaling relation.  %could mention something about X-ray data quality improving.  In this paper we describe new measurements of the SZ effect towards eight galaxy clusters made with the Sunyaev-Zel'dovich Infrared Experiment (SuZIE).  The SuZIE~II receiver makes simultaneous measurements of the SZ effect in three frequency bands, centered at 145\\,GHz, 221\\,GHz, and 355\\,GHz, spanning the null of the SZ thermal effect.  In combination with previous measurements from the SuZIE instrument, we form a sample of 15 clusters for which we determine the central Comptonization, $y_0$, and integrated SZ flux, $S$. We then use this sample to construct scaling relations between $y_0$ and $S$ and the X-ray temperature and compare these results to theoretical predictions. The layout of this paper is as follows: in \\S\\ref{sec:obs} and~\\S\\ref{sec:szanalysis} we describe the observations and multi-frequency data analysis. In \\S\\ref{sec:145analysis} we describe the analysis to determine the SZ flux. In \\S\\ref{sec:y0compare} we compare our SZ measurements to other experiments, particularly from the OVRO and BIMA interferometers. Finally, in \\S\\ref{sec:scaling} we construct SZ scaling relations and compare these measurements with theoretical predictions. ",
        "conclusions": "We report new measurements of the SZ effect from clusters of galaxies in three frequency bands. We use the multi-frequency measurements to measure the SZ spectrum and the central Comptonization in each cluster.  We combine these new measurements with previous cluster observations to construct a sample of 15 clusters of galaxies detected with the SuZIE experiment. For the entire set of clusters we use the 145 GHz frequency band to calculate a central Comptonization, $y_0$, and an integrated SZ flux, $S(r_{2500})$.  We find that the calculated central Comptonization is much more sensitive to the assumed spatial model for the intra-cluster gas than the calculated integrated SZ flux.  The calculated central decrement depends significantly on the assumed spatial distribution of the intra-cluster (IC) gas. For the case of A1835 the calculated central Comptonization can vary by a factor of two depending on which of two different published IC gas models is assumed.  This result is not surprising considering the fact that SuZIE~II does not significantly resolve any of the observed clusters. This effect causes the calculated central Comptonization to have a particularly large systematic uncertainty in cooling flow clusters because of their large cooling core which makes the standard Beta model an inadequate fit to the spatial distribution of the gas. However, our measurements of the integrated SZ flux, $S(r_{2500})$, negligibly depend on the assumed spatial distribution of the IC gas because $r_{2500}$ is well-matched to our beam-size for most of the observed clusters.  %This effect is believed to be more pronounced in cooling flow %clusters because of their large cooling core which makes the %standard Beta model an inadequate fit to the spatial distribution %of the gas.  Ten of the clusters in our sample overlap with published measurements from the BIMA and OVRO interferometers.  For these clusters, we compare the calculated central Comptonization from BIMA and OVRO to those from SuZIE and find that the SuZIE calculated central Comptonizations are generally higher, significantly so in the cooling flow clusters.  If we compare the SuZIE 145 GHz results to the BIMA and OVRO results, SuZIE measures a central Comptonization $\\sim12\\%$ higher in the non-cooling flow clusters, and $\\sim60\\%$ higher in the cooling flow clusters.  We attribute this difference to the large systematic uncertainty in the calculated central Comptonization from the assumed intra-cluster gas model which, as expected, is more pronounced in our cooling flow sub-sample.  We use the central Comptonization and integrated SZ flux results from the SuZIE 145 GHz data to construct SZ scaling relations with the X-ray temperature, $T_X$. We construct a $y_0$--$T_X$ scaling relation and find a slope significantly different than what is expected for self-similar clusters.  However, we believe that this result is questionable because of the large systematic uncertainty in the central Comptonization.  This conclusion is supported by the significantly discrepant scaling relations derived for the cooling flow and non-cooling flow sub-samples. For the $S(r_{2500})$--$T_X$ scaling relation we find a slope which is consistent with the expectation for self-similar clusters. In constructing this relation, we find that using X-ray temperatures which do not account for the presence of the cooling flow significantly biases the best-fit relation. We detect a redshift evolution of the $S(r_{2500})$--$T_X$ scaling relation consistent with standard cluster formation theory for which the density of the cluster scales with the mean density of the universe.  If we assume that the X-ray temperature is a good indicator of the mass of the cluster, as suggested from X-ray measurements, our results imply that the the integrated SZ flux will be a good indicator of the cluster mass as well, a promising result for future SZ cluster surveys.  We thank Jim Bartlett, Gilbert Holder, and Phil Marshall for many useful discussions.  The SuZIE program is supported by a National Science Foundation grant AST-9970797, NASA grant NAG5-12973, and a Stanford Terman Fellowship awarded to SEC. This work was partially carried out at the Infrared Processing and Analysis Center and the Jet Propulsion Laboratory of the California Institute of Technology, under a contract with the National Aeronautics and Space Administration. We would like to thank the CSO staff for their assistance with SuZIE observations. The CSO is operated by the Caltech Submillimeter Observatory under contract from the National Science Foundation.  \\newpage"
    },
    "0404/astro-ph0404547_arXiv.txt": {
        "abstract": "We propose a new model for the origin and evolution of the Local Group of Galaxies (LGG) which naturally explains the formation of the Magellanic Clouds and their large orbital angular momenta around the Galaxy. The basic idea is that an off-center hydrodynamical collision occurred some 10 Gyr ago between the primordial gas-rich Andromeda galaxy and the similar Galaxy, and compressed the halo gas to form the LGG dwarf galaxies including the Magellanic Clouds. In this model, new-born dwarf galaxies can be expected to locate near the orbital plane of these two massive galaxies. In order to see the reality of this model, we reexamine the two-dimensional sky distribution of the LGG members and the Magellanic Stream, we confirm an earlier and widely-discussed idea that they align along two similar great circles, each with an angular width of $\\sim $30$^{\\circ}$, and the planes of these circles are approximately normal to the line joining the present position of the sun and the Galactic center. Further we make a three-dimensional distribution map of these objects, and observe it from various directions. A well-defined plane of finite thickness is found, within which most of the member galaxies are confined, supporting the existence of the above circles on the sky. Thus we could determine the orbital elements of M31 relative to the Galaxy through reproducing the well-studied dynamics of the LMC and the SMC around the Galaxy. The expected proper motion of M31 is $(\\mu _{l},\\ \\mu_{b})=(38\\ \\mu{\\rm as~yr}^{-1}, -49\\ \\mu{\\rm as~yr}^{-1})$. Probable orbital motions of the other dwarfs are also determined, and the corresponding proper motion for each object is given to compare with observations in near future. ",
        "introduction": "In spite of extensive observational/theoretical studies of the Magellanic Clouds, their origins remain unsolved for a long time. Did they initially form as satellite galaxies of the Galaxy, or did they fall into the Galaxy in the course of dynamical evolution of the LGG \\citep{byr94}. A more realistic hint to this problem can be presented by studies of the origin of the Magellanic Stream (MS). Tidal models have been successfully introduced to the dynamics of the Galaxy-LMC-SMC system for reproducing the geometrical and dynamical structure of the MS \\citep{mur80,lin82,gar94,gar96}. In such models, orbits of the Large Magellanic Cloud (LMC) and the Small Magellanic Cloud (SMC) can be traced back in time over the entire past period of $\\sim 10$ Gyr: The orbital plane is approximately perpendicular to the line joining the present position of the sun and the Galactic center, and they are viewed to move counterclockwise along a nearly great circle centered on $(l, b)=(0^{\\circ},0^{\\circ})$ or the Galactic center.  However, the tidal model in which the Magellanic Clouds formed in the neighborhood of the Galaxy cannot explain large orbital angular momenta of the Magellanic Clouds around the Galaxy (see, for example, \\citet{fic94}). According to the theory of the Magellanic Stream, the apo-Galactic distance of the LMC orbit is $R_\\mathrm{LMC} \\sim 200$ kpc and its transverse velocity is $V_\\mathrm{LMC} \\sim $100 km s$^{-1}$ some 10 Gyr ago, implying that if the LMC mass is $M_\\mathrm{LMC} \\sim 2 \\times 10^{10}\\MO$, the orbital angular momentum of the Magellanic Clouds $L_\\mathrm{LMC}$ at that time is given by \\begin{eqnarray} L_\\mathrm{LMC} & \\sim & M_\\mathrm{LMC}R_\\mathrm{LMC}V_\\mathrm{LMC} \\nonumber \\\\ & \\sim & 2 \\times 10^{10}\\MO \\times 200{\\rm ~kpc~}\\times 100 {\\rm ~km~s}^{-1} \\nonumber \\\\ & = & 4\\times 10^{14} \\MO {\\rm ~kpc~km~s}^{-1}, \\end{eqnarray} while the spin of the Galactic disk $L_\\mathrm{disk}$ in flat rotation with velocity $V_{0}$ is given by \\begin{eqnarray} L_\\mathrm{disk} & \\sim & \\int r V_{0} dm \\nonumber \\\\ & \\sim & \\int_{0}^{R_\\mathrm{disk}} \\int_{0}^{2\\pi}\\int_{-h/2}^{h/2} \\rho_{0}r^{2}V_{0}drd\\phi dz \\nonumber \\\\ & = & \\frac{2\\pi}{3}\\rho_{0}V_{0}R^{3}h, \\end{eqnarray} where $\\rho_{0}$ is the density of the disk, $h$ the thickness, and $R_\\mathrm{disk}$ the outer radius. In this estimate, we assume $\\rho_{0}$ and $h$ to be constant. Since the total mass of the disk $M_\\mathrm{disk}$ is given by $M_\\mathrm{disk}= \\pi R_\\mathrm{disk}^{2} h \\rho_{0}$, the spin of the Galactic disk $L_\\mathrm{disk}$ is estimated as \\begin{eqnarray} L_\\mathrm{disk} & \\sim & \\frac{2}{3}M_\\mathrm{disk}R_\\mathrm{disk}V_{0} \\nonumber \\\\ & \\sim & \\frac{2}{3}\\times2\\times 10^{11} \\MO \\times 15 {\\rm kpc} \\times 220{\\rm ~km~s}^{-1} \\nonumber \\\\ & = & 4\\times 10^{14} \\MO{\\rm ~kpc~km~s}^{-1}, \\end{eqnarray} if we adopt $V_{0}=220$ km s$^{-1}$ and $M_\\mathrm{disk}=2 \\times 10^{11} \\MO$ for $R_\\mathrm{disk}=15$ kpc. We find that the orbital angular momentum of the LMC, $L_\\mathrm{LMC}$, is comparable to the spin of the Galactic disk $L_\\mathrm{disk}$. It should be also born in mind that these angular momenta make a right angle against each other. They cannot be explained in terms of the tidal interaction in the LGG \\citep{got78}. Further it has also been widely recognized that the morphological type of dwarf galaxies near their massive parent galaxy are generally dwarf spherical or spheroidal, but that the Magellanic Clouds are morphologically classified as dwarf irregulars in spite of their close location to the Galaxy and are exceptionally massive \\citep{ein74} and gas-rich.  These facts are difficult to understand only in terms of the gravitational interaction among the three-body system, the Galaxy-LMC-SMC. Then \\citet{fuj99} and \\citet{saw99} conducted numerical simulations of the Magellanic Stream by taking into account the gravitational effect of the nearby massive galaxy M31 which is approaching us with velocity of 120 km s$^{-1}$ if we adopt the standard rotation velocity of the LSR ($V_{0}=220$ km s$^{-1}$) and the solar motion ($v_{\\odot}=16.5$ km s$^{-1}$, $l_{\\odot}=53^{\\circ}$, $b_{\\odot}=25^{\\circ}$). With the orbital elements of M31 assumed by \\citet{kah59}, \\citet{fuj99}, and \\citet{saw99}, it is found that the binary orbits of the LMC and the SMC at their early times are much disturbed compared with those determined by \\citet{mur80} and \\citet{gar94}, but no successful result has been obtained for solving the angular momentum problem even through the four-body dynamics of the Galaxy-M31-LMC-SMC.    The present authors come to take a view point that some questions about the Galaxy-LMC-SMC system must be answered through a more-global model for the LGG in which the Galaxy, the LMC and the SMC should be treated as three objects among more-than-forty members of the LGG. For this reason we begin with in the next section our reexamination of the sky (two-dimensional) and spatial (three-dimensional) distributions of the member galaxies of the LGG, which would give key ideas for the dynamics and the origin of the LGG.   In concluding this section, we note that many studies have tried to determine the past orbits of the LGG members under their various assumptions. They are conducted mostly in virial scheme and, therefore, the results are rather stochastic. Here we should refer to three important works by \\citet{mis85}, \\citet{pee93}, and \\citet{lyn99}. They have pioneered the global LGG dynamics, although many basic differences are present from our ideas and results, particularly concerning the origin of the LGG. ",
        "conclusions": "On the bases of the ring-like distributions of the Local Group of Galaxies and the Magellanic Stream on the plane of sky (seen in the two-dimensional space), and of their coplanar distribution seen also in the three-dimensional space, we have proposed a dynamical model for the origin and evolution of the dwarf members of the LGG: The primordial dwarfs, the primordial LMC and SMC as well, were driven to form in the high density region of gas generated hydrodynamically by the off-center collision between the primordial gas-rich M31 and the primordial gas-rich Galaxy some 10 Gyr ago. Newborn dwarf galaxies are scattered on the orbital plane of these two massive galaxies and they are trapped either to the potential well of the Galaxy or to that of M31, while other some number of dwarfs are left behind from such trappings. The latter objects are observed as if linking the Galaxy and M31 groups.  The orbital motion of M31 (relative to the Galaxy) is determined so that it successfully reproduces the well-studied dynamics and history of the LMC/SMC system. The most important model scenario is that the LMC/SMC are born from the high-density gas of the formation site, together with other dwarf members of the LGG simultaneously. Thus the orbital plane of M31 must be parallel to that of the LMC/SMC, roughly normal to the line joining the present position of the sun and the Galactic center. Both M31 and the LMC/SMC revolve round the Galaxy counterclockwise seen from the sun, on the unclosed elliptical orbit.  Thus we could have an answer to a long-standing question what the origin of the LMC/SMC system is and where their large orbital angular momenta come from. Moreover interestingly, we could identify the birth places of M31 and the LMC/SMC on the sky and in the space, if precise age of the universe is given.  We are unable to demonstrate theoretically the validity of our dynamical model only from the data about the dwarf members: their positions on the sky $(l, b)$, radial velocities and radial distances measured from the Galactic center. Therefore, instead, we have searched for the dwarf members' past orbits that follow the model scenario. Figures \\ref{fig-12} and \\ref{fig-13} can actually show that the dwarf members start from their common formation site, move on the orbital plane (figure \\ref{fig-13}), and arrive at the present positions, reproducing the ring-like distribution in figures \\ref{fig-3} and \\ref{fig-4}.  From the results in Table \\ref{tab-2} we find finally that more than sixty percent of the LGG members (22 out of 35) have $\\bigcirc $, suggesting our dynamical model for the LGG is very realistic.   The dwarf members marked with $\\times$ in Table \\ref{tab-2} do not have their orbits that can be traced back in time to the formation site in our model scenario. They could be objects that plunged into the spherical LGG region from outside, or their orbits have been bent gravitationally by other nearby massive galaxies. We could also consider that these objects are intrinsic dwarfs born directly from the expanding universe and left behind from their merger to M31 and/or to the Galaxy, although we have again a question why such cosmologically early-formed objects tend to gather in the LGG disk.  On the other hand, if we relax our model conditions in such a way that the LGG dwarfs are formed due also to the tidal disturbance in gas around M31 and the Galaxy, then the formation site depicted in figure \\ref{fig-5}(a) would be largely extended on their orbital plane. See a recent HST picture and its caption for dwarf formation in the Stepfan's Quintet \\citep{wea01}. The number of $\\bigcirc $ in Table \\ref{tab-2} will increase, although it becomes difficult to predict the dwarfs' proper motions.  A small mass spiral, M33, has been treated as a test-particle in determining the orbits of the LGG members. It is, of course, not correct in a sense that the mass of M33 is three to four times as large as that of the LMC/SMC system. It is necessary to join this spiral to our four-body problem, but we do not yet have a key phenomenon and/or theory how to include M33 in a new five-body dynamics of the Galaxy-M31-M33-LMC-SMC in the LGG region. A fragmentary chain of hydrogen gas is observed as if bridging M33 and M31 \\citep{bli99}, and this may be usable to determine the M33 orbit. However, we remain open to a conclusion if it is in the same category as the Magellanic Stream. Even if the gravity of M33 is taken into account, however, we believe that our conclusions do not change so significantly, because the total mass of the Galaxy and M31 dominates that of all other members and control them gravitationally.   We did not use the data about the tidal-cut-off radii of the LGG dwarfs, because, as figures \\ref{fig-12} and \\ref{fig-13} show, they could be located close to the Galaxy, or M31 or to both before they become condensed to dwarf galaxies. Even if the pericenter distance of a dwarf's orbit is calculated formally from the tidal-cut-off-radii data, we cannot know what it means in our LGG model.  \\citet{yoo02} found a group of metal-poor star-clusters (counted as seven at the moment) and showed that these clusters align on the plane perpendicular to the line joining the present position of the sun and the Galactic center, in quite a similar way to that discussed about the LGG members in figures \\ref{fig-1} and \\ref{fig-2}. As they suggest, if these metal-poor clusters were born originally in the LMC and have been captured recently to the Galaxy, some of the globular clusters in the Galactic halo are regarded as tightly related to the LGG dynamics that the LMC was driven to form at the off-center collision between M31 and the Galaxy.    We have pursued backward in time M31 and the LMC/SMC in an expanding flat universe. When the present model proves to be realistic and is improved to be more quantitative, we must reexamine various parameters applied so far. For example, a claim may be made that the lifetime of 10.4 Gyr for the LGG model is a little too small and $\\sim 12$ Gyr more suitable. The former age has been obtained by use of the masses of the Galaxy and M31 which are applied uniformly in their series of papers by \\citet{mur80}, \\citet{gar94}, \\citet{fuj99} and \\citet{saw99}. However, if we take a-little-less masses for the Galaxy \\citep{pra03} and M31, or their smaller radii of dark matter halos, it is possible that the latter age is more reasonable. At present, however, we must wait for a moment to conclude which is more realistic and more contribute to understanding the origin and evolution of the LGG.   We have only briefly touched upon some problems such as the tidal effect on the disks of the Galaxy and M31, the dark matter mass-fraction among the LGG dwarfs, and etc. The extensive studies on these problems would contribute to obtaining new hints to improve and develop more realistically our LGG model.  \\bigskip  We would like to thank Professor Y. Kumai for his critical comments and discussions about the present work, and one of us (M.F) deeply thanks Vice-President of Nagoya University K. Yamashita for his hospitality and research support during my stay at his U-Lab., Department of Physics and Astrophysics."
    },
    "0404/astro-ph0404017_arXiv.txt": {
        "abstract": "We present the preliminary results of the INTEGRAL survey of the Large Magellanic Cloud. The observations have been carried out in January 2003 (about 10$^6$ s) and January 2004 (about 4$\\times10^{5}$ s). Here we concentrate on the bright sources LMC X-1, LMC X-2, LMC X-3 located in  our satellite galaxy, and on the serendipitous detections of the Galactic Low Mass X-ray Binary EXO 0748--676 and of the Seyfert 2 galaxy IRAS 04575--7537. ",
        "introduction": "\\begin{figure*}[ht!] \\centerline{\\psfig{figure=jemx2.ps,angle=0,width=14cm}} \\caption{Mosaics of the JEM-X observations of the LMC field performed in January 2003.} \\label{jemx-im} \\end{figure*}   The Large Magellanic Cloud (LMC), with an angular extent of $\\sim$10$^{\\circ}$$\\times$5$^{\\circ}$ well matched to the fields of view of the IBIS \\citep{ibis} and SPI \\citep{spi} instruments, is an ideal target to exploit the multiplexing capabilities of INTEGRAL.  Being the closest galaxy (d$\\sim$54 kpc), the LMC has always been a primary target for all  X-ray and $\\gamma$-ray missions. It contains a few  luminous  X-ray binaries, that have been extensively studied in X-rays,  but for which little information has been obtained above $\\sim$20-30 keV. Besides giving information on the brightest sources, an INTEGRAL survey of the LMC  should allow studies of diffuse emission, and to discover new fainter hard X-ray sources.  INTEGRAL observed the LMC for about 1 Msec in January 2003 and for about 400 ksec in January 2004. Fig.\\ref{jemx-im} shows the images of the LMC field obtained with the JEM-X instrument \\citep{jemx} in the energy ranges 3-5 keV and 5-10 keV.  Here we present the preliminary results on the brightest sources detected in this field with IBIS: the black hole candidate LMC X-1,  the Low Mass X-ray Binary EXO 0748--676, and the AGN IRAS 04575--7537. The two latter sources do not belong to the LMC, but were serendipitously detected in the large IBIS field of view.  As can be seen in Fig.\\ref{jemx-im}, also LMC X-2, LMC X-3 and LMC X-4\\footnote{we will not report on LMC X-4 since the data rights for this source belong to a different group} were clearly detected. LMC X-2 is a Low Mass X-ray Binary containing a neutron star, and the only Z-type source known out of our Galaxy \\citep{smale}. LMC X-3 is a persistent black hole candidate \\citep{cowley} which spends most of the time in a soft spectral state, with a very faint power law tail at high energies. Due to their  soft spectrum and relatively low flux, LMC X-2 and LMC X-3 were not detected in the preliminary analysis of the IBIS data carried out to date. ",
        "conclusions": ""
    },
    "0404/astro-ph0404221_arXiv.txt": {
        "abstract": "Interstellar magnetic fields exist over a broad range of spatial scales, extending from the large Galactic scales ($\\sim 10$~kpc) down to the very small dissipative scales ($\\ll 1$~pc).  In this paper, we use a set of 490 pulsars distributed over roughly one third of the Galactic disk out to a radius $R \\simeq 10$~kpc (assuming $R_\\odot = 8.5$~kpc) and combine their observed rotation and dispersion measures with their estimated distances to derive the spatial energy spectrum of the Galactic interstellar magnetic field over the scale range $0.5 - 15$~kpc. We obtain a nearly flat spectrum, with a 1D power-law index $\\alpha=-0.37\\pm0.10$ for $E_{\\rm B}(k)=C\\, k^{\\alpha}$ and an rms field strength of approximately $6\\, \\mu$G over the relevant scales.  Our study complements the derivation of the magnetic energy spectrum over the scale range $0.03 - 100$~pc by \\citet{ms96b}, showing that the magnetic spectrum becomes flatter at larger scales.  This observational result is discussed in the framework of current theoretical and numerical models. ",
        "introduction": "On the whole, the physical properties of the interstellar gas are better established and understood than those of the interstellar magnetic field. The spatial power spectrum of interstellar HI has been measured using observations of the 21-cm line in several small areas across the sky. Emission measurements suggest a smooth power-law behavior with a 3D spectral index $\\simeq -3$ for spatial scales $\\ga 5$~pc \\citep[e.g.,][]{cd83,gre93,dms+01}, while 21-cm absorption measurements yield a 3D spectral index $\\simeq -2.75$ in the scale range $0.01 - 3$~pc \\citep{ddg00}.  The spatial structure of interstellar free electrons has been probed using observations of Galactic pulsars and extragalactic compact sources. Relying on interstellar scintillation data, \\citet{ars95} showed that the spatial power spectrum of the interstellar free-electron density in the nearby interstellar medium (ISM) could be approximated by a single power law with a 3D spectral index $\\simeq -3.7$ (very close to the Kolmogorov value of $-11/3$) for scales ranging from $\\sim 10^{10}$~cm to $\\sim 10^{15}$~cm.  By combining their scintillation data with rotation measure (RM) fluctuation measurements and with gradients in the average electron density, they were able to extend the range of their observed nearly-Kolmogorov spectrum up to a scale $\\sim 10^{20}$~cm ($\\sim 30$~pc).  Studies of pulsar and extragalactic radio-source RMs have yielded a wealth of observational information on the spatial structure of the Galactic interstellar magnetic field. However, most studies to date have focused on the large-scale field structure \\citep{sk80,hq94,rl94,id98,hmlq02} and only a few have provided estimates for the rms amplitude and the characteristic scale length of the turbulent magnetic field.  For instance, by adopting a single-cell-size model for the turbulent field and analyzing the residuals with respect to their best-fit model for the large-scale field, \\citet{rk89} obtained a turbulent field strength of $\\simeq 5 \\ \\mu$G and a cell size of $\\simeq 55$~pc.  \\citet{os93}, who did not resort to large-scale field models, organized their pulsar data by pairs of pulsars seen in nearly the same direction and interpreted them in the framework of a single-cell-size model for the random field. With an assumed cell size in the range $10 - 100$~pc, they obtained a random field of amplitude $4 - 6 \\ \\mu$G.  As \\citet{rk89} themselves acknowledged, the turbulent magnetic field cannot be satisfactorily described by a single scale length. \\citet{scs84} and \\citet{sc86} were the first authors to look into the question of how the turbulent magnetic energy in the ISM is actually distributed in the scale range $\\sim 0.01 - 100$~pc. They used the structure function of RMs of extragalactic sources to investigate variations in the composite quantity $n_e \\, {\\bf B}$, where $n_e$ is the free-electron density and ${\\bf B}$ is the magnetic field. Their study suggests that the 3D spatial power spectrum of fluctuations in $n_e \\, {\\bf B}$ at scales $\\ga 0.01$~pc can be described by a power law of index $\\sim -3.1$ with an outer scale estimated at $\\la 90$~pc. Later, \\citet{ms96b} performed a combined analysis of the structure functions of RM and emission measure (EM) to separate the fluctuations in electron density and magnetic field in the scale range $\\sim 0.03 - 100$~pc. They measured the RMs of 38 extragalactic sources located in a small area of the sky near Galactic coordinates $l=144^\\circ$, $b=21^\\circ$, where the EMs were deduced from measured H$\\alpha$ intensities. For scale sizes up to $\\sim 4$~pc, the structure functions of RM and EM were consistent with a 3D-turbulence model in which both electron density and magnetic field fluctuations have Kolmogorov spectra.  Between $\\sim 4$~pc and $\\sim 80$~pc, the structure functions were consistent with 2D turbulence. \\citet{ms96b} also obtained an rms amplitude of $\\sim 1 \\ \\mu$G for the turbulent magnetic field on scales up to 4~pc.   \\begin{figure*} \\includegraphics[angle=270,width=150mm]{distrib.ps} \\caption{Distribution projected on the Galactic plane of the 490 pulsars with well-determined RMs viewed from the North Galactic pole.  Pulsars with a positive RM are denoted with crosses, those with a negative RM are denoted with circles, and the symbol area is proportional to the RM. Superimposed onto the pulsar distribution is an approximate outline of the four known spiral arms.\\label{distr}} \\end{figure*}  It is clear that the magnetic energy spectrum does not abruptly stop at the estimated outer scale of the turbulence, $\\sim 80$~pc. The spectral energy of the magnetic field will extend continuously up to the largest Galactic scales, as a result of an inverse cascade of magnetic helicity \\citep[e.g.,][]{pfl76} and/or due to other physical processes involving compression and shearing at large scales (e.g., Parker 1971; Kulsrud 1986). % turbulent field and the more structured field at larger scales describe the present-day properties of the Galactic magnetic field and provide the necessary observational reference for theoretical models. Knowledge of the complete magnetic energy spectrum can offer a solid observational test for dynamo and other theories for the origin of Galactic magnetic fields.  The purpose of the present paper is to determine the portion of the magnetic spectrum that lies above the strictly turbulent domain, thereby filling the observational gap between the large Galactic scales and the small turbulent scales. Pulsars are unique in providing a direct measure of the average line-of-sight component of the interstellar magnetic field, weighted by the local electron density, from the ratio of RM and dispersion measure (DM). Here we will analyse the available pulsar RM and DM data in combination with their estimated distances to deduce the interstellar magnetic energy spectrum $E_{\\rm B}$ across the scale range $0.5 - 15$~kpc. In \\S~2, we present our pulsar data and explain how they are used to construct the interstellar magnetic energy spectrum, in \\S~3, we review the assumptions underlying our derivation, and in \\S~4, we discuss the significance of our results and compare them with theoretical predictions. ",
        "conclusions": "We have derived the interstellar magnetic energy spectrum over the scale range $0.5 - 15$~kpc, based directly on measurements of pulsar RMs and DMs. The outer scale is a significant fraction of the size of the Galaxy and so contributions from the so-called large-scale (or regular) magnetic field are included. The spectral region covered by our derived spectrum has never been studied observationally before, yet it constitutes a crucial piece of the complete spectrum, whose knowledge is important for our understanding of the behavior and evolution of Galactic magnetic fields.  \\begin{figure}[t] \\includegraphics[angle=270,width=87mm]{Eb_k.ps} \\caption{Composite magnetic energy spectrum in our Galaxy. The thick solid line is the large-scale spectrum derived in this paper. The thin solid and dashed/dotted lines give the Kolmogorov and 2D-turbulence spectra, respectively, inferred from the \\citet{ms96b} study. Their Kolmogorov spectrum, after adjusting to our definition of $k$, can be written as $E_{\\rm B}(k) = C_{\\rm K} \\, (k/k_0)^{-5/3}$ with $C_{\\rm K} = 9.5 \\times 10^{-13}$~erg~cm$^{-3}$~kpc and $k_0=1$~kpc$^{-1}$. The 2D-turbulence spectrum is uncertain (see main text); it probably lies between the dashed ($E_{\\rm B}(k) \\propto k^{-2/3}$) and dotted ($E_{\\rm B}(k) \\propto k^{-5/3}$) lines, which are both scaled to match the Kolmogorov spectrum at $k_{\\rm 3D} = (4~{\\rm pc})^{-1}$. \\label{gs}} \\end{figure}   In Figure~\\ref{gs}, we plot our measured spectrum with that inferred from the \\citet{ms96b} study at smaller scales.  This spectrum has a Kolmogorov slope ($E_{\\rm B}(k) \\propto k^{-5/3}$) up to a scale $l_{\\rm 3D} \\sim 4$~pc and an uncertain shape consistent with 2D turbulence between $l_{\\rm 3D}$ and $l_{\\rm 2D} \\sim 80$~pc. If, as suggested by \\citet{ms96b}, the 2D turbulence is in the form of thin sheets with thickness $\\sim l_{\\rm 3D}$, we could assume that fluctuations with scales larger than $l_{\\rm 3D}$ persist only if their projected wavelength normal to a sheet, $\\lambda \\ \\cos \\theta$, is less than the sheet's thickness $l_{\\rm 3D}$. The fraction of acceptable wave vectors at a given wavenumber $k$ is then $l_{\\rm 3D} / \\lambda = k / k_{\\rm 3D}$. If the acceptable wave vectors preserve a Kolmogorov spectral energy, the magnetic spectrum in the wavenumber range $k_{\\rm 2D} - k_{\\rm 3D}$ is simply the Kolmogorov spectrum multiplied by the above factor, i.e., a $k^{-2/3}$ spectrum.  On the other hand, if 2D turbulence means that the spectral energy is simply redistributed without loss into two dimensions, then the Kolmogorov spectrum $\\propto k^{-5/3}$ extends out to $k_{\\rm 2D}$.  The actual situation probably lies between these two extremes.  The two important features emerging from the composite spectrum (Figure~\\ref{gs}) are (1) a change in slope from $\\alpha = -5/3$ at the small ($\\la 4$~pc) scales to $\\alpha \\simeq -0.37$ at the large ($\\ga 0.5$~kpc) scales and (2) a possibly significant discontinuity between $\\sim 80$~pc and $\\sim 0.5$~kpc.  How do these two features fit with theoretical predictions?  According to our theoretical understanding of the ISM, turbulent energy is injected into the ISM by stellar winds and supernova explosions on scales $\\sim 10 - 100$~pc.  Magnetic energy is then transferred, via a so-called direct cascade, toward smaller and smaller scales. In addition, because interstellar turbulence is helical, there is also an inverse cascade of magnetic helicity toward larger and larger scales, making it possible to amplify or maintain a large-scale magnetic field \\citep{pfl76}.  Because the quantity that is transferred as a whole in the domain of the inverse cascade is magnetic helicity ($H_{\\rm B}(k) \\propto E_{\\rm B}(k) / k$) rather than magnetic energy, the corresponding spectrum is flatter than in the domain of the direct cascade.  Therefore, our finding a flatter spectrum at the larger scales is consistent with the inverse cascade theory.  However, the value obtained for the spectral index ($\\alpha \\simeq -0.37$) differs from current theoretical estimates.  In the theory of 3D MHD helical turbulence developed by \\citet{pfl76}, a numerical solution of the evolution equation for the magnetic energy spectrum yields $E_{\\rm B}(k) \\propto k^{-1}$ at scales larger than the injection scale. As shown by these authors, the same spectral slope is obtained by applying a dimensional Kolmogorov-type argument to the inverse cascade of magnetic helicity. Direct numerical simulations do not yet convey an unambiguous or consistent picture.  \\citet{bp99} obtained numerical spectra consistent with the $k^{-1}$ theoretical prediction.  In contrast, in the simulations of \\citet{bra01}, a wave of spectral magnetic energy propagates toward smaller $k$ under a $k^{-1}$ envelope, but no $k^{-1}$ spectrum actually develops. Instead, magnetic energy builds up at the largest scales and a secondary peak appears at the injection scale. \\citet{mb02} showed that the actual shape of the magnetic energy spectrum depends in fact quite sensitively on the level of fractional kinetic helicity with which the turbulence is forced. According to their simulations, high levels of kinetic helicity lead to double-peaked spectra similar to that obtained by \\citet{bra01}, whereas low levels of kinetic helicity give rise to a single peak at the small resistive scale.  It is clear that one has to be extremely cautious when trying to apply numerical spectra to the real Galactic magnetic field. Current simulations are still highly idealized (e.g., uniform background medium, incompressible or isothermal gas, periodic boundary conditions), and they neglect potentially important factors such as the true nature of turbulent forcing, the presence of large-scale shear, and the actual boundaries of the Galactic disk.  Another difficulty is that the saturated spectra obtained numerically emerge after resistivity comes into play, whereas the Galaxy is probably not yet in the resistive regime (Blackman, private communication).  The discontinuity appearing between the <100pc and >0.5kpc scales in the composite spectrum of Figure~\\ref{gs} can be partly explained by the fact that the different parts of the spectrum apply to different Galactic regions.  As discussed in \\S~3, our portion of the spectrum was obtained from a magnetic energy average over about one third of the inner Galactic disk, while that of \\citet{ms96b} was obtained from a small high-latitude region in the outer Galaxy, where the fields are presumably much weaker than in the inner disk.  However, if the 2D-turbulence portion of the spectrum is as flat as $k^{-2/3}$, the observed discontinuity is so severe that it suggests a genuine excess in the Galactic magnetic energy spectrum at large scales. This excess compared to the amount expected from the inverse cascade alone could indicate that most of the energy input to the large-scale magnetic field occurs directly at large scales.  This interpretation would be consistent with the standard Galactic dynamo theory, in which the large-scale magnetic field is amplified or maintained through the combined action of small-scale helical turbulence, presumably via the inverse cascade leading to the so-called alpha-effect, and large-scale shear, typically associated with the Galactic differential rotation \\citep[e.g.,][]{par71,vr71}. It is widely believed that the large-scale shear acts much more efficiently than the helical turbulence \\citep[e.g.,][]{par71,rss88,bdm+92,fs00,bla00}, especially when the spiral structure of the Galaxy and the streaming motions along spiral arms are taken into account \\citep[e.g.,][]{rre99,eovu00}. A more extreme point of view would be to consider that the large-scale magnetic field is not affected at all by the small-scale turbulent field and that all the energy contained in the large-scale field is directly injected at large scales. Various scenarios have been proposed in this spirit, such as protogalactic collapse and subsequent shearing by the Galactic differential rotation with rapid ambipolar diffusion \\citep{kul86b} or protogalactic collapse and subsequent excitation of spiral arms and bars resulting in nonaxisymmetric, not-perfectly-azimuthal motions \\citep{lc97}."
    },
    "0404/astro-ph0404151_arXiv.txt": {
        "abstract": "The strongest spectroscopic dust extinction feature in the Milky Way, the broad absorption bump at 2175~\\AA, is generally believed to be caused by aromatic carbonaceous materials -- very likely a mixture of Polycyclic Aromatic Hydrocarbon (PAH) molecules, the most abundant and widespread organic molecules in the Milky Way galaxy. In this paper we report identifications of this absorption feature in three galaxies at $1.4 \\lesssim z \\lesssim 1.5$ which produce intervening \\mgii\\ absorption toward quasars discovered by the Sloan Digital Sky Survey (SDSS). The observed spectra can be fit using Galactic-type extinction laws, characterized by parameters [$R_V$, $E(B-V)$] $\\simeq$ [0.7, 0.14], [1.9, 0.13], and [5.5, 0.23], respectively, where $R_V\\equiv A_V/E(B-V)$ is the total-to-selective extinction ratio, $E(B-V)\\equiv A_B-A_V$ is the color-excess. These discoveries imply that the dust in these distant quasar absorption systems is similar in composition to that of Milky Way, but with a range of different grain size distributions. The presence of complex aromatic hydrocarbon molecules in such distant galaxies is important for both astrophysical and astrobiological investigations. ",
        "introduction": "The space between the stars of the Galaxy and external galaxies is filled with gaseous ions, atoms, molecules and tiny dust grains. These interstellar grains, spanning a wide range of sizes from a few angstroms to a few micrometers, are important for the evolution of galaxies, the formation of stars and planetary systems. So far, one of the most best-studied properties of interstellar dust is its obscuration of starlight. The size and composition of interstellar dust are mainly inferred from, respectively, the extinction curve and its spectral features. The strongest spectroscopic interstellar extinction feature in the Galaxy is the broad \\bump\\ absorption bump. This feature was first detected by the Aerobee rocket observations \\citep{St65}. In this work we report the detection of the 2175~\\AA\\ dust exinction bump from individual intervening absorption systems in the spectra of three Sloan Digital Sky Survey \\citep[SDSS;][]{York00} quasars.  The review of \\citet{dr03} provides an excellent summary of the properties of the 2175~\\AA\\ dust extinction bump and theoretical constraints on its carrier. This feature is seen in extinction curves along lines of sight in the Milky Way (MW), the Large Magellanic Cloud (LMC), and some regions of the Small Magellanic Cloud (SMC). Extinction curves in the SMC bar region lack the 2175~\\AA\\ feature. The central wavelength of the feature varies by only $\\pm0.46\\%~(2\\sigma)$ around 2175~\\AA, while its FWHM varies by $\\pm12\\%~(2\\sigma)$ around 469~\\AA. Given its substantial band strength, the carrier responsible for the feature must contain at least one of the most abundant elements C, O, Mg, Si or Fe. Although the exact carrier is unknown, Draine (2003) states that ``it now seems likely that some form of graphitic carbon is responsible'', most likely polycyclic aromatic hydrocarbons (PAHs; Joblin, Leger, \\& Martin 1992; Li \\& Draine 2001; Draine 2003 and references within).  Several previous detections have been reported of the observed \\bump\\ feature in distant galaxies using quasar spectra. \\citet{Pit00} presented a good summary of the state of the field as of three years ago, and ruled out several reported detections, which will not be discussed here.  The highest redshift \\bump\\ bump observation to date is that of \\citet{Vernet01}, who claimed a probable detection in a composite spectrum of radio galaxies at $z\\sim2.5$, although they acknowledged that the detection is difficult to confirm due to the low signal-to-noise ratio (SNR) and blending with \\feii\\ emission at $\\lambda>2300$~\\AA.  Additionally, the dust responsible for the feature does not appear to be located in the host galaxies proper, but rather in the narrow emisson-line regions of these radio galaxies where the physical conditions may be quite different, perhaps leading to different carriers for the bump.  The only previous detection of this feature from an individual intervening absorption system was that of \\citet{Cohen99}, who detected the \\bump\\ feature in a damped Ly$\\alpha$ absorber at redshift $z=0.524$ toward the BL Lac object AO~0235+164 at $z=0.94$. They found a lower dust-to-gas ratio than in the Galaxy and mentioned that the \\bump\\ feature ``suggests there are differences from the average Galactic curve.''\\footnote{The average Galactic curve is parameterized by $R_V=3.1$, where $R_V\\equiv A_V/E(B-V)$ is the ratio of total to selective extinction and $E(B-V)\\equiv A(B)-A(V)$ is the color excess.}  Such differences appear to be rather common when the \\bump\\ feature is detected with data of sufficient quality to model the extinction curve involved. \\citet{Ma97} detected the \\bump\\ feature in the composite absorption spectrum of 96 intervening \\mgii\\ absorption systems at redshifts $0.2 < z < 2.2$. The strength of the feature was roughly consistent with a standard Galactic dust-to-gas ratio. \\citet{Falco99} reported detections in several $z \\lesssim 1$ galaxies responsible for the gravitational lensing of background quasars. Many of the estimated extinction curves do not match the standard $R_V=3.1$ Galactic curve (e.g., $R_V=7.2\\pm 0.1$ was found for a $z_l=0.68$ spiral galaxy). \\citet{Toft00} detected the \\bump\\ feature from the lensing galaxy of B~1152+199 at $z=0.44$ and fitted its extinction curve by a Galactic-type extinction law with $1.3\\lesssim R_V \\lesssim 2.1$ and $0.9 \\lesssim E(B-V) \\lesssim 1.1$. \\citet{Motta02} detected a strong \\bump\\ bump in a lensing galaxy at $z=0.83$. Their data are well fitted by a standard $R_V=2.1\\pm 0.9$ Galactic extinction curve, leading them to suggest that the lens galaxy of SBS 0909+532 contains dust like that of the Galaxy. \\citet{Wuck03} may have marginally detected the \\bump\\ feature in a lensing galaxy (later identified with a damped Ly$\\alpha$ absorber) at $z=0.93$. They found that using extinction curves with a significant 2175~\\AA\\ bump reproduces the data better than curves without this feature. Recently \\citet{mu04} reported that the dust in the $z_l=0.68$ lens galaxy of B~0218+357 shows the \\bump\\ bump but produces a very flat ultraviolet extinction curve with $R_V=12\\pm 2$.  Given the range of extinction curves found in these extragalactic systems, large and well-selected samples extending to higher redshifts are needed to characterize the diversity and evolution of dust properties in the early universe. As a step towards this goal, we present the \\bump\\ feature identified in three individual intervening absorption systems in the spectra of three SDSS quasars. ",
        "conclusions": "We report three direct spectroscopic detections of the \\bump\\ dust absorption feature in quasar absorption systems at redshifts $1.4 \\lesssim z \\lesssim 1.5$. These are the first detections of this feature in individual \\mgii\\ absorption systems.  From fitting the composite SDSS quasar spectra reddened by a CCM Galactic extinction law to the observed spectra, we derive best-fit reddening parameters $E(B-V) \\approx 0.14$, $R_V \\approx 0.7$ for SDSS J1446+0351, $E(B-V) \\approx 0.13$, $R_V \\approx 1.9$ for SDSS J1459+0024 and $E(B-V) \\approx 0.23$, $R_V \\approx 5.5$ for SDSS J0121+0027. The various $R_V$ values in these systems compared to the average Galactic value of $R_V = 3.1$ indicates a wide range of dominant grain sizes among intervening absorption systems. We found that although the intrinsic slope of the quasar continuum in each system is unknown, we can still meaningfully constrain the $R_V$ value in each system. If the presently favored PAH model for the \\bump\\ feature carrier is correct, we have detected complex organic molecules in the young universe, about 9 billion years ago.  The presence of PAHs at such large lookback times may be of astrobiological interest."
    },
    "0404/astro-ph0404198_arXiv.txt": {
        "abstract": "In cold molecular clouds submillimetre emission lines are excited by the ambient radiation field. The pumping is dominated by the cosmic microwave background (CMB). It is usual in molecular line radiative transfer modelling to simply assume that this is the only incident radiation field.  In this paper, a molecular line transport code and a dust radiative transfer code are used to explore the effects of the inclusion of a full interstellar radiation field (ISRF) on a simple test molecular cloud. It is found that in many galactic situations, the shape and strength of the line profiles that result are robust to variations in the ISRF and thus that in most cases, it is safe to adopt the CMB radiation field for the molecular line transport calculations. However, we show that in two examples, the inclusion of a plausible radiation field can have a significant effect on the the line profiles. Firstly, in the vicinity of an embedded massive star, there will be an enhanced far infared component to the radiation field. Secondly, for molecular clouds at large redshift, the CMB temperature increases and this of course also alters the radiation field. In both of these cases, the line profiles are weakened significantly compared to a cloud exposed to a standard radiation field. Therefore this effect should be accounted for when investigating prestellar cores in massive star forming regions and when searching for molecular clouds at high redshift. ",
        "introduction": "Molecular line profiles from prestellar and protostellar objects can potentially yield dynamical information about the collapse process that leads to the formation of stars. In the last decade or so, the common practice has been to obtain line profiles from an infall candidate and then to use a radiative transfer code, together with a dynamical model to interpret the data. Inferences are then made about the validity of one or other of the competing star formation models. However, it is becoming clear that the interpretation of these line profiles can be fraught with difficulty. \\citet{rawlings&yates01} used a self-consistent chemical and dynamical model of collapsing star-forming cores to explore the effects of abundance variations. They showed that the line profiles are very sensitive to the assumed values of the free parameters in the chemical models.  \\citet{ward-thompson&buckley01} have presented a sensitivity analysis of the line profiles to various free parameters, in the context of modelling the Class 0 sources, NGC1333-IRAS2 and Serpens SMM4. The {\\sc stenholm} radiative transfer code (see, e.\\@g.\\@~\\citealt{heaton.et.al93}) was used, adopting analytical fits to the \\citet{shu77} collapse model for the velocity and density profiles, and an optically thin thermal balance model for the dust temperature.  The following free parameters were investigated: the impact parameter, or beam offset from the central position, the beam size (or assumed source distance), the infall velocity, the fractional abundances of the tracer species (assumed not to vary with position), the turbulent velocity and the systemic rotational velocity. One of the most important (general) findings was that the qualitative shape of the line profiles - in particular the velocity separation between the red- and blue-shifted peaks in the self-reversed line profiles is highly sensitive to the turbulent velocity.  This results from the broadening of the absorption profiles of the foreground envelope, whilst the total integrated flux increases with the turbulent velocity dispersion as the effective optical depth of the emitting gas is reduced.  In this paper, the results of numerical experiments to investigate the sensitivity of molecular line profiles to the radiation field incident on a realistic test globule are reported. Firstly, the effects on molecular line profiles of changing this field from the CMB to a more realistic ISRF is investigated. Secondly, the ISRF is modified to simulate the radiation field in the vicinity of embedded massive stars. Finally the ISRF is modified by increasing the CMB temperature, to simulate the radiation field in which molecular clouds at high-redshift will be found. Our study lacks the more self-consistent approach of the combined hydrodynamical/chemical/radiative transfer models of Rawlings and Yates (2001), but serves to show the sensitivities and hence to highlight the diagnostic strengths (and weaknesses) of line profile analyses.  The discussion is limited to two well known molecular tracers used to diagnose collapsing cores: HCO$^+$ (J=3$\\to$2) and $^{13}$CO (J=2$\\to$1, 3$\\to2$). This allows direct comparison with other work; previous studies (eg. Rawlings and Yates, 2001) have shown that the HCO$^+$ transition is particularly likely to be strongly self-reversed and asymmetric, whilst the effects on the CO lines has implications for high-redshift molecular observations. The discussion is also limited to zero beam offset positions so as to allow a simple comparison of the effects discussed in this paper. ",
        "conclusions": "It has been argued that the intensity and shape of submillimetre molecular line profiles as modelled by radiative transfer codes of prestellar and protostellar cores are sensitive to the long-wavelength ambient radiation field illuminating the exterior of the cloud. Some caution is therefore suggested in the modelling of observational results - the ambient environment must somehow be constrained before fitting profiles and making conclusions about the dynamics and chemical composition of the system.  It should be noted that the effect described here is not unknown to workers in computational radiative transfer (it was discussed at the benchmarking exercise that led to the paper by Van Zadelhoff et al 2002), in preparation). However, this is the first time that the astrophysical consequences of this effect have been explored in any detail. In future papers on star-formation, we will - following the preliminary study of \\citealt{ward-thompson&buckley01}) - address the issue of how the velocity structure in general, and the microturbulent velocity in particular, affects the line profiles. It is important to realise that it is essential to characterise correctly the microphysics of potential infall sources. This will eventually allow more robust inferences to be made from line profile data than is currently possible."
    },
    "0404/astro-ph0404367_arXiv.txt": {
        "abstract": "We observed seven central stars of planetary nebulae (CSPN) in the Large Magellanic Cloud (LMC) with the \\emph{Far Ultraviolet Spectroscopic Explorer} (FUSE), and performed a model-based analysis of these spectra in conjunction with \\emph{Hubble Space Telescope} (HST) spectra in the UV and optical range to determine the stellar and nebular parameters.  Most of the objects show wind features, and they have effective temperatures ranging from 38 to 60~kK with mass-loss rates of $\\simeq 5\\E{-8}$~\\Msunyr.  Five of the objects have typical LMC abundances. One object (SMP LMC~61) is a [WC4] star, and we fit its spectra with He/C/O-rich abundances typical of the [WC] class, and find its atmosphere to be iron-deficient.  Most objects have very hot ($T \\gtrsim 2000$~K) molecular hydrogen (\\Htwo) in their nebulae, which may indicate a shocked environment.  One of these (SMP LMC~62) also displays \\OVI\\ 1032-38 nebular emission lines, rarely observed in PN. ",
        "introduction": "\\label{sec:intro}  With respect to the study of planetary nebulae (PN) systems, those in the Large Magellanic Cloud (LMC) are important for two primary reasons.  Attempts to describe the evolution of Galactic PN are hindered by large relative uncertainties in their distances, which carry over into physical parameters such as the stellar luminosity and radius of the central star of the PN (CSPN), and the size and ionized mass of the PN itself.  This obstruction is removed for the PN of the LMC, all of which lie at essentially at the same distance, allowing the physical parameters to be scaled to absolute values. Additionally, the lower metallicity of the LMC relative to the Milky Way allows the role of metallicity in low/intermediate-mass stellar evolution to be assessed.  The largest impact a higher metallicity is expected to have on a star's evolution is a more efficient radiative driving during its windy phases: the asymptotic giant branch (AGB) phase, the post-AGB phase, and/or perhaps a Wolf-Rayet phase [WR].  This would increase the star's mass-loss rate, and thus slow its evolution during these periods.  Such an effect would manifest itself in the form of different relative population ratios in galaxies of varying metallicities.  These implications carry over into galaxy evolution, through chemical evolution and the dynamic interactions between the star's ejected material and the surrounding interstellar medium (ISM).  Although these effects are not as dramatic on an individual basis as those of a massive star, the large fraction of stars that evolve through the CSPN phase make their contribution to the Galactic chemical evolution significant (see, \\eg, \\citealp{marigo:01} for a discussion).  Characterization of an individual PN system requires an understanding of both the PN and its central star.  Several non-LTE (NLTE) spectroscopic studies of Galactic CSPN in the optical range have been carried out, both of CSPN without winds using plane-parallel analyses (\\eg, \\citealp{mendez:85,herrero:90}) and of CSPN with winds using spherical codes (\\eg, \\citealp{leuenhagen:98,koesterke:97,demarco:98}).  The sample of known LMC PN for which HST spectroscopy exists includes only the brightest and rather compact objects.  The nebular continuum of compact (high density) PN typically masks the light of the central star for wavelengths longwards of $\\sim$1215~\\AA\\ \\citep{bianchi:97}, complicating the task of characterizing the two components if relying solely on UV and optical data.  Characteristics of a large sample of LMC CSPN have been inferred from photoionization models of their nebular spectra (\\eg, \\citealp{dopita:91a,vassiliadis:96s,vassiliadis:98s}).  A handful have been analyzed using nebular continuum models in conjunction with modeling of the central star in the UV (\\eg, \\citealp{bianchi:97,dopita:97}).  With the \\emph{Far-Ultraviolet Spectroscopic Explorer} (FUSE) , it is now possible to observe LMC CSPN at far-UV wavelengths where the spectrum is unaffected by nebular contamination.  Far-UV/UV analyses of Galactic CSPN with winds have been carried out by \\citet{koesterke:98b} and \\citet{herald:04b}.  The far-UV is where these hot stars emit most of their observable flux and often exhibit their strongest photospheric and/or wind features. \\citet{herald:04c} performed an analysis of the far-UV spectra of the Galactic CSPN K1-26, and derived a significantly higher temperature (120 vs. 65~kK) than had resulted from analysis of optical spectra \\citep{mendez:85}, illustrating the value of far-UV-based analyses.  Motivated by the above considerations, we observed a sample of seven CSPN with FUSE as part of Bianchi's cycle 2 program B001. The FUSE spectra allow us to separate the nebular and central star components and to characterize the physical parameters of each through modeling. The FUSE range also includes strong absorption due to molecular hydrogen (\\Htwo), which is also modeled concurrently.  This paper is arranged as follows.  The observations and data reduction are described in \\S~\\ref{sec:obs}.  A comparison of the spectra of the objects is presented in \\S~\\ref{sec:description}.  Our models and parameter determinations are described in \\S~\\ref{sec:modeling}.  The implications of our results are discussed in \\S~\\ref{sec:discussion} and our conclusions in \\S~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have analyzed FUSE observations of seven central stars of bright, compact planetary nebulae in the LMC.  Most objects display definite wind features, and we determined their stellar parameters using the stellar atmosphere code CMFGEN to analyze their FUSE far-UV and HST UV spectra.  We also modeled the nebular continua (which contributes to the UV flux) and the atomic and molecular hydrogen absorption along the sight-line (which severely affects the far-UV spectra).  By virtue of their membership of the LMC, the uncertainties in the distances is small.  This is a great advantage over Galactic CSPN, where the distance is the largest source of uncertainty in the analysis.  The objects have a spread of effective temperatures between 35 and 70~kK, with mass-loss rates of $\\Mdot \\sim 5\\E{-8}$~\\Msunyr\\ (the one [WC] star has a mass-loss rate an order of a magnitude larger).  Terminal wind velocities generally increase with increasing effective temperatures, and range between 700---1300~\\kms, a factor of two lower than Milky Way counterparts.  Radii decrease with increasing temperature, as expected.  Their luminosities are similar ($L \\sim 4000$~\\Lsun), and the parameters of the CSPN fall on the He-burning evolutionary tracks of \\citet{vassiliadis:94}, from which we infer post-AGB ages in good agreement with the estimated dynamical ages.  We modeled five of the objects with typical LMC abundances ($z = 0.4$~\\Zsun).  For these, we investigated the effects of the lower LMC metallicity on the wind acceleration by calculating the modified wind momentum.  Contrary to what is expected based on radiative driven wind theory, we did not find any substantial deviation from the wind momentum-luminosity relation that holds for Galactic O and B stars, and roughly holds for Galactic CSPN stars, but the scatter is comparable to the expected difference.  Whether this is because the wind momentum does not as simply depend on the luminosity for CSPN as for O and B stars (as suggested by \\citealp{tinkler:02}) or because of some other effect is unclear.  The one [WC] star of our sample, SMP LMC~61, was modeled using typical He/C/O enriched abundances (products of He-burning).  We also determined a sub-LMC iron abundance (also found by \\citealp{stasinska:04}), perhaps a product of $s$-process nucleosynthesis during a thermally pulsating AGB or post-AGB phase \\citep{lugaro:03,herwig:03}.  This is particularly interesting, as iron opacity is thought to play a key role in the driving of WC winds \\citep{lamers:02,crowther:02}. SMP LMC~61 has a wind terminal velocity about half that of Galactic [WC4] stars from the sample of \\citet{acker:03}.   Our analysis of the far-UV and UV spectra also provided insight into the circumstellar environment of these objects.  Our measured \\HI\\ column densities are higher than those predicted by typical interstellar gas-to-dust relations using our derived reddenings, which implies a significant amount of circumstellar \\HI, presumed to have once been a part of the progenitor object.  We calculated some simple shell models which implied most of this material lies outside the ionized radius of the nebula of these stars, with the exception of SMP LMC~62, for which it seems to lie within.  This object also displays nebular \\OVI\\ \\doublet 1032-38, (such features have been observed in the Galactic CSPN NGC~2371 by \\citealp{herald:04b}).  The high resolution FUSE spectra revealed that these objects have very hot ($T \\gtrsim 2000$~K) molecular hydrogen in the circumstellar environment.  These temperatures may be due to the proximity of the nebular gas to the star, or perhaps to shocks.   In summary, our FUSE observations have allowed us to derive a set of stellar and wind parameters for young CSPN in the LMC that are unhampered by the distance uncertainties that plague Galactic studies. In addition to revealing the flux of the hot star (which is obscured at longer wavelengths by the nebular flux), the far-UV observations also revealed hot molecular hydrogen surrounding these young CSPN."
    },
    "0404/astro-ph0404373_arXiv.txt": {
        "abstract": "{ We present in this paper a detailed study of a new sample of large angular size FR I and FR II radio galaxies and compare the properties of the two classes. As expected, a pure morphology based distinction of FR Is and FR IIs corresponds to a break in total radio power. The radio cores in FR Is are also weaker than in FR IIs, although there is not a well defined break power. We find that asymmetry in the structure of the sample members must be the consequence of anisotropies in the medium where the lobes expand, with orientation playing a minor role. Moreover, literature data and our observations at kiloparsec scales suggest that the large differences between the structures of FR I and FR II radio galaxies must arise from the poorly known central kiloparsec region of their host galaxies. We analyze the sub-sample of giant radio galaxies, and do not find evidence that these large objects require higher core powers. Our results are consistent with giant radio galaxies being the older population of normal FR I and FR II objects evolving in low density environments. Comparing results from our sample with predictions from the radio luminosity function we find no evidence of a possible FR II to FR I evolution. Moreover, we conclude that at $z\\sim 0.1$, one out of four FR II radio sources has a linear size above 500 kpc, thus being in an advanced stage of evolution (for example, older than $\\sim 10$ Myr assuming a jet-head velocity of 0.1c). Radio activity seems to be a short-lived process in active galaxies, although in some cases recurrent: five objects in our sample present signs of reactivation in their radio structures. ",
        "introduction": "Radio galaxies, with linear sizes reaching up to several megaparsecs, are possibly the largest individual objects in the Universe.  It is widely accepted that they originate from highly energetic non-thermal processes occurring in the nucleus of the so-called active galaxies (Blandford \\& Rees \\cite{blandford}; Rees \\cite{rees1}). According to the standard model of active galactic nuclei (AGNs), a super-massive black hole with a mass between $10^6$ and $10^9 M\\odot$, resides in the center of the active galaxy, powered by an accretion disk surrounded by a torus formed by gas and dust. In about 10\\% of these AGNs, there is intense synchrotron radio emission produced in a bipolar outflow of relativistic particles expelled perpendicularly to the plane of the disk and extending to distances reaching the megaparsec scales.  The reason why an AGN presents or not powerful radio emission is matter of strong debate.  While there is increasing evidence about the existence of super-massive black holes in the center of AGNs, and even at the nuclei of non active galaxies (Macchetto \\cite{macchetto}; Kormendy \\& Gebhardt \\cite{kormendy}), it has been argued that the presence or not of intense radio emission might be due to the rotation velocity of the black hole (e.g. Wilson \\& Colbert \\cite{wilson}; Cavaliere \\& D'Elia \\cite{cavaliere}), or to its total mass and the efficiency of accretion (McLure \\& Dunlop \\cite{mclure}; Dunlop et al. \\cite{dunlop2}).  Considering a natural evolutionary sequence of radio galaxies, jets emanating from the center of activity start boring their way through the interstellar medium first, reaching the galactic halo and in some large cases the intergalactic medium. Finally, the lobes of radio galaxies which have ceased their central activity expand and disappear in the external medium. Radio sources representing these different phases of evolution are currently known adding support to this scenario: compact symmetric objects (CSOs; Wilkinson et al. \\cite{wilkinson}) are thought to be young radio galaxies (e.g. Owsianik \\& Conway \\cite{owsianik}), while the giant radio galaxies (GRGs; defined as those with a projected linear size\\footnote{We assume that H$_{0}=50$km s$^{-1}$ Mpc$^{-1}$ and q$_{0}=0.5$} $\\ge$ 1 Mpc) are probably old objects at the latter stages of evolution (Ishwara-Chandra \\& Saikia \\cite{ishwara}). Relic radio sources found in clusters of galaxies might correspond to the last detectable emission from ``dead'' radio galaxies (e.g. Komissarov \\& Gubanov \\cite{komissarov}; Slee et al. \\cite{slee}).  However, the degree of influence of parameters other than the age (e.g. source power, conditions of the external medium) in the evolution of radio galaxies is not clear . For example, although there is observational evidence supporting the young source scenario for CSOs, it has also been argued that CSOs are short lived objects which never reach the size of their big relatives (Readhead et al. \\cite{readhead}). At the other extreme, GRGs could be the result of normal radio galaxies expanding in very low density environments permitting them to reach their overwhelming sizes.  But they could also result from very powerful core activity, or both conditions must apply for a radio galaxy to become a giant. Complicating the previously outlined evolutionary sequence, some radio galaxies seem to wake up after a dormant phase of absence or much lower activity (e.g. Lara et al. \\cite{lara}).  Moreover, the presence of super-massive objects in many non--active galaxies argue in favor of activity as a short transition period in most, if not all, (elliptical) galaxies, and that the ``menace\" for future activity is present at the center of every galaxy.  This paper is the last of a series of three devoted to the study of a sample of large angular size radio galaxies which try to address some of these open questions. Definition of the sample and radio maps in one side, and images and spectroscopic data on the other, were presented by Lara et al. (\\cite{paperI}, hereafter Paper I) and (\\cite{paperII}, Paper II), respectively. The sample, covering a sky area of 0.842 steradians\\footnote{Papers I and II erroneously mention a sky area of $\\pi$ steradians. However, results and discussion in those papers are not affected by this value.} and spectroscopically complete at the 80\\%, consists of 84 radio galaxies selected from the NRAO\\footnote{National Radio Astronomy Observatory} VLA\\footnote{Very Large Array, operated by the NRAO} Sky Survey (NVSS; Condon et al. \\cite{nvss}) under the following selection criteria (see Paper I for details): declination above $+60^{\\circ}$, total flux density at 1.4 GHz greater than 100 mJy and angular size larger than 4$\\arcmin$. All sources in the sample have redshift below 0.75. In this paper we present the general results of the study of the sample, with distinction between Fanaroff-Riley type I (FR I; Fanaroff \\& Riley \\cite{fanaroff}) and II (FR II) radio galaxies. Special attention is devoted to show the properties of GRGs, of which this sample contains 37 members. ",
        "conclusions": "We present in this paper, the last of a series of three, the general properties of a sample of 84 radio galaxies selected from the NVSS, with total flux density at 1.4 GHz $\\ge 100$ mJy, declination above $+60^{\\circ}$ and angular size larger than $4'$. This study is based on radio (Paper I) and optical (Paper II) observations of the sample members. The main conclusions of our work can be summarized as follows: \\begin{itemize}  \\item From a pure morphological distinction of the sample members according to its FR I or FR II type morphology, we confirm that both groups are separated by a break in the total radio luminosity, at $\\log P_{t}$(1.4 GHz)$= 25.2$ (W/Hz).  \\item There is a small population of radio sources (7\\% in our sample) which show simultaneously properties of FR I and FR II radio galaxies, with a mean total radio power in the region of transition between the FR I and FR II families.   \\item We find that asymmetry in the structure of FR I and FR II radio galaxies can be explained as due to anisotropies in the medium through which the jet propagates. Arguments based on orientation effects would require too high jet-head advance velocities.  \\item FR I radio galaxies have in general lower power cores than FR II radio galaxies, although the distinction is not as clear as in the total radio power. There is a correlation, independent of the FR I - FR II dichotomy, between the core and the total radio power, consistent with that derived by Giovannini et al. (\\cite{giov88}, \\cite{giov01}). The prominence of the radio cores in our sample (evidenced through the $P_{CN}$ orientation indicator) can be explained as a result of the evolution of extended radio galaxies. The evolution of extended radio galaxies implies a total power decrease in GRGs of a factor 10--100, in agreement with Kaiser et al.(\\cite{kaiser}).  \\item Considering the known properties of FR I and FR II radio galaxies at parsec, subkiloparsec and megaparsec scales, and the properties of their optical spectra, we suggest that the central kiloparsec region plays a crucial role in explaining the FR I - FR II dichotomy, where differences in the radio structure between the two families seem to appear and whose properties might determine the presence or not of emission lines. We note however that the properties of this region are surely linked with the properties of the accretion disk which ultimately determines the source of ionization and the different core radio power in FR I and FR II sources, but apparently not the jet production mechanism.  \\item The luminosity tracks predicted by the model of Kaiser et al. (\\cite{kaiser}) explains well the dearth of high luminosity and large size radio galaxies. Our study and other recent searches of GRGs, even if much more sensitive and systematic than previous works on large size radio galaxies, do not find high luminosity giant sources.  \\item We do not find evidence of the GRG phenomenon as being due to stronger core radio power. Our observations are consistent with GRGs as the older population of normal FR I and FR II radio galaxies expanding in low density environments.  \\item A compact radio core is detected in 100\\% of the sample members, indicating that core activity is present in all objects. However, we find in our sample 5 objects with signs at kpc scales of having passed through different phases of activity. Nearby galaxies are found in two of these ``reactivated'' objects, but more detailed observations are required to study the role of galaxy merging in triggering different cycles of activity.  \\item Comparing the FR I and FR II abundances as a function of redshift with the predictions of the RLF, we do not find evidence of a possible FR II to FR I evolution of radio galaxies.  \\item From the RLF and our sample, we find that about one out of four FR II type radio galaxies with $z\\sim 0.1$ have linear sizes larger than 500 kpc, thus being in an advanced stage of evolution.   \\end{itemize}"
    },
    "0404/astro-ph0404145_arXiv.txt": {
        "abstract": "We present an update on the results of our monitoring observations of the X-ray remnant of supernova (SN) 1987A with the {\\it Chandra X-Ray Observatory}. As of 2002 December, we have performed a total of seven observations of SN 1987A, which allows us to monitor the details of the earliest stage of the supernova remnant evolution in X-rays. The high angular resolution images from the latest data reveal developments of new X-ray bright spots in the northwestern and the southwestern portions of the remnant as well as changes on the eastern side. The observed soft X-ray flux is increasing more rapidly than ever, and the latest 0.5$-$2 keV band flux ($f_X$ $\\sim$ 6 $\\times$ 10$^{-13}$ ergs cm$^{-2}$ s$^{-1}$) is four times brighter than three years earlier when this monitoring began. The overall X-ray emission is primarily from the blast wave shock with $kT$ $\\sim$ 2.4 keV. As the blast wave approaches the dense circumstellar material, the contribution from the decelerated slow shock ($kT$ $\\sim$ 0.22 keV) to the observed X-ray emission is becoming significant. The increase of this slow shock contribution over the last two years is particularly noticeable in the western half of the remnant. These results indicate that the shock front is now reaching the main body of the inner circumstellar ring and that SN 1987A will be a complete ring with dramatic brightening in coming years. Based on the best-fit two-shock spectral model, we derive approximate densities of the X-ray-emitting regions ($n_e$ $\\sim$ 235 cm$^{-3}$ for the fast shock and $n_e$ $\\sim$ 7500 cm$^{-3}$ for the slow shock). There is no direct observational evidence to date for a neutron star associated with supernova remnant 1987A. We obtain an upper limit on the observed X-ray luminosity of any embedded point source ($L_X$ $\\le$ 1.5 $\\times$ 10$^{34}$ ergs s$^{-1}$) in the 2$-$10 keV band. The X-ray remnant continues to expand linearly at a rate of 4167 km s$^{-1}$. ",
        "introduction": "OBSERVATIONS \\& DATA REDUCTION}  The seven {\\it Chandra} observations of SNR 1987A are presented in Table~\\ref{tbl:tab1}. We first screened all data sets with the flight timeline filter and turned off the pixel randomization for the highest possible angular resolution. We then corrected the spatial and spectral degradation of the ACIS data caused by radiation damage, known as the charge transfer inefficiency (CTI; Townsley et al. 2000), with the methods developed by Townsley et al. (2002a), before further standard data screening by status, grade, and energy selections. The expected effects of the CTI correction include an increase of the number of detected events and improved event energies and energy resolution \\citep{town00,town02a}. ``Flaring'' pixels were removed and {\\it ASCA} grades (02346) were selected. Photons between 0.3 keV and 8.0 keV were extracted for the data analysis. Lightcurves around the source region were examined for possible contamination from variable background emission and no severe variability was found. The pileup fraction was relatively small ($<$10\\%) and thus was ignored in the analysis.  We then applied the ``sub-pixel resolution'' method \\citep{tsunemi01} to improve the angular resolution of the images to better than the CCD pixel size. A typical improvement in the angular resolution by $\\sim$10\\% is expected from this method \\citep{mori01}. The angular size of SNR 1987A is small (the inner ring is only about 1$\\farcs$6 across; e.g., Burrows et al. 1995; Jakobsen et al. 1991), and the ACIS detector pixel size (0$\\farcs$492) is not adequate to fully resolve the remnant. In order to further improve the effective angular resolution of the ACIS images, we deconvolved the images using a maximum likelihood algorithm \\citep{rich72,lucy74} as described in our previous works \\citep{bur00,park02}. For the image deconvolution, we used 0$\\farcs$125 sky pixels for the data sets with low photon statistics ($\\la$1000 source counts) and 0$\\farcs$0625 pixel size for the high statistics data sets ($\\ga$1000 source counts). ",
        "conclusions": ""
    },
    "0404/gr-qc0404102_arXiv.txt": {
        "abstract": "Understanding quantum theory in terms of a geometric picture sounds great. There are different approaches to this idea. Here we shall present a geometric picture of quantum theory using the de-Broglie--Bohm causal interpretation of quantum mechanics. We shall show that it is possible to understand the key character of de-Broglie--Bohm theory, the quantum potential, as the conformal degree of freedom of the space--time metric. In this way, gravity should give the causal structure of the space--time, while quantum phenomena determines the scale. Some toy models in terms of tensor and scalar--tensor theories will be presented. Then a few essential physical aspects of the idea including the effect on the black holes, the initial Big--Bang singularity and non locality are investigated. We shall formulate a quantum equivalence principle according to which gravitational effects can be removed by going to a freely falling frame while quantum effects can be eliminated by choosing an appropriate scale. And we shall see that the best framework for both quantum and gravity is Weyl geometry. Then we shall show how one can get the de-Broglie--Bohm quantum theory out of a Weyl covariant theory. Extension to the case of many particle systems  and spinning particles is discussed at the end. ",
        "introduction": "In this century, physicists have been departed from 19th century physics, in two ways. The first was the generalization and bringing the old idea of \\textit{frame independence} or \\textit{general covariance}, in a manifest form. The result of this effort was the pioneer general relativity theory, in which the gravitational effects of matter are identified with the geometry of the space--time.  The enigmatic character of this theory is just the above-mentioned property, i.e. the interconnection of gravity and general covariance. When one tries to make a general covariant theory, one is forced to include gravity.  The main root of this interconnection is the \\textit{equivalence principle}. According to the equivalance principle, it is possible to go to a frame in which gravity is locally absent, and thus the special theory of relativity is applicable locally. Now using the general covariance and writing down anything in a general frame, we will get the general relativity theory\\cite{weinberg}.  The second was the investigation of the quantal behavior of matter, that leads to the \\textit{quantum theory},  according to which a great revolution  appeared in physics. In order to explain the atomic world, the quantum theory threw out two essensial classical concepts, the \\textit{principle of causality} and \\textit{the dogma of formulation of physics in terms of motion in  space--time} (motion dogma). The first one is violated during a measurement process, while the second does not exist at any time.  After the appearance of quantum mechanics, it was proven that not only do the ordinary particles show quantal behavior but mediators of the fundamental forces also do so. In this way quantum electrodynamics, quantum chromodynamics and quantum flavor dynamics were born. But the construction of quantum gravitodynamics or quantum gravity, and its application to cosmology, is considerably very problematic\\cite{ash}. These difficulties may be mainly divided into two categories. Some of them are related to the conceptual problems of the standard quantum mechanics, while others are specific to gravity, and are in fact related to the classical features of gravity theory. The first category includes the measurement problem and the meaning of the wave function of the universe, while the vanishing of the hamiltonian which leads to the time independence of the wave function and nonrenormalizability, belong to the second category.  From a fundamental physical viewpoint, in contrast to the general theory of relativity which is the best theory for gravity, standard quantum mechanics is not the only satisfactory way of understanding the quantal behavior of matter.  One of the best theories explaining the quantal behavior of matter but remaining faithful to the principle of causality and the motion dogma, is the de Broglie--Bohm quantum theory.\\cite{bohm} According to this theory, all the enigmatic quantal behavior of the matter results from a self-interaction of the particle. In fact, any particle which exerts a \\textit{quantum force} on itself can be expressed in terms of a \\textit{quantum potential} and which is derived from the particle wave function.  The celebrated property of the de Broglie--Bohm quantum theory is the following property. \\textit{At anytime, even when a measurement is done, the particle is on the trajectory given by Newton's law of motion, including the quantum force}. During a measurement, the system is in fact a many-body system (including the particle itself, the probe particle, and the registrating system particles). When one writes down the appropriate equation of motion of all the particles,\\footnote{There is a consistent de Broglie--Bohm quantum theory for a many particle system} and when one considers the very fact that we know nothing about the initial conditions of the registerating system particles, one sees how the projection postulate of quantum mechanics came about\\cite{bohm}. Accordingly the result of any measurement is one of the eigenvalues of the operator related to the measured quantity with some calculable probability distribution. The de-Broglie--Bohm quantum theory of motion, is a causal theory which although behaves as the Copenhagen quantum mechanics at the statistical level, it has non of the conceptual difficulties of the standard quantum mechanics. It is well proved that the causal theory reproduces all the results of the orthodox quantum theory\\cite{bohm}, as well as predicting some new results (such as time of tunneling through a barrier\\cite{cush}) which in principle lets the experiment to choose between the orthodox and the causal quantum theories. Perhaps the most important point about the causal theory is that it presents a causal deterministic description of the reality. So it looks very natural to make a quantum theory of gravity in the spirit of the de-Broglie--Bohm quantum theory of motion. In the standard form of this theory, the classical gravity should be viewed as a field. Then it is possible to construct Bohmian metric trajectories . This is what is essentially done in \\cite{gil}.  A way of struggling with quantum gravity is to use the \\textit{minisuperspace} of the conformal degree of freedom of the space-time metric\\cite{nar}. This approach has several fruitfull results. It admits non-perturbative calculations, and it is very useful for studying quantum cosmology. Because the isotropic and homogeneous space-time used in   cosmology is conformally flat. In addition, by including the effects of the back--reactions of the quantum variable (i.e. the conformal factor) on the background metric, one arrives at some extended form of Einstein's equations. These semi-classical equations lead to non--singular cosmological solutions and they have the correct classical limit. In this approach, by merely quantizing one degree of freedom of the space-time metric, and by considering the back--reaction effects, the time independence problem is solved. This is achieved because of the extension of Einstein's equations. But it must be noted that the physical meaning of the quantum variable, i.e. the conformal factor, is not clear in this approach. The non-singularity of the results of the above approach rests on the theorem\\cite{kem} which states that \\textit{for any singular metric , there is some appropriate conformal factor, in such a way that conformal metric is non-singular}.  The present work tries to combine the de-Broglie--Bohm quantum theory of motion and gravity in a very different way. The foundation of this approach is the de-Broglie remark\\cite{deb} that \\textit{the quantum theory of motion for relativistic spinless particles is very similar to the classical theory of motion in a conformally flat space-time}. The conformal factor is related to the Bohm's quantum potential. We shall present a generalization and an appropriate formulation of this remark. That is to say, we geometrize Bohmian mechanics according to the de-Broglie remark. Then, it can be seen that the effects of gravity on geometry and the quantum effects on the geometry of the space-time are highly coupled. In fact there are two contributions to the background metric the gravitational quantal effects of matter which constitute the energy--momentum tensor. Since in the evaluation of the quantal part the background metric is used, the gravitational and quantal contributions to the background metric are so highly coupled that no one without the other has any physical significance.  It must be pointed out here that as aresult the conformal factor is meaningless as the enesemble density goes to zero and the geometry looses its meaning at this limit. This is a desired property, because it is in accord with Mach's principle, which states that for an empty universe the space--time should be meaningless. In subsection(\\ref{s1})\\cite{geo} the authors, as a first step towards the formulation of the above conclusion, introduced the quantum conformal degree of freedom via the method of Lagrange multipliers. In this way there are a set of equations of motion describing the background metric, the conformal degree of freedom and the particle trajectory. A corollary of this theory is that one can always work in a gauge (classic gauge) in which no quantum effect be present or in a gauge (quantum gauge) in which the conformal degree of freedom of the space--time metric is identified with the quantum effect This, in its turn, leads to dramatic departures from the classical prediction, when both the effects of gravity and quantum on geometry are considerable, i.e. around those areas of the space-time which are singular according to the classical theory.  As a different approach in ref \\cite{mot} the authors symmetrized the Brans--Dicke theory by a conformal transformation. And arrive to a particle interpretation suggesting that the quantum aspect of matter can be geometrized.  In \\cite{conf}, the conformal transformation was applied only  to the space--time metric. Other quantities like mass, density and so on were assumed to posses no transformation. This is because the above conformal transformation which incorporates the quantum effects of matter into a specific conformal factor, is in fact a scale transformation. As the conformal transformation is more general than scale  transformation which is used in \\cite{geo}, it seems preferable  to make a conformal transformation, in  which all physical quantities are transformed, instead of making only a scale transformation. In reference \\cite{conf}, it is shown that by the conformal transformation the equation of motion would be transformed to an equation in which there is no quantum effects. As a result, the geodesic equation would be changed to the one without the quantum force. This means that it is possible to have two identical pictures for investigating the quantal effects of matter in the curved space--time background. According to the first picture, the space--time metric contains only the gravitational effects of matter. The quantum effects affect the path of the particles via the quantum force. In the second picture, the space--time metric is related to the previous by a conformal factor and contains the gravitational and  quantal effects of matter.  This shows that the quantum as well as the gravitational effects of matter have geometrical nature. The second picture mentioned above provides a unified geometrical framework for understanding the gravitational and quantum forces. Accordingly, we call the conformal metric as the physical metric (containing both gravity and quantum) and the other metric is the background metric (including only gravity).  The above-mentioned theory,\\cite{geo} has a problem.  In this theory, it is assumed that one deals with an ensemble of similar particles with density. In Bohm's theory, the quantum potential exists for a single particle as well as for an ensemble. In the case of a single particle, the interpretation of the quantum potential is in terms of an hypothetical ensemble.  Note that in the above theories, the ensemble is a real one, not an hypothetical one, because, the energy--momentum tensor of the ensemble is appeared and has physical effects. As we shall show in subsection (\\ref{s2})\\cite{qgg}, we have solved this problem and the theory would work both for a single particle and for an ensemble.  In subsection (\\ref{s2}) we shall show that it is possible to make a pure tensor theory for quantum gravity. As a result we shall show that the correct quantum conformal degree of freedom would be achieved, and that the theory works for a particle as well as for a real ensemble of the particle under consideration and that it includes the pure quantum gravity effects. We shall do all of these by trying to write the quantum potential terms in terms of geometrical parameters, not in terms of ensemble properties.  The important point about both references\\cite{geo} and \\cite{conf} is that in order to fix the relation of the conformal degree of freedom of the space--time metric and the quantum potential, the method of lagrange multiplier is used and in this way they are a little artificial. In subsection (\\ref{s3}) we shall show\\cite{st} that in the framework of the scalar--tensor theories, it is possible to write an action principle, in which both gravitational and quantum contributions to the geometry are included and that the conformal degree of freedom of the space--time metric is fixed at the level of the equations of motion not needing the method of lagrange multiplier.  Next in subsection (\\ref{s4})\\cite{ss} we attend to the double scalar case because in some theories such as superstring and Kaluza--Klein, it is more useful\\cite{bar}. In both of these  theories the gravitational interaction includes two other fields in addition to the metric field. In string frame(or Jordan frame)one of them is coupled nonminimally to gravity as in the Brans--Dicke theory and the other is coupled minimally to gravity, but has a nontrivial coupling with the first scalar field. Note that, in these theories one can couple both the scalar fields minimally to gravity by a conformal transformation (Einstein frame). As a result, the question that the physical interpretation must be presented in which frame, is an open problem\\cite{far}. On the other hand we shall show that using two scalar fields, one can relax this preassumption and on the equations of motion, the correct form of quantum potential will be achieved.  In subsections (\\ref{s5}), (\\ref{s6}) and (\\ref{s7}) some general solutions are obtained. using these, the important question that if this quantum gravity theory leads to some new results, is investigated. That solusions are used for black holes and bigbang in subsections (\\ref{s8}) and (\\ref{s9})\\cite{qgg}. In subsection (\\ref{s10})\\cite{clu} we are interested in investigating whether this theory has anything to do with the cluster formation or clustering of the initial uniform distribution of matter in the universe. The problem of cluster formation is an important problem of cosmology and there are several ways to tackle with it\\cite{bor}. Here we don't want to discuss those theories, and our claim is not that the present work is a good one. Here we only state that \\textit{the cluster formation can also be understood in this way}. It is a further task to decide if this work is in complete agreement with experiment or not.  A special aspect of the quantum force is that it is highly nonlocal. This property, is an experimental matter of fact \\cite{bell}. Since the mass field represents the conformal degree of freedom of the physical metric, quantum gravity is expected to be highly nonlocal. In the subsection (\\ref{s11})\\cite{non} this is shown explicitley for a specific problem.  From a different point of view it has been believed for a long time that the long range forces (i.e. electromagnetism and gravity) are different aspects of a unique phenomena. So they must be unified. Usually it is proposed that one must generalize Einstein's general relativity theory to have a geometrical description of electromagnetic fields. This means to change the properties of the manifold of general relativity. Using higher dimensional manifolds\\cite{wes}, changing the compatibility relation between the metric and the affine connection\\cite{we} and using a non-symmetric metric\\cite{ein} are some examples of the attempts towards this idea. In all the above approaches, the additional degrees of freedom correspond  to the components of the electromagnetic potential. The second idea leads to the Weyl's gauge invariant geometry. Apart from the electromagnetic aspects of Weyl geometry, it has some other applications. Some authors believe that Weyl geometry is a suitable framework for quantum gravity. E.g. in a series of papers \\cite{odin} a succesful approach to Weyl quantum gravity and conformal sector in quantum gravity is presented. The authors have used an effective theory based on integrated conformal anomaly dynamics, in the infrared region. They also have considered a sigma model action which is the most general version of a renormalizable theory in four dimmensions. They have investigated the phase structure and the infrared properties of conformal quantum gravity and then extend its results to higher derivative quantum gravity.  Also in ref\\cite{whe} a new quantum theory is proposed on the basis of Weyl picture which is purely geometric. The observables are introduced as zero Weyl weight quantities. Moreover any weightful field has a Weyl conjugate such as complex conjugate of the state vector in quantum mechanics. By these dual fields, the probability can be defined. These are the elements of a consistent quantum theory which is equivalent to the standard quantum mechanics. Moreover it is shown that the quantum measurement and the related uncertainty  would emerged from Weyl geometry naturally. In this theory when the curl of Weyl vector is zero, we arrive at the classical limit. By noting the transformation relation of Weyl vector, it is concluded that the change of length scale is only a quantum effect.  One more approach to geometrize quantum mechanics can be found in \\cite{wood}. Here a modified Weyl--Dirac theory is used to join the particle aspects of matter and Weyl symmetry breaking. This is also a geometrization of quantum mechanics. Also one can find the relation of quantum potential, the basic character of Bohm's theory, to the fundamental geometric properties, especially to the curvature of the space-time using Weyl geometry in \\cite{san}. Furthermore in \\cite{sid} Sidharth considers the geometrical interpretation of quantum mechanics from the point of view of non-commutative non-integrable geometry.  In the present work we shall look at the conformal invariance at the quantum level. Does the quantum theory lead us to any characteristic length scale and thus break the conformal symmetry? Or conversely the quantum effects lead us to a conformal invariant geometry? In section (\\ref{s12}) we shall discuss these questions in the context of the causal quantum theory proposed by Bohm\\cite{bohm} and use our new way of geometrization of quantum mechanics introduced in here. We emphasize that what we shall show that our specific geomerization of quantum mechanics procedure (based on Bohmian quantum mechanics) can be better understood in the Weyl framework. This is different from Weyl quantum gravity approaches like those of \\cite{odin}.  We shall show that the Weyl vector and the quantum effects of matter are connected. We shall see how the conformal symmetry emerges naturally by considering quantum effects of matter. Finally in section (\\ref{s13}) we show that the Weyl--Dirac theory is a suitable framwork for identification of  the conformal degree of freedom of the space--time with the Bohm's quantum mass.  From a similar perspective, Quiros and et all\\cite{quir} discuss the space--time singularity by the geometrical dual representation in general relativity. On this basis they emphasis on the Weyl integrable geometry as a consistent framework to describe the gravitational field.  Finally in section (\\ref{s20}) we shall investigate possible extension of our results in two ways. First analyzing the case of many--particle systems and second, inclusion of spin. ",
        "conclusions": "We have used an approach which is different from other existing ones, which try to combine the gravitational and quantal effects.  By the investigation of the quantum effects of matter in the framework of Bohmian mechanics, we have shown that the motion of a particle with quantum effects is equivalent to its motion in a curved space-time.  We have investigated the coupling of purely gravitational effects and purely quantal effects of the particle, by considering a general background space-time metric. The use of the de-Broglie--Bohm quantum theory of motion, instead of the standard Copenhagen quantum mechanics, has at least three advantages. First, that the inherent problems of the standard quantum mechanics are not present. Second, that the conceptual problems of the standard quantum gravity like the meaning of the wave-function, are circled. Finally, the equivalence of quantum effects of matter and a curved space-time, which is our most important result, is achieved through this point of view. This leads directly to the minisuperspace of conformal degree of freedom, in which, the conformal factor has now a clear physical meaning. Other problems like, time- dependence which is necessary to understand the evolution of the universe, consideration of the back--reaction effects, and so on, are also handled in the present theory. An important property of this theory is that the conservation laws are the same as those of the classical theory. As it is shown, if one applies this theory to the case of quantum cosmology, it is possible to solve the equations of motion nonperturbatively (i.e. exactly). One sees that there is no singularity at small times (provided one assumes that the quantum coupling constant of radiation is negative). It is smeared out by the quantum effects. Two remarks are in order here. First, in principle, the above model can be applied to the non flat Friedmann universes and similar results would emerge. Second, in the present theory, the quantum effects are only those of matter. The effect of these quantal behaviours on the background space-time is achieved via the modified Einstein equations, i.e. via the back--reaction effects. This means that if one removes the matter, the quantum effects of the background geometry would disappear. In order to generalize the present theory to a fully theory of quantum gravity, the quantum effects of gravity must be included.  As explained before the keystone of Bohm's theory is the \\textit{quantum potential}. Any particle is acted upon by a quantum force derived from the quantum potential. The quantum potential is itself resulted from some self-field of the particle, the wave function. Since the quantum potential is related only to the norm of the wave function and because of Born's postulate asserting that the ensemble density of the particle under consideration is given by the square of the norm of the wave function, the quantum potential is obtained resulted from the ensemble density. The non understandable point of Bohm's theory is just this. How does a particle know about its hypothetical ensemble? When the hypothetical ensemble is a real one, i.e.~when there are is large number of similar particles just like the particle under consideration, quantum potential can be understood. It is a kind of interaction between the particles in the real ensemble. But when one deals with only one particle the quantum potential is interaction with the other hypothetical particles!!!  On the other hand, quantum potential is highly related to the conformal degree of freedom of the space--time metric. In fact, the presence of the quantum force is just like having a curved space--time which is conformally flat and the conformal factor is expressed in terms of the quantum potential. In this way one sees that quantum effects are in fact geometric effects. Geometrization of the quantum theory can be done successfully, but still, there is the problem of the ensemble noted above.  Here we have shown that if one tries to geometrize the quantum effects in a purely metric way, the ensemble problem would be overcome. In addition it provides the framework for bringing in the purely quantum gravity effects.  A point about the geodesic equation must be noted here. In the background metric, this equation resembles the geodesic equation in Brans--Dicke theory. Consideration of the matter quantum effects, leads to the physical metric in which a particle moves on the geodesic of Branse--Dicke theory written in Einstein gauge. This point supports the suggestion that the discussion of quantum gravity requires a scalar--tensor theory. Previously this was suggested when discussing Bohmian quantum gravity \\cite{cov}.  Next it is shown that it is possible to write a scalar--tensor theory which automatically leads to the correct equations of motion. This has the advantage that the conformal factor would be fixed by the equations of motion, and not by introducing a lagrangian multiplier by hand. We see that  the matter distribution determines the local curvature of the space--time (in confirmity with Mach's principle). Furthermore, from the matter equation of motion one can see that the cosmological constant (a large scale structure constant) and the quantum potential\\footnote{In Bohmian quantum mechanics, observable effects of quantum potential may appear at both large and small scales, depending on the shape of the ensemble density\\cite{bohm}.} are coupled together. This is another manifestation of the Mach's principle.  Also  we have constructed a scalar--tensor theory with two scalar fields, for which the equations of motion lead to the correct form of quantum potential. We have not only shown that quantum effects are geometrical in nature, but also \\textit{derive} the form of quantum potential. This specific form for quantum potential is a result of the equations of motion.  It must be noted that in this theory both the scalar fields interact with the cosmological constant.  Hence, the presence of the cosmological constant (even very small) is essential in order the theory works.  Note that the interaction between the cosmological constant and quantum potential represents a connection between the large and small scale structures.  We presented a toy model, and investigated its solutions. Also it is shown that the initial singularity is removed by quantum effects.  Finally we saw that one can formulate a \\textit{generalized equivalence principle} which states that gravitation can be removed locally via an appropriate coordinate transformation, while quantum force can be removed either locally or globally via an appropriate scale transformation. So the natural framework of quantum and gravity is Weyl geometry. The most simplest Weyl invariant action functional is written out. It surprisingly leads to the correct Bohm's equations of motion. When it applied to cosmology it leads to time decreasing cosmological and gravitational constants. A phenomena which is good for describing their small values."
    },
    "0404/astro-ph0404235_arXiv.txt": {
        "abstract": "We present near-IR spectra of $\\Ha$ emission from 13 galaxies at $z\\sim2$ in the GOODS-N field.  The galaxies were selected primarily because they appear to have elongated morphologies, and slits were aligned with the major axes (as determined from the rest-frame UV emission) of 11 of the 13.  If the galaxies are elongated because they are highly inclined, alignment of the slit and major axis should maximize the observed velocity and reveal velocity shear, if present. In spite of this alignment, we see spatially resolved velocity shear in only two galaxies.  We show that the seeing makes a large difference in the observed velocity spread of a tilted emission line, and use this information to place limits on the velocity spread of the ionized gas of the galaxies in the sample: we find that all 13 have $v_{0.5} \\leq 110$ \\kms, where $v_{0.5}$ is the velocity shear (half of the velocity range of a tilted emission line) that would be observed under our best seeing conditions of $\\sim0\\farcs5$.  When combined with previous work, our data also indicate that aligning the slit along the major axis does not increase the probability of observing a tilted emission line.  We then focus on the one-dimensional velocity dispersion $\\sigma$, which is much less affected by the seeing, and see that the elongated subsample exhibits a significantly {\\it lower} velocity dispersion than galaxies selected at random from our total $\\Ha$ sample, not higher as one might have expected. We also see some evidence that the elongated galaxies are less reddened than those randomly selected using only UV colors.  Both of these results are counter to what would be expected if the elongated galaxies were highly inclined disks.  It is at least as likely that the galaxies' elongated morphologies are due to merging subunits. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404003_arXiv.txt": {
        "abstract": "We present velocity dispersion measurements of 14 globular clusters in NGC~5128 (Centarus~A) obtained with the MIKE echelle spectrograph on the 6.5m Magellan Clay telescope. These clusters are among the most luminous globular clusters in \\cena\\ and have velocity dispersions comparable to the most massive clusters known in the Local Group, ranging from 10 -- 30 \\kms. We describe in detail our cross-correlation measurements, as well as simulations to quantify the uncertainties. These 14 globular clusters are the brightest \\cena\\ globular clusters with surface photometry and structural parameters measured from the {\\it Hubble Space Telescope}. We have used these measurements to derive masses and mass-to-light ratios for all of these clusters and establish that the fundamental plane relations for globular clusters extend to an order of magnitude higher mass than in the Local Group. The mean mass-to-light ratio for the \\cena\\ clusters is $\\sim 3 \\pm 1$, higher than measurements for all but the most massive Local Group clusters. These massive clusters begin to bridge the mass gap between the most massive star clusters and the lowest-mass galaxies. We find that the properties of \\cena\\ globular clusters overlap quite well with the central properties of nucleated dwarf galaxies and ultracompact dwarf galaxies. As six of these clusters also show evidence for extratidal light, we hypothesize that at least some of these massive clusters are the nuclei of tidally stripped dwarfs. ",
        "introduction": "\\label{sec:intro}  Globular clusters provide valuable snapshots of the formation history of galaxies and their large sizes and luminosities make them the most readily observable sub-galactic constituents. In addition, globular clusters exhibit surprisingly uniform properties that suggests a common formation mechanism. They are well-fit by single-mass, isotropic King models \\citep{king66}, which describe clusters in terms of scale radii, central surface brightness, and core velocity dispersion. Detailed studies of globular clusters in our Galaxy have shown that in fact they only inhabit a narrow range of the parameter space available to King models \\citep{djorgovski95,mclaughlin00} and other globular cluster systems in the Local Group also appear to approximately follow the same relations \\citep{djorgovski97,dubath97a,dubath97b,barmby02,larsen02}. These parameter correlations trace a globular cluster fundamental plane that is analogous to, but distinct from, the fundamental plane for elliptical galaxies \\citep{dressler87,djorgovski87,bender92,burstein97}.  Globular cluster studies that include internal kinematics have been confined to the Local Group due to the faint apparent magnitudes of more distant extragalactic globular clusters. These studies have therefore only included the globular cluster systems of spiral and dwarf galaxies and not the globular cluster systems of large ellipticals. Yet the globular cluster systems of ellipticals are a particularly interesting regime as they probe both a new morphological type and one likely to have exhibited a different and more complex formation history. The globular cluster systems of elliptical galaxies, as well as many spirals such as our own, have bimodal color distributions suggestive of multiple episodes of formation \\citep[\\eg][]{kundu01,larsen01a}. Models for the formation of these globular cluster systems posit that one of these populations may be the intrinsic population of the galaxy and subsequent mergers resulted in the second population either as the result of a new episode of globular cluster formation \\citep*{schweizer87,ashman92,forbes97} or accretion of globular cluster systems from other galaxies, including accretion of globular clusters by tidal stripping from other members of a cluster of galaxies \\citep*{cote98}.  NGC~5128 (Centarus~A), as the nearest large elliptical galaxy, is arguably the best source for extending detailed globular cluster studies outside of the Local Group. While \\cena\\ is the central galaxy of a large group, rather than a giant elliptical in a cluster, it likely had a similar formation history to its larger cousins. The most relevant similarity is the strong evidence for a recent, gas-rich merger \\citep[for a recent review of \\cena\\ see][]{israel98}. Estimates of the size of the \\cena\\ globular cluster population suggest that it has a total of $\\sim 2000$ clusters, approximately a factor of 3 more than the entire Local Group \\citep{harris84}. Simple scaling arguments suggest that \\cena\\ should possess a number of extremely massive globular clusters and therefore is not only a good target for study of the globular cluster system of an elliptical galaxy, but also for study of how well the fundamental plane relations established locally apply to more massive globular clusters. A recent photometric and spectroscopic study of \\cena\\ by \\citet*{peng04a,peng04b} concluded that the metal-rich globular clusters may have a mean age of $5^{+3}_{-2}$ Gyr, while an analysis of their photometric data yields a metallicity range of $-2.0$ through $+0.3$  \\citep{yi04}  The most massive globular clusters can also be used to explore connections between the formation processes for star clusters and galaxies. While fundamental plane studies \\citep[\\eg][]{burstein97} clearly illustrate a significant mass gap between the most massive Galactic globular clusters and the least massive dwarf galaxies, there have been encroachments into this gap from both sides. For many years, studies have speculated that at least some globular clusters may be the remains of tidally-stripped dwarf galaxy nuclei \\citep{zinnecker88,freeman93,bassino94}. Two of the most massive globular clusters in the Local Group, $\\omega$Cen in our Galaxy and G1 in M31, have been interpreted as the nuclei of tidally stripped dwarfs \\citep{meylan01,gnedin02,bekki03b}. From the galaxy side, recent studies of nucleated dwarf galaxies in the Virgo cluster \\citep*{geha02} and ultracompact dwarf galaxies in the Fornax cluster \\citep{drinkwater03} show some similarities between these least-massive galaxies and the most massive globular clusters. This mass gap may thus reflect the scarcity of the most massive globular clusters and the difficulty of kinematic measurements for the least massive, lowest surface-brightness dwarf galaxies, rather than a physical separation.  In this paper we present velocity dispersion measurements for 14 globular clusters in \\cena. These globular clusters were selected from the {\\it Hubble Space Telescope} (\\hst) study of \\citet{harris02} and therefore have well-measured structural parameters. These data are combined to estimate masses for these clusters, masses that are among the largest known for any star clusters and comparable to the nuclei of the lowest-mass galaxies. In the next sections we describe the observations, data processing, and velocity dispersion measurements. Analysis of these measurements is described in \\S5 and the potential link between star clusters and galaxies is explored in \\S6. Our results are summarized in the final section. Throughout this paper we adopt the distance of $3.84 \\pm 0.35$ Mpc for \\cena\\ determined by \\citet{rejkuba04} from the brightness of the tip of the red giant branch and the Mira period-luminosity relation. ",
        "conclusions": "\\label{sec:discuss}  The strong dynamical and photometric similarities of these massive globular clusters to the nuclei of dwarf galaxies has interesting implications for models which posit stripped dwarf galaxy nuclei as the origin of the most massive globular clusters. The most massive Local Group clusters, such as $\\omega$Cen in our own Galaxy and G1 in M31, have both been discussed and modeled as stripped dwarf nuclei \\citep{freeman93,meylan01, gnedin02,bekki03b}. The position of $\\omega$Cen and G1 in Figures~\\ref{fig:fp} and \\ref{fig:kappa} show that the \\cena\\ clusters occupy a similar region of the fundamental plane and $\\kappa-$space. All of these massive clusters have comparable masses, mass-to-light ratios, and central surface brightnesses.  Support of the interpretation of these clusters as tidally-stripped galaxy nuclei is provided by the tentative detection of extratidal light by \\citet{harris02} for six of the globular clusters (marked with an $x$ in Table~\\ref{tbl:params} and \\ref{tbl:mass}, {\\it partly filled circles} in the figures). In the context of the stripped-dwarf model, it is tempting to view this extratidal light as the last vestiges of the extended dwarf envelope around these nuclei, although extratidal light may also be present due to the evaporation of cluster stars by two-body relaxation. Deeper observations of these and other massive clusters would be extremely valuable to confirm and quantify this extratidal light. Such observations would also better quantify the ellipticities of these globular clusters, which are comparable to those of $\\omega$Cen and G1 but greater than those of dwarf nuclei.  These globular clusters may also simply represent the upper end of the globular cluster mass function in \\cena. The similarities between these globular clusters and the most massive globular clusters in the Local Group are more established than the interpretation of the better-studied Local Group clusters as tidally-stripped dwarf nuclei. In addition, there is evidence for yet more massive young star clusters that demonstrate that there is overlap between the masses of the most massive star clusters and the least massive dwarf galaxies. The most massive star cluster known is W3, a cluster with $\\sigma = 45$ \\kms\\ in the merger remnant NGC~7252 \\citep{maraston04}. The inferred mass of this cluster is nearly $10^8$ \\msun, even more massive than the nuclei of the dwarf galaxies studied by \\citet{geha02}. Even if we were to make the extreme suggestion that all of the \\cena\\ globular clusters are in fact relic dwarf nuclei, the existence of W3 demonstrates there is overlap between the masses of galaxies and star clusters.  On the other side of the mass gap, the ultra-compact dwarf galaxies in the Fornax cluster \\citep{drinkwater03} are actually {\\it less} massive than W3 and overlap with the most massive \\cena\\ clusters. These dwarfs have velocity dispersions of $\\sigma = 24 - 37$ \\kms, effective radii of 10 -- 30 pc, masses in the range $10^{7-8}$ \\msun, and mass-to-light ratios of 2 -- 4 in solar units \\citep{drinkwater03}. They are thus very comparable to the larger and more massive of the \\cena\\ globular clusters. Numerical models for these ultra-compact dwarfs have shown that they can form from nucleated dwarfs that have been stripped of their envelopes by tidal forces in a cluster \\citep{bekki03a}. This stripping process can explain the ultra-compact dwarfs in Fornax, although it will operate over approximately a factor of 2 smaller radius in a smaller group like \\cena\\ due to decreased tidal shear. The radius for the tidal stripping of \\cena\\ globular cluster progenitors is likely to be on order 10~kpc \\citep[][see their Figure~7]{bekki03a}, comparable to the projected distances of 5 -- 23~kpc for some of the \\cena\\ globular clusters in this study. It is therefore plausible that tidal stripping could have transformed at least some nucleated dwarf galaxies into these \\cena\\ globular clusters.  Another aspect of the tidally-stripped dwarf model is that their dark matter halos cannot be too cuspy as otherwise the relatively concentrated dark matter core will be too effective at retaining the stellar envelope. In their study, \\citet{bekki03a} used the dark matter profile of \\citet{salucci00} (originally proposed by \\citet{burkert95} for dwarf galaxies), rather than the more commonly adopted profile of \\citet*{navarro96}, because it has a flatter core. However, this requirement for a relatively flat dark matter core may remove one possible explanation for the higher mass-to-light ratio of these massive globular clusters. While a residual dark matter halo from their dwarf galaxy past is a possible explanation for the larger mass-to-light ratios of these \\cena\\ clusters, a dark matter profile with a flat, low-density core will also be more efficiently stripped away \\citep{bekki03a}. An alternative explanation for the high mass-to-light ratios of these massive globular clusters is if they formed in a starburst with a truncated stellar initial mass function \\citep{charlot93}. A mass function with relatively more low-mass stars would produce a higher mass-to-light ratio.  The detailed shapes and kinematics of the nuclei of nucleated dwarfs and the most massive globular clusters may provide one way to further investigate the potential connection between these two populations. The most massive globular clusters have significant ellipticities \\citep{harris02}, while this does not appear to be the case for the nuclei of nucleated dwarfs \\citep{geha02}. However, the same tidal forces that strip a dwarf envelope may also induce significant ellipticities. The importance of rotational flattening and anisotropies in the velocity distribution may also serve to distinguish between dwarf nuclei and globular clusters.  The hypothesis that some of the most massive globular clusters are the nuclei of galaxies offers an appealing explanation for recent evidence of an intermediate-mass black hole in G1 \\citep*{gebhardt02}. The mass of this black hole falls on the same $M_{BH} - \\sigma$ relationship for galaxies and suggests that the formation mechanisms for black holes in star clusters and galaxy spheroids are similar. If G1 is instead simply a tidally-stripped, nucleated dwarf galaxy, the problem is reduced in complexity and only one physical mechanism for black hole growth is required to explain the $M_{BH} - \\sigma$ relation.  An alternate interpretation of the similarity between the most massive globular clusters and the nuclei of nucleated dwarfs is that the later are simply star clusters that have migrated to or formed at the centers of these dwarfs. The properties of these objects would then be probes of massive star clusters in different environments, rather than of actual overlap between the properties of the most massive star clusters and the least massive galaxies. This interpretation also explains their similar location in $\\kappa-$space, although stands in contrast to the simple explanation for the intermediate-mass black hole in G1."
    },
    "0404/astro-ph0404529_arXiv.txt": {
        "abstract": " ",
        "introduction": "The Alpha Magnetic Spectrometer (AMS) is a cosmic ray (CR) experiment that will operate on the International Space Station (ISS) for three years, measuring the particle spectra in the rigidity range from 0.2 GV to 2 TV.  The first version of the detector, called AMS-01, was flown on board of the space shuttle Discovery in 1998, from June 2 to 12.  The second version of the detector, called AMS-02, will operate on the ISS starting from year 2007.  AMS-01 collected more than hundred million events, and the results were published about the CR protons and antiprotons, helium, electrons and positrons, and the upper limit on the antimatter ratio (for a review, see Ref.~\\cite{ams01}).  The physics goals of the AMS experiment are: (1) the search for antinuclei among the CR, that would be a smoking gun in favor of the existence of cosmic antimatter in form of anti-stars (figure~\\ref{fig1}, left panel, shows the upper limits on the antihelium to helium ratio); (2) the search for signatures of the supersymmetric dark matter candidates, looking at the antiproton and positron spectra (figure~\\ref{fig1}, right panel, shows the expected distortion of the positron to electron ratio that would be produced by the annihilation of 65 GeV/$c^2$ neutralinos \\cite{lamanna03}); (3) the high-statistics study of the CR particle spectra up to 1--2 TV rigidity.  \\begin{figure}[t] \\includegraphics[width=0.44\\textwidth]{antiHe-fraction.eps} \\hfill \\includegraphics[width=0.48\\textwidth]{e+e-bump.eps} \\caption{Upper limits on the CR antihelium to helium ratio (left panel), and a possible signature of the neutralino annihilation in the ratio between CR positrons and electrons (right panel).}\\label{fig1} \\end{figure}  \\begin{figure}[t!] \\centering \\includegraphics[width=0.4\\textwidth]{ams-02-3d.eps} \\caption{The AMS-02 detector.}\\label{fig2} \\end{figure}  The AMS-02 detector (figure~\\ref{fig2}) is able to make multiple measurements of the energy lost by the particles traversing it, of their track and of their velocity \\cite{ams02}.  The basic technique for particle recognition lays on the measurement of the energy loss, giving the particle charge absolute value, of the track curvature, giving the particle rigidity, and of the velocity value and direction, that gives the charge sign together with the curvature measurement.  The anticoincidence system (ACC) vetoes tracks crossing the magnet, hence all particles must traverse the transition radiation detector (TRD), the time of flight system (TOF), the silicon tracker and the proximity focusing \\v{C}erenkov detector (RICH).  About half of them will also produce a shower in the electromagnetic calorimeter (ECAL), placed below the RICH.  RICH and TOF are used to distinguish electrons and positrons from antiprotons and protons down to 0.2 GeV, whereas TRD and ECAL allow for the e/p separation up to 300 GeV, that is the upper energy limit for the antiproton and positron measurements.  Nuclei can be separated from antinuclei up to the maximum detectable rigidity, i.e.~up to 2 TV.  Finally, gamma rays can be detected in the largely unexplored energy range above few tens of GeV.  This paper focuses on the ``ordinary'' CR astrophysics items that can be studied with the help of the AMS-02 detector.  They are necessary steps before anyone will trust the possible exotic results of the measurement, but they are also important for the understanding of our Galactic environment. ",
        "conclusions": ""
    },
    "0404/astro-ph0404286_arXiv.txt": {
        "abstract": "{ We present the results of stellar photometry of the polar-ring galaxy \\object{NGC~660} using the Hubble Space Telescope's archival data obtained with the Wide Field Planetary Camera 2. The final list of the resolved stars contains 550 objects, a considerable part of which are blue and red supergiants belonging to the polar ring. The analysis of the Colour Magnitude Diagram for polar ring stars shows that it is best represented by the isochrones with metallicity Z = 0.008. The process of star formation in the polar ring was continuous and the age of the youngest detected stars is about 7~Myr. ",
        "introduction": "NGC~660, known for a long time as a peculiar galaxy containing two inclined dust lanes \\citep{benv76}, was included by \\citet{whit90} in their catalogue of Polar-Ring Galaxies (PRGs), candidates and related objects as a possible candidate (C\\,-\\,13). The existence of two different kinematic systems was established from radio H{\\sc i} \\citep{gott90} and CO \\citep{comb92} observations, and the galaxy was classified as a kinematically confirmed PRG \\citep{arnab93}.  An extensive study of NGC~660 was performed by \\citet{vanDr95}. The authors give a detailed description of the host galaxy (which is spiral, while host galaxies in PRGs are usually lenticular or elliptical) and its ring, which is not polar but is inclined to the disc of the host galaxy by $55\\degr$. The authors  notice that the ring is very massive and  may be stable \\citep{sparke86}. Multicolour optical observations made with Schmidt telescope with low resolution ($5\\farcs5$) show that the ring is blue ($V-I_{\\rm c}~=1\\fm0$). From comparison with models of single burst star formation, the authors claim that the age of the ring is about 2~Gyr. However, many H{\\sc ii} regions \\citep{young88, armus} observed in the ring  and the high mass of molecular gas sufficient for active star formation \\citep{comb92} may point to more recent star formation. Detecting blue supergiants in the ring may confirm this suggestion.  The galaxy is one of the nearest PRGs (D$\\sim13$~ Mpc, H$_{\\rm 0}$~=75~km/s/Mpc) and resolving its ring into stars probably may be achieved using HST images. Our study of PRGs NGC~2685 and NGC~4650A whose distances are further than NGC~660 confirms this possibility (Karataeva et al.,\\,2004, Paper I). The HST archive contains several images of NGC~660, unfortunately not very deep. These images form the observational basis for our work devoted to resolving the ring of the galaxy into stars.  \\begin{figure}[htb] \\vbox{\\special{psfile=\"0080fig1.ps\" hscale=70 vscale=70 hoffset=-30 voffset=-290 angle=0}% \\vspace{9.5cm} } \\caption{DSS2 $8\\arcmin \\times 11\\arcmin$ image of NGC~660 with WFPC2 footprint overlaid. \\label{f:N660_ima} } \\end{figure} ",
        "conclusions": "We have resolved the ring of NGC~660 into stars. Some blue and red supergiants were found with ages of about 7$\\div$9~Myr. The relatively low metallicity of these stars is in accordance with the data for polar rings in other PRGs.  It is very important to reach the region of red giants in CMDs. This might give not only the information about ring age but also assist in choosing the model of ring formation because the merging scenario \\citep{bourncomb03} predicts the existence around the galaxy of a halo from old/intermediate age stars. Longer exposures with HST may allow us to construct such CMDs."
    },
    "0404/astro-ph0404079_arXiv.txt": {
        "abstract": "s{The influence of neutrino trapping (NT) on the early cooling evolution of hot proto quark stars (PQS) with initial temperatures in the range $T\\sim 40~$ MeV is studied. Within a simplified model for the neutrino transport it is shown that the time for reaching neutrino opacity temperature of $T_{opac}\\sim 1$ MeV is about 10 sec. This is an order of magnitude larger than without NT and of the same order as the duration long gamma ray bursts. } ",
        "introduction": "\\label{sec:intro} Gamma ray bursts (GRBs) are among  the most intriguing phenomena in the Universe, see \\cite{Piran} and references therein. If the energy is emitted isotropically, the measured energy release is of the order of $10^{53}\\div 10^{54}$~erg and it is a puzzle to explain the engine of a GRB. However, there is now a compelling evidence that the gamma ray emission is not isotropic, but displays a jet-like geometry. When the emission is collimated within a narrow beam a  smaller energy, of the order of $10^{52}$~erg, would be sufficient for the GRBs \\cite{Frayl} but their sources are not yet understood. There is growing evidence for a connection of GRBs to supernovae now from emission features in some GRB afterglows, e.g. GRB 990707 \\cite{Amati}, GRB 991216 \\cite{Piro}, GRB 000214 \\cite{Antonelli} and, most recently GRB 030329 \\cite{Fynbo}. As in the realm of a supernova explosion a compact star is likely to be born, it has been conjectured (see \\cite{Berezhiani} and references therein) that a phase transition from hadronic to deconfined quark matter might power the GRB. However, although the energy release might be in the right order of magnitude, the collapse timescale is too short ($\\sim $ several ms) to explain long GRBs with a duration of several tens of seconds. Recently, it has been suggested \\cite{Berezhiani} that the deconfinement transition in a compact star might be an example for a nucleation process of quark matter droplets which is a quantum tunneling process between metastable states, with a sufficient delay, depending on the surface tension of the quark matter droplet. This approach has been criticized \\cite{DNV} since during the supernova collapse the protoneutron star can be heated up to temperatures of the order of the Fermi energy $T\\sim \\varepsilon_F \\approx 30 \\div 40$ MeV so that the thermal fluctuations would dominate over quantum ones and make the phase transition sufficiently fast without the delay claimed in Ref. \\cite{Berezhiani}.  In the present contribution we consider the cooling evolution of a hot protoneutron star above the neutrino opacity temperature $T_{opac}\\sim 1$ MeV \\cite{Reddy}, so that the neutrino mean free path is by orders of magnitude smaller than the size of the star. As it has recently been estimated \\cite{Aguilera} and also reported at this conference \\cite{Agui_Korea}, the neutrino untrapping transition occuring when the star cools below $T_{opac}$ might serve as an engine of a GRB. The question to be considered here is whether the energy release by neutrino emission can be sufficiently delayed due to neutrino trapping so that the typical duration of a long GRB could be explained.  \\begin{figure}[th] \\psfig{figure=Graph15e.eps,width=\\textwidth,angle=0} \\caption{Schematic representation of the evolution of the composition of compact star matter with increasing density. $Y_i=n_i/n_B$ is the fraction of particles of species i per baryon and $n_B/n_0$ is the baryon density in units of the nuclear saturation density $(n_0 = 0.16~ {\\rm fm}^{-3})$. \\label{Graph15e}} \\end{figure} ",
        "conclusions": ""
    },
    "0404/astro-ph0404553_arXiv.txt": {
        "abstract": "{We have calculated the \\submm spectral indices of 26 Herbig Ae/Be stars, for which we can determine the infrared spectral energy distribution (SED). We find a clear correlation between the strength of the ratio of the near- to mid-infrared excess of these sources, and the slope of the \\submm energy distribution. Based on earlier multi-dimensional modeling of disks around Herbig Ae stars, we interpret this as a correlation between the geometry of the disk (flared or self-shadowed) and the size of the grains: self-shadowed disks have, on average, larger grains than their flared counterparts. These data suggest that the geometry of a young stellar disk evolves from flared to self-shadowed. ",
        "introduction": "Herbig Ae/Be (HAEBE) stars are the somewhat more massive analogues of the T Tauri stars, which are low-mass young-stellar objects. The spectral energy distribution (SED) of HAEBE stars is characterized by the presence of an infrared (IR) flux excess, due to circumstellar dust and gas. The geometry of this circumstellar matter is believed to be disk-like (e.g. Mannings \\& Sargent 1997, 2000; Fuente et al. 2003).  Meeus et al. 2001 (henceforth M01) classified 14 isolated HAEBE sample stars into two groups, based on the shape of the SED. \\textit{Group I} contains the sources in which a rising mid-IR (20--100~$\\mu$m) flux excess is observed; these sources have an SED that can be fitted with a power-law and a black-body continuum. \\textit{Group II} sources have more modest mid-IR excesses; their IR SEDs can be reconstructed by a power-law only. M01 suggested phenomologically that the difference between the two groups is related to the disk geometry.  Dullemond 2002 (henceforth D02) and Dullemond \\& Dominik 2004 (henceforth DD04) have modelled young stellar disks with a self-consistent model based on 2-D radiative transfer coupled to the equation of vertical hydrostatics. The model consists of a disk with an inner hole ($\\sim$0.5 AU), a puffed-up inner rim and an outer part. The outer part of the disk can be flared (as in Chiang \\& Goldreich 1997), but can also lie entirely in the shade of the inner rim. The SEDs of flared disks display a strong mid-IR flux excess, while self-shadowed disks have a much steeper slope towards long wavelengths. D02 can explain the difference in SED shape in HAEBEs (as expressed by the classification of M01) as the result of a different disk geometry; flared-disks systems have group I SEDs, sources with self-shadowed disks can be linked to group II objects.  In this paper we calculate the spectral index of the \\submm (350--2700~$\\mu$m) SED for a sample of 26 HAEBE stars. We look for a connection between this index and the geometry of the circumstellar disk. In Section~\\ref{OandA} we classify the sample sources and present the computed \\submm spectral indices. Interpretations of the results obtained in the latter section are given in Section~\\ref{interp}. We summarize our conclusions in the final Section~\\ref{concl}. ",
        "conclusions": "}  We showed in this paper that there is a significant difference between the \\submm spectral indices of HAEBE stars with different IR SEDs. This could very well mean that the grain size distribution of the cold grains in a HAEBE star depends on the disk geometry.  The HAEBE systems with large IR excesses (group I, i.e. flaring-disk objects) have smaller cold grains than systems with more moderate IR excesses (group II, i.e. self-shadowed-disk sources). In other words, it appears that \\textit{the flaring geometry of a circumstellar disk can only persist when a sufficiently large population of small grains exists in the disk}.  The circumstellar-disk models of DD04 agree qualitatively with the results presented here. Disks in which the vast majority of the grains have grown to larger sizes, lead automatically to self-shadowed geometries. Models with small grains result in a mid-infrared SED akin to group I sources. Models with grain sizes larger than $\\sim$1 mm not only display mid-infrared SEDs more resembling those of group II sources, but also change their \\submm slope from larger values of $n$ to a black-body character ($n=3$). DD04 note that the \\submm slope changes with somewhat smaller grain sizes than are required to alter the mid-IR SED from group I to group II. Therefore, according to the models, the spectral index can have all values larger than 3 in group I sources, and all group II sources should have a black-body \\submm slope. This is consistent with the observed trend in which group I sources are observed to have $3.17<n<4.59$ and group II sources have $n$ very close to 3. The DD04 models predict that grain growth resulting in a change of $n$ from values similar to those found in the ISM ($> 3$) to the optically thin black-body value ($n=3$) occurs before the disk becomes self-shadowed. However the DD04 models do not model the dynamical processes leading to grain growth. At present the time scale of this transition from larger $n$ to $n=3$ is unknown, so we are unable to estimate the fraction of group I sources that should have $n=3$. Therefore we have no reason to assume a priori that our distribution of (sub-)mm-spectral-index values is incompatible with the interpretation outlined in Section~\\ref{gg}.   In the present paper, we report a clear correlation between the shape of the mid-IR SED and the \\submm slope in our sample. Previous authors have shown that variations in \\submm slope within systems with proto-planetary disks are caused by variations in dust particle size. D02 have shown that the shape of the mid-IR SED is a proxy for disk geometry (flared versus self-shadowed). In accordance with these results, we interpret our find as a correlation between disk geometry and grain size. If one assumes that in the course of disk evolution grains will grow, this offers the first observational indications to the idea that flaring disks evolve into self-shadowed disks."
    },
    "0404/astro-ph0404309_arXiv.txt": {
        "abstract": "We present observations of the unusual microlensing event OGLE~2003--BLG--235/MOA~2003--BLG--53. In this event a short duration ($\\sim$7 days) low amplitude deviation in the light curve due a single lens profile was observed in both the MOA and OGLE survey observations. We find that the observed features of the light curve can only be reproduced using a binary microlensing model with an extreme (planetary) mass ratio of $0.0039^{+11}_{-07}$ for the lensing system. If the lens system comprises a main sequence primary, we infer that the secondary is a planet of about 1.5 Jupiter masses with an orbital radius of $\\sim3$ AU. ",
        "introduction": "Gravitational microlensing occurs when a foreground object passes through or  very near the line of sight of a background source star generating a well known symmetric light curve profile. If the foreground  lens object is a star with an orbiting planet, then the presence of the planet  may be detectable via a brief disturbance in the single lens light curve \\citep{mp, gl}. This effect can potentially be utilized to detect  planets with masses ranging from those of gas giants right down to terrestrial planets \\citep{br}.  The short timescales of these deviations, ranging from a few days for giant  planets to hours for terrestrial planets, and their unpredictability, present  considerable challenges in any observational program. While some encouraging results have been obtained \\citep{benn99, alb00, rhie, bond02, jp}, no firm detections of planets by microlensing have previously been obtained.  In this Letter we report observations, obtained by OGLE and MOA, of the event OGLE~2003--BLG--235/MOA~2003--BLG--53 (hereafter O235/M53) that was independently detected in both survey programs. We observed a 7 day deviation that was strongly detected in both surveys. We show that an extreme mass ratio binary microlensing model best reproduces the observed features in the light curve. ",
        "conclusions": "In Section~3 we concluded that the observed light curve of O235/M53 is best described by a binary lensing model with an planetary mass ratio of $q=0.0039$. Our definition of the planetary nature of the secondary lens by means of the mass ratio is optimal when the mass ratio can be measured, but it is useful to consider other possible definitions. Another potential dividing line between planets and brown dwarfs is the Solar metalicity threshold for sustained Deuterium burning at 13.6 M$_\\mathrm{J}$, although Deuterium burning itself has little relevance for planet formation.  The situation in the case of O235/M53 was helped by the measurement of finite source effects. If the lens is a main sequence star, then as shown in the previous section, it must be an M dwarf with a $\\sim 1.5\\,$M$_\\mathrm{J}$ planetary companion\\footnote{The only other M dwarf star known to have planetary companions is Gliese 876 \\citep{marcy}.}. There is a non-negligible chance that the lens is a white dwarf and a much smaller chance that it is a neutron star, but in both cases, a planetary companion below the nominal 13.6 M$_\\mathrm{J}$ threshold is required. Only in the unlikely case of a massive black hole primary, could the secondary be outside the range traditionally associated with a planet.  There are some prospects for follow-up observations of this event. Our measurements of the finite source effects imply a proper motion of the lens with respect to the source of $\\mu = \\theta_\\mathrm{E}/t_\\mathrm{E} = 3.1 \\pm 0.4~\\mathrm{mas/yr}$. High resolution imaging carried out $\\sim 10\\,$years from now with JWST or adaptive optics systems should be able to resolve the lens and source stars providing direct measurements \\citep{han,alcock01} of the color and brightness of the lens, as well as confirmation of the proper motion measurement.  We present these observations as a demonstration of the planetary microlensing phenomenon.  The power of microlensing is in its ability to acquire statistics on many systems \\citep{gest}. These include planets in wide orbits, very low mass planets, and even planets in other galaxies \\citep{covone, bond02}. The challenge now to the microlensing community is to develop effective strategies to find more planetary microlensing events.  Numerical photometry of OGLE~2003--BLG--235/MOA~2003--BLG--53 is available from the websites for OGLE \\url{http://ogle.astrouw.edu.pl} and MOA \\url{http://www.physics.auckland.ac.nz/moa}."
    },
    "0404/astro-ph0404415_arXiv.txt": {
        "abstract": "We provide the first direct lifting of the mass/anisotropy degeneracy for a cluster of galaxies, by jointly fitting the line of sight velocity dispersion and kurtosis profiles of the Coma cluster, assuming an NFW tracer density profile, a generalized-NFW dark matter profile and a constant anisotropy profile. We find that the orbits in Coma must be quasi-isotropic, and find a mass consistent with previous analyses, but a concentration parameter 50\\% higher than expected in cosmological $N$-body simulations. We then test the accuracy of our method on realistic non-spherical systems with substructure and streaming motions, by applying it to the ten most massive structures in a  cosmological $N$-body simulation. We find that our method yields fairly accurate results on average (within 20\\%), although with a wide variation (factor 1.7 at 1$\\,\\sigma$) for the concentration parameter, with decreased accuracy and efficiency when the projected mean velocity is not constant with radius. ",
        "introduction": "The kinematical (``Jeans'') analyses of near-spherical structures using the spherical Jeans equation for stationary systems: \\[ \\frac{\\rm d}{{\\rm d} r}  (\\nu \\sigma_r^2) + \\frac{2 \\beta}{r} \\nu \\sigma_r^2 = - \\nu \\frac{G\\,M}{r^2} \\ , \\] where $\\nu$ is the 3D density distribution of the tracer population, suffer from the fact that this one equation involves two unknowns: the mass distribution $M(r)$ and the velocity anisotropy $\\beta = 1 - \\sigma_\\theta^2/\\sigma_r^2$ (we make use of the symmetry of spherical systems yielding $\\sigma_\\phi = \\sigma_\\theta$). In some cases of simple anisotropy profiles $\\beta(r)$, the Jeans equation can be inverted to yield the radial dispersion as a single integral over radii of $\\nu M$ times some kernel, and one can insert this solution into the equation that links the line-of-sight velocity dispersion to the radial one, to obtain after some algebra, the line-of-sight dispersion as a single integral over $\\nu M$ and some other kernel.  In what follows, we attempt to lift the mass/anisotropy degeneracy by adding a second equation, namely the 4th order Jeans equation for constant anisotropy ({\\L}okas 2002) \\[ \\frac{\\rm d}{{\\rm d} r}  (\\nu \\overline{v_r^4}) + \\frac{2 \\beta}{r} \\nu \\overline{v_r^4} + 3 \\nu \\sigma_r^2 \\frac{{\\rm d} \\Phi}{{\\rm d} r} =0 \\ , \\] where we express the line of sight \\emph{kurtosis} as a double integral over $\\nu M$ and $M$ and a third kernel (the details are given in {\\L}okas \\& Mamon 2003). ",
        "conclusions": ""
    },
    "0404/astro-ph0404505_arXiv.txt": {
        "abstract": "We have performed a systematic analysis of the dynamics of different galaxy populations in galaxy groups from the 2dFGRS. For this purpose we have combined all the groups into a single system, where velocities $v$ and radius $r$ are expressed adimensionally. We have used several methods to compare the distributions of relative velocities of galaxies with respect to the group centre for samples selected according to their spectral type (as defined by Madgwick et al., 2002), $b_j$ band luminosity and $B\\!-\\!R$ colour index.  We have found strong segregation effects: spectral type I objects show a statistically narrower velocity distribution than that of galaxies with a substantial star formation activity (type II-IV). Similarly, the same behavior is observed for galaxies with colour index $B\\!-\\!R\\!>\\!1$ compared to galaxies with $B\\!-\\!R\\!<\\!1$. Bright ($M_{b}<-19$) and faint ($M_{b}>-19$) galaxies show the same segregation. It is not important once the sample is restricted to a given spectral type.  These effects are particularly important in the central region ($R_p<0.5\\;R_{vir}$) and do not have a strong dependence on the mass of the parent group. These trends show a strong correlation between the dynamics of galaxies in groups and star formation rate reflected both by spectral type and by colour index. ",
        "introduction": "Galaxy properties can be affected by several mechanisms in groups or clusters. The fact that different galaxies can be modified to a different extent, could give rise to observable segregational effects. By studying these effects, we may obtain valuable information on the way in which these mechanisms act on galaxies and drive their evolution.  The morphology density relation \\citep{oelmer,dress80,andreon} is the best known segregational effect. Early type galaxies are more concentrated in denser regions, and lie closer to the centres of the clusters than late type galaxies. More recently, the clustering properties of galaxies have been found to be dependent on the characteristics of spectral features \\citep{julian,mardom,biviano,madg03c}, and luminosity \\citep{benoist,norberg01,norberg02,stein,adami,girardi}.  Several mechanisms have been proposed to explain galaxy transformations. Their relevance are quite different according to the environmental conditions \\citep{balogh} and so, their importance depends on the mass of the clusters, and perhaps on the history of galaxy clustering \\citep{gnedin.b}. Some effects are more effective in dense regions like rich clusters, whereas in groups of galaxies other mechanisms play the most important role.  Ram pressure \\citep{gunn-gott} can inhibit star formation by exhausting the gas present in galaxies that move fast in the intergalactic medium of rich clusters.  Similarly, galaxy harassment \\citep{moore} can produce significant changes in the star formation rate of a galaxy.  These effects are not expected to be important in poor clusters or groups, where the velocity dispersion is lower, instead, effects such as mergers or tidal interactions can be dominant in these environments \\citep{gnedin.a}.  Besides affecting galaxy properties, such as star formation, luminosity and colour index, some of the physical processes listed above may also produce changes on the dynamics of the galaxy with respect to the cluster centre.  In turn, the efficiency of some of these mechanisms to produce significant changes on a galaxy, depends on its dynamical behavior.  For example, the effects of galaxy interactions are stronger for galaxies moving slowly with respect to the cluster centre.  This suggests that the dynamical properties of galaxies in groups and clusters may be related with the star formation efficiency, colours or luminosities of galaxies\\citep{menciff,moore2}. Segregation effects of galaxy velocities in clusters are predicted theoretically \\citep{menciff,gnedin.b}, in semi--analytical models \\citep{menci}, and has been reported in rich clusters \\citep{sodre}.  The relation between the dynamical properties and the luminosity of a galaxy has been observed in rich clusters, by e.g. \\citet{whitmore,adami} and \\citet{stein}, who find evidence for velocity segregation. In agreement with these findings, theoretical studies \\citep{fyf} and numerical simulations \\citep{yepes} show similar trends.  However, it is not clear how to interpret these results. Some authors propose different orbit shapes for galaxies with different morphologies or luminosities.  In this scheme, early--type galaxies have quasi--isotropic orbits, while late--type galaxies move in nearly radial orbits \\citep{biviano77,adami}. However, other models have been proposed that contradict this statement \\citep{amelia}.  Theoretical works predict virialized systems with a Maxwellian velocity distribution \\citep{saslaw,ueda} so it has been proposed that early and late-type galaxies have Gaussian velocity distributions.  However, observations in rich clusters do not support these hypothesis (e.g. Colles \\& Dunn, 1996)  The purpose of this paper is to explore for a possible difference in the dynamical behavior of galaxies with different spectral types, luminosities and colour indexes.  The outline of this paper is as follows.  In section 2 we describe the data sample used in our work, and in section 3, the method used in our analysis.  Section 4 presents the results of our search for velocity segregation in spectral type, luminosity and colour. Finally, in section 5 we present a discussion of our results and future perspectives. ",
        "conclusions": "\\begin{figure} \\includegraphics[width=\\columnwidth]{fig8.eps} \\label{smooth} \\caption{Normalized velocity distributions convolved with a Gaussian of width $85\\;Km\\,s^{-1}$.  Inside box shows the binned case for comparison.} \\end{figure}%  We find a statistically significant difference of the distributions of group centric line of sight velocities, normalized to the group mean velocity dispersion, for samples of galaxies selected by spectral type, luminosity or colour index. Given the large data set analysed, we have been able to investigate the dependence of this velocity segregation on group properties and galaxy--group centric distance.  Spectral type I objects, corresponding to passively star forming galaxies, show a statistically narrower velocity distribution than that of galaxies with a substantial star formation activity (types III-IV). Similarly, samples of galaxies with greater colour index ($B\\!-\\!R\\!>\\!1$) have a larger fraction of small velocities ($v<1$) compared to galaxies with $B\\!-\\!R\\!<\\!1$.  These two trends show a strong correlation between galaxy dynamics in groups and star formation, reflected both by spectral type and by colour index.  The velocity distribution of luminous galaxies (typically brighter than $M_b=-19$) also show a larger fraction of small velocities, although we notice that once the galaxies are restricted to a given spectral type, there is a less significant segregation. Thus, luminosity is not likely to be a primary parameter determining galaxy dynamics in groups. Our results suggest that the observed luminosity segregation might be related to the fact that the slowest objects, of early spectral type, are on average more luminous than star forming galaxies.  There are several mechanisms that may produce dynamical segregations of galaxies in groups and clusters. Ram pressure can effectively remove the existing gas in the galaxies and transform star forming into passively star forming objects.  This mechanism affects most strongly those galaxies with large velocities with respect to the intra--cluster medium, and then it should produce a dynamical segregation with opposite trends to the observed one.  Moreover, since our analysis concerns groups and small clusters of galaxies, ram pressure is not expected to be significant.  Our results indicates that the velocity segregation effects are nearly independent of group virial mass.  This fact also suggests that ram pressure is not important since its effects are stronger in more massive systems, and with higher velocity dispersion. Thus, it is unlikely that ram pressure may explain the observed correlations.  Mergers on the other hand, are effective to generate spheroidal objects with a low star formation rate.  Galaxy encounters are expected to lower the original cluster-centric relative velocities of each galaxy, with respect to the velocities of the galaxies prior to the merger event, so that they can act effectively in generating the observed trends.  In a similar fashion, tidal interactions may effectively remove a substantial amount of gas from disks of galaxies, and then, are also effective in truncating star formation.  Early type objects generated through this mechanism, would be biased to smaller velocities since interactions are expected to be more effective in slow encounters.  Furthermore, these are generally brighter and redder objects.  It has been suggested that morphological transformation of galaxies takes place in systems which have a density threshold larger than the density of groups and poor clusters of galaxies \\citep{moore, gray}. Our results indicate that dynamical segregation of passively star forming galaxies is a generic feature of systems of galaxies, irrespective of global properties. However, the fact that segregation effectively occurs in the inner regions of groups indicates that density might be an important parameter in determining the observed effects."
    },
    "0404/astro-ph0404396_arXiv.txt": {
        "abstract": "The high eccentricities of the known extrasolar planets remain largely unexplained. We explore the possibility that eccentricities are excited in the outer parts of an extended planetary disk by encounters with stars passing at a few hundreds of AU. After the encounter, eccentricity disturbances propagate inward due to secular interactions in the disks, eventually exciting the innermost planets. We study how the inward propagation of eccentricity in planetary disks depends on the number and masses of the planets and spacing between them and on the overall surface-density distribution in the disk. The main governing factors are the large-scale surface-density distribution and the total size of the system. If the smeared-out surface density is approximated by a power-law $\\Sigma(r)\\propto r^{-q}$, then eccentricity disturbances propagate inward efficiently for flat density distributions with $q\\la 1$. If this condition is satisfied and the size of the planetary system is 50 AU or larger, the typical eccentricities excited by this mechanism by field star encounters in the solar neighborhood over 5 Gyr are in the range 0.01-0.1. Higher eccentricities ($>$0.1) may be excited in planetary systems around stars that are formed in relatively dense, long-lived open clusters. Therefore, this mechanism may provide a natural way to excite the eccentricities of extrasolar planets. ",
        "introduction": "One of the remarkable features of the $\\sim 120$ known extrasolar planets is their relatively high eccentricities\\footnote{A catalog of extrasolar planets is maintained by Paris Observatory (http://www.obspm.fr/planets)}, most far larger than seen in the giant planets of the solar system. Many mechanisms to produce these high-eccentricity orbits have been offered. The most popular involve planet-planet gravitational scattering \\citep{rasi96,weid96,lin97,ford01,yu01,chia02, terq02}, either through close encounters, physical collisions, secular interactions or resonant interactions. Other eccentricity excitation mechanisms include interactions with the gaseous protoplanetary disk \\citep{gold80,papa01,gold03,ogil03}, or interactions with the planetesimal disk \\citep{murr98}.  All of these proposed mechanisms have shortcomings. Planet-planet scattering through close encounters appears to create the wrong eccentricity distribution---there may be too many planets on low-eccentricity orbits, formed by physical collisions or tidal captures during the scattering process (\\citealt{ford01,gold03}, but see \\citealt{ford03}). Discussions of eccentricity excitation through mean-motion resonances assume that planet-disk interactions induce evolution of semimajor axes leading to resonant capture, but neglect the effects of these interactions on eccentricity. Interactions with the gaseous planetesimal disk can either damp or excite eccentricity, depending sensitively on the formation of gaps at Lindblad resonances, nonlinear saturation of corotation resonances \\citep{ogil03}, and the rate of viscous diffusion across resonances \\citep{gold03}, all of which are difficult to calculate reliably. Excitation through interactions with the planetesimal disk may require more massive disks than are indicated by observations and may only be efficient for companions with masses much higher than those of most extrasolar planets \\citep{papa01}.  There is strong and growing evidence that most forming stars are surrounded by extended disks, with radii of hundreds of AU. This evidence includes CO observations of massive gas disks around young stars, disks detected in Orion and other star-forming regions by optical emission from their photoionized surfaces, the infrared excess at wavelengths of 25--100$\\mu$ discovered by IRAS around bright nearby stars such as Vega, dust disks found optically around nearby stars such as $\\beta$ Pictoris, and imaging by the NICMOS camera on the Hubble Space Telescope of near-infrared radiation scattered by dusty disks around stars such as HR 4796 (see \\citealt{koer01} and \\citealt{zuck01} for reviews). It may well be that the dust around older stars such as Vega and $\\beta$ Pictoris is ``second-generation'' dust, released in collisions between larger objects (planetesimals or comets) that were formed by the accretion of primordial or first-generation dust particles. Thus the absence of dust around older stars may arise either because the dust has been dispersed or destroyed or because it has been incorporated into larger bodies such as planets.  Indirect evidence for planets at large distances from the host stars comes from the morphologies of the dust debris disks of Vega and $\\varepsilon$ Eri. Disk asymmetries and dust concentrations are often interpreted as resulting from the dynamical influence of an unseen massive planet. Modeling suggests that the planetary companion has a semimajor axis $a>30$ AU in the Vega system \\citep{wiln02} and about 55--65 AU in the $\\varepsilon$ Eri system \\citep{ozer00}. Formation of giant planets by core accretion is difficult at these distances, and it has been suggested that planets can migrate outward to distances of up to tens of AU through interactions with the gas disk in young systems \\citep{vera04}.  We propose to explore the possibility that the eccentricities of the extrasolar planets arise because of secular interactions in a long-lived extended planetary or protoplanetary disk. We suggest that eccentricities are excited in the outer part of the disk by a passing star and propagate inward through the disk, somewhat like a wave. We explore how the inward propagation of eccentricity in disks depends on the overall surface-density distribution.  In Section \\ref{sec_secular} we review the definitions and the approach of secular perturbation theory. In Section \\ref{sec_excit} we describe the initial excitation of the eccentricity and in Section \\ref{sec_prop} we study the propagation of eccentricities via secular interactions. We discuss typical eccentricities produced by this mechanism in Section \\ref{sec_disc} and conclude in Section \\ref{sec_conc}. ",
        "conclusions": "\\label{sec_conc}  The eccentricities of most extrasolar planets are much larger than those of giant planets in the Solar System. This observation requires explanation because planets that form from a protoplanetary disk are expected to have nearly circular orbits. In this paper we explore the possibility that close stellar encounters can excite planetary eccentricities.  Encounters with nearby stars can be dynamically important for the outer parts of circumstellar disks and planetary systems. \\citet{koba01} studied pumping up of eccentricities and inclinations of planetesimals in a close stellar encounter. \\citet{kala01} and \\citet{larw01} suggested that encounters with nearby stars can be responsible for asymmetries in the $\\beta$ Pictoris disk. \\citet{hurl02} found that most planets are stripped from their host stars by encounters with neighbor stars in the dense centers of globular clusters.  In the solar neighborhood or an in open cluster the eccentricities of short-period planets cannot be excited directly by perturbations from nearby stars to any significant values because such planets are tightly bound to their host star. We suggest here that if there are other planets or a significant mass in smaller bodies in the same system on larger orbits, the outer planets can be excited to high eccentricities and then transfer this excitation to inner planets through secular perturbations.  The extent to which the eccentricity of the innermost planets can be excited depends on the total size of the planetary system and on its smeared-out surface density distribution. If the smeared-out surface density distribution is parametrized by a power law, $\\Sigma(r)\\propto r^{-q}$, then the proposed mechanism is important for systems with $q\\la 1$ and with total size $a_N\\ga 50$ AU. At least some gas disks around young stars show flat density distributions with $q\\simeq 1$ \\citep{dutr96,wiln00} and extend out to tens or hundreds of AU. Theoretical models of steady-state protoplanetary disks also suggest flat surface density distributions \\citep{bell97}. It appears that the maximum lifetime of such gas disks is $10^6-10^7$ years \\citep{hais01}; some of this material may dissipate, but much or most of it may survive as planets or large planetesimals. Outward migration of planets can also yield massive planets at a few tens of AU \\citep{vera04}. It has been suggested that the asymmetries in the dust debris disks in some systems may be due to unseen planets at radii $30-60$ AU \\citep{ozer00, wiln02}.  If the mass distribution is flat and the planetary system is large, the inner planets can be excited to high eccentricities (0.1$-$1, eq. \\ref{eps_open}) due to encounters in rich birth clusters. The typical eccentricities that can be excited by the described mechanism from field star encounters in the solar neighborhood are smaller (0.01$-$0.1, eq. \\ref{eps_solar}). A small fraction of stars (10\\%) experience very close encounters with field stars (within 200 AU) in which excited eccentricities can be much higher.  All the calculations were done under a simplifying assumption that the perturber is moving in the same plane as the planets. \\citet{koba01} argued that the eccentricities excited by a passing star are insensitive to the inclination of the perturber relative to the planetary disk and lie in the range determined by coplanar prograde and retrograde encounters. Therefore, our conclusions about the eccentricity excitation and propagation are likely to be independent of the inclination of the perturber. In the case when the trajectory of the perturber is inclined relative to the plane of the planetary system, inclinations of the outer planets in the system can be excited and then propagate inward, similarly to propagation of eccentricities that we discussed in this paper. Warps in circumstellar disks like those seen in $\\beta$ Pictoris \\citep{kala95, wahh03} may be signatures of inclination waves propagating through the system after the initial perturbation.  In this work our focus was to find the conditions for efficient propagation of eccentricity disturbances. In our simulations, the masses and semimajor axes of successive planets follow geometric progressions. Clearly, these special configurations do not represent the full range of properties of real systems. Therefore, we do not discuss the distribution of eccentricities resulting from our simulations in relation to the observed distribution.  \\bigskip We would like to thank Marc Kuchner and Roman Rafikov for useful discussions and the referee for helpful comments on the manuscript. ST acknowledges the support of NASA grants NAG5-10456 and NNG04GH44G.  \\newpage"
    },
    "0404/astro-ph0404169_arXiv.txt": {
        "abstract": "\\citet[LBK]{Lynden-Bell.Kalnajs:72} presented a useful formula for computing the long-range torque between spiral arms and the disk at large. The derivation uses second-order perturbation theory and assumes that the perturbation slowly grows over a very long time (the {\\em time-asymptotic limit}).  This formula has been widely used to predict the angular momentum transport between spiral arms and stellar bars between disks and dark-matter halos.  However, this paper shows that the LBK time-asymptotic limit is {\\em not} appropriate because the characteristic evolution time for galaxies is too close to the relevant dynamical times.  We demonstrate that transients, not present in the time-asymptotic formula, can play a major role in the evolution for realistic astronomical time scales. A generalisation for arbitrary time dependence is presented and illustrated by the bar--halo and satellite--halo interaction.  The natural time dependence in bar-driven halo evolution causes quantitative differences in the overall torque and qualitative differences in the physical- and phase-space location of angular momentum transfer.  The time-dependent theory predicts that four principal resonances dominate the torque at different times and accurately predicts the results of an N-body simulation.  In addition, we show that the Inner Lindblad Resonance (ILR) is responsible for the peak angular momentum exchange but, due to the time dependence, the changes occur over a broad range of energies, radii and frequencies.  We describe the implication of these findings for the satellite--halo interaction using a simple model and end with a discussion of possible impact on other aspects secular galaxy evolution. ",
        "introduction": "\\label{sec:intro}  A near-equilibrium galaxy only evolves through the excitation of some non-axisymmetric structure such as spiral arms, bars, infalling satellites, etc.  The resulting ``waves'' are likely actors in triggering star formation, mediating inward gas flow thereby fueling AGN, heating the stellar disk, among other observable phenomena.  In addition, the overall evolution of near equilibrium galaxies is caused by the collective transport of energy and angular momentum from these same structures.  An accurate theory of secular evolution is essential for developing our understanding of these interrelated dynamics.  The first step is predicting the evolution of the non-axisymmetric distortions.  The underlying dynamical mechanism for the stellar and dark matter components is understood.  At the level of individual orbits, an excitation such as a bar, a spiral pattern, or orbiting satellite produces a periodic distortion on all orbits, much like a driven harmonic oscillator.  Although the torque distorts a typical orbit, this distortion averages to zero over many orbital periods for most orbits.  However, for nearly closed orbits\\footnote{A closed orbit is one with integral commensurabilities between the orbital frequencies and the pattern frequency.}, the net torque does not vanish.  An adiabatic invariant is broken by the vanishing precession frequency of the closed orbit; clearly, there is no adiabatic invariant for a degree of freedom with no motion.  The perturbation then couples to the oscillatory distortion for these orbits and net torque is transferred to or from these orbits.  This coupling between the forced oscillation and the perturbation will be second order in the perturbation strength. In time, the perturbation itself and the underlying equilibrium is slowly changed by this torque.  This leads to a finite measure of orbits that transfer angular momentum to and from the pattern \\citep[for additional discussion, see][hereafter Paper 2]{Weinberg.Katz:04}.  Although the first-order perturbation may dominate the instantaneous changes to phase space, the second-order changes describe the net changes that will persist after many dynamical times.  This process was described mathematically by \\citet[hereafter LBK]{Lynden-Bell.Kalnajs:72} in the limit that the change in the perturbation is very slow but still fast enough that the orbital perturbations remain linear.  LBK derived a formula describing the exchange of angular momentum between a spiral pattern and the rest of the disk assuming that the perturbation began infinitely long in the past, the time-asymptotic limit.  In this limit, any vestigial response from the formation of the pattern is gone.  The LBK approach has been widely applied to estimate secular evolution \\citep[e.g.][]{Goldreich.Tremaine:79b,Weinberg:85,Zhang:98,Athanassoula:03}. \\cite{Carlberg.Sellwood:85} took the first step towards a time-dependent generalisation of the LBK formalism for studying secular disk evolution.  They included time-dependent transients but evaluated the secular changes in the infinite time limit.  This is appropriate for studying the cumulative effect of short-lived transients.  However, the perturbation theory may be fully generalised to treat arbitrary time-dependent perturbations over arbitrary intervals of time.  This allows us to treat the long-lived bar and satellite interactions over the full age of a galaxy with some surprising results.  In attempting a detailed description of the angular momentum exchange between a rotating bar and a dark halo in an N-body simulation (see Paper 2), I found a significant discrepancy between the predictions of the second-order LBK perturbation theory and the N-body simulation. This discrepancy persisted for perturbations over a wide range of bar amplitudes and scaled with time and perturbation amplitude as expected from second-order linear theory, implying that its source is not a break down in linearity.  Rather, the problem is due to the assumption of an infinitely slow growth of the perturbation.  Not only is the number of characteristic dynamical times in a galactic age modest, the growth of a bar, arms and of course satellites is most likely a small fraction of the galactic age.  Because patterns often appear over several orbital periods and the total number of orbital periods available are small, transients may be significant and the time-asymptotic limit does not apply.  I will show in this paper that a finite-time-limit generalisation of the LBK formula gives quantitatively and qualitatively different results.  Quantitatively, we will see that the overall torque is smaller than that computed from the LBK formula for a rotating bar.  Qualitatively, different commensurabilities are important at different epochs of the evolution. Because the location of the resonances in physical space are important for the long-term evolution of galaxy, including the finite-time response is important for understanding this evolution.  In this limit, the time dependence is more than simply a transient but embodies the galaxy's evolutionary history.  I will present a new, generalised secular evolution formula and describe the important details of the dynamics in \\S\\ref{sec:method}.  The new torque formula replaces the delta function in the LBK formula by an integral over the time-dependent perturbation and is similarly straightforward to apply.  Two specific applications are presented in \\S\\ref{sec:examples}: bar--dark halo coupling and satellite--dark halo coupling.  N-body simulations will demonstrate the importance of the finite-time response to the net torque and illustrate the discrepancy with the LBK formula.  In \\S\\ref{sec:barslow}, we will see that the ILR dominates the bar evolution and use this new formalism to explicitly predict the location of the angular momentum deposited. The same dynamics is then applied to sinking satellites in \\S\\ref{sec:sinksat}.  We will conclude with \\S\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  This paper shows the \\citet[LBK]{Lynden-Bell.Kalnajs:72} secular torque formula is not a quantitatively accurate description of secular evolution in galaxies.  The LBK formula assumes that the growth time of a perturbation driving secular evolution is infinitely long compared the characteristic dynamical time.  However, not only is the number of characteristic dynamical times in a galactic age modest, the growth of a spiral arms, a bar or the decay of a large dwarf satellite has evolutionary time scales of order 100 Myr to 1 Gyr, uncomfortably close to the equilibrium galaxy lifetime of 10 Gyr or smaller.  We have seen that the finite time effects yield a different distribution of angular momentum in the halo and a different total torque than for a time-asymptotic system assumed by LBK.  Similarly, it is unlikely that time-dependent effects will affect time-asymptotic perturbation theory predictions in planetary systems \\cite{Goldreich.Tremaine:79b, Goldreich.Tremaine:80} where the secular evolution time scale is much larger than the dynamical time scale.  In \\S\\ref{sec:method}, I presented generalised the LBK formula for perturbations of finite duration.  The new formula shares many of the features of the LBK formula; the main difference is that the Dirac delta function $\\delta(\\ldo - m\\Omega_p)$ is replaced with a time integral (see eq. \\ref{eq:LBK1}).  Over astronomically-relevant time scales, the finite-time generalisation shows that transients from the formation history of the system play an important role in the overall dynamical evolution.  For example, Figures \\ref{fig:pattern} and \\ref{fig:lbk} illustrate the total torque and differential contributions of the principal resonances during slow down of a bar in dark matter halo.  The correct result is the net contribution for approximately 8 resonances (not all shown in Fig. \\ref{fig:lbk}) and therefore these resonances must be accurately represented in an N-body simulation to obtain the correct result (see Paper 2).  The most important general finding is that the history of galaxy evolution can not be ignored in understanding and predicting a particular galaxy's evolution.  These dynamics are illustrated in two cases: the bar--halo interaction and the satellite--halo interaction. The details of the bar--halo interaction have been recently described in Paper 2 and explored in a self-consistent simulation in Paper 3 \\citep{Holley-Bockelmann.etal:04}.  Both papers demonstrate that bar--halo coupling through the ILR agrees with the predictions presented here.  Both the LBK formula and the time-dependent formula predict that the ILR resonance dominates the overall angular momentum transfer. However, the time-dependent evolution spreads the resonant interaction over a broader range of lower energies that are more populated in phase space and therefore have a larger overall evolutionary consequence.  The time-asymptotic ILR is located at smaller energies and radii and would be inaccessible to most N-body simulations.  The time-asymptotic approximation is especially poor for the ILR resonance in a cuspy halo where the time-dependent spread in frequency corresponds to a large region in radius and energy.  This is strong encouragement to attempt an understanding of the dynamical interactions prior to simulation.  One expects a similar halo response in the satellite--halo interaction because the quadrupole perturbation from a rotating bar shares much in common with the quadrupole perturbation from an orbiting satellite.  Indeed, we show that this is the case.  However, because the symmetries, the radial profile, and the distribution of frequencies for the two perturbations are different, the ILR no longer dominates in the satellite interaction. Nonetheless, the ILR does play a role during the early stages of orbital evolution and this has been verified in N-body simulation \\citep[see][]{Choi.etal:04}.  The aggregate effect of this mechanism from a population of substructure will be the subject of a later investigation.  The wakes excited by interactions between satellites, disks and dark matter halos give rise to a variety of observable consequences.  For example, large-scale transport of angular momentum throughout the disk and to the halo may will tend to decrease the disk scale length.  The response of the galaxy halo to the surrounding group environment can propagate features to the inner galaxy, exciting disk waves and subsequent star formation.  In addition, the interaction between the satellite and disk will drive bending modes and the energy in these modes will heat the disk. Detailed simulations for these processes of disparate scales are difficult.  Time-asymptotic secular perturbation theory has often been used to estimate evolution and compare with simulation.  However, we have found that the difference in amplitude and significance between the time-dependent and time-asymptotic responses for the bar--halo interaction (Fig. \\ref{fig:lbk}) is dramatic.  If these differences are representative of the difference in these other processes, our current picture of secular evolution must be revisited by considering the time-dependent evolution of these ``waves'' in the context of a galaxy's environmental history."
    },
    "0404/astro-ph0404443_arXiv.txt": {
        "abstract": "{  We present 2D Monte Carlo radiative transfer simulations of prestellar cores. We consider two types of asymmetry: {\\it disk-like} asymmetry, in which the core is denser towards the equatorial plane than towards the poles; and {\\it axial} asymmetry, in which the core is denser towards the south pole than the north pole. In both cases the degree of asymmetry is characterized by the ratio $e$ between the maximum optical depth from the centre of the core to its surface and the minimum optical depth from the centre of the core to its surface. We limit our treatment here to mild asymmetries with $e = 1.5\\,$ and $2.5\\,$. We consider both cores which are exposed directly to the interstellar radiation field and cores which are embedded inside molecular clouds.  \\hspace{0.3cm} The SED of a core is essentially independent of the viewing angle, as long as the core is optically thin. However, the isophotal maps depend strongly on the viewing angle. Maps at wavelengths longer than the peak of the SED (e.g. 850 $\\mu$m) essentially trace the column-density. This is because at long wavelengths the emissivity is only weakly dependent on temperature, and the range of temperature in a core is small (typically $T_{\\rm max}/T_{\\rm min} \\la 2$). Thus, for instance, cores with disk-like asymmetry appear elongated when mapped at 850 $\\mu$m from close to the equatorial plane. However, at wavelengths near the peak of the SED (e.g. 200 $\\mu$m), the emissivity is more strongly dependent on the temperature, and therefore, at particular viewing angles, there are characteristic features which reflect a more complicated convolution of the density and temperature fields within the core.  \\hspace{0.3cm} These characteristic features are on scales $1/5$ to $1/3$ of the overall core size, and so high resolution observations are needed to observe them. They are also weaker if the core is embedded in a molecular cloud (because the range of temperature within the core is then smaller), and so high sensitivity is needed to detect them. {\\it Herschel}, to be launched in 2007, will in principle provide the necessary resolution and sensitivity at 170 to 250 $\\mu$m. ",
        "introduction": "\\hspace{1.2em} Prestellar cores are condensations in molecular clouds that are either on the verge of collapse or already collapsing (e.g. Myers \\& Benson 1983; Ward-Thompson, Andr\\'e \\& Kirk 2002). They represent the initial stage of star formation and their study is important because theoretical models of protostellar collapse suggest that the outcome is very sensitive to the initial conditions. Prestellar cores have been observed both in isolation and  in protoclusters. Isolated prestellar cores (e.g. L1544, L43 and L63; Ward-Thompson, Motte \\& Andr\\'e 1999) have extents $\\ga 1.5\\times 10^4$ AU and  masses $0.5-35~{\\rm M}_{\\sun}$ (see also Andr\\'{e}, Ward-Thompson, \\& Barsony 2000). On the other hand, prestellar cores in protoclusters (e.g. in $\\rho$ Oph and NGC2068/2071) are generally smaller, with extents $\\sim 2-4 \\times 10^3$ AU and masses $\\sim 0.05-3~{\\rm M}_{\\sun}$ (Motte, Andr\\'e  \\& Neri 1998, Motte et al. 2001).  Many authors have modelled prestellar cores with Bonnor-Ebert (BE) spheres, i.e. equilibrium isothermal spheres in which self-gravity is balanced by gas pressure (Ebert 1955, Bonnor 1956). For example, Barnard 68 has been modelled in this way by Alves, Lada \\& Lada (2001).  However, it is evident from 850 $\\mu$m continuum maps of prestellar cores, which essentially trace the column-density through a core, that prestellar cores are not usually spherically symmetric (e.g. Motte et al., 1998; Ward-Thompson et al., 1999; Kirk et al., 2004). Indeed, statistical analyses of the projected shapes of a large sample of cores (Jijina, Myers \\& Adams, 1999) suggest that prestellar cores do not even have spheroidal symmetry and are better represented by triaxial ellipsoids (Jones, Basu \\& Dubinski, 2001; Goodwin, Ward-Thompson \\& Whitworth, 2002).  This is not surprising, given the highly turbulent nature of star-forming molecular clouds, and the short time-scale on which star formation occurs (e.g. Elmegreen, 2000). Normally, prestellar cores are formed -- and then either collapse or disperse -- so rapidly that they do not have time to relax towards equilibrium structures. Even in more quiescent environments where cores can evolve quasi-statically, the combination of magnetic and rotational stresses is likely to produce significant departures from spherical symmetry.  For instance, SPH simulations of isothermal, turbulent, molecular clouds by Ballesteros-Paredes, Klessen \\& Vazquez-Semadeni (2003) show (a) that most of the cores that form are transient and non-spherical; and (b) that, despite this fact, the column density can, more often than not, be adequately fitted with the Bonnor-Ebert profile, although the parameters of the fit depend on the observer's viewing angle. Likewise, the FD simulations of magnetic, isothermal, turbulent, molecular clouds reported by Gammie et al. (2003) produce transient, triaxial cores.  Similarly, evolutionary models of individual prestellar cores predict the formation of flattened (oblate spheroidal) structures, either due to rotation (e.g. Matsumoto, Hanawa \\& Nakamura 1997), or due to flattening along the lines of a bipolar magnetic field (e.g. Ciolek \\& Mouschovias, 1994). Other models invoke a toroidal magnetic field to create prolate equilibrium cores (e.g. Fiege \\& Pudritz, 2000). Triaxial structures can be generated with a suitable combination of rotation and magnetic field.  It is therefore important to investigate, by means of radiative transfer modelling, how the intrinsic asymmetries inherent in the formation and evolution of a core might translate into observable asymmetries on continuum maps of cores. Previous continuum radiative transfer modelling of prestellar cores has examined non-embedded BE spheres (Evans et al., 2001, Young et al., 2003) and embedded BE spheres (Stamatellos \\& Whitworth, 2003a), using 1D (spherically-symmetric) codes. Zucconi, Walmsley \\& Galli (2001) have used an approximate, semi-analytic method to model non-embedded, magnetically flattened prestellar cores, in 2D. We have recently developed a Monte Carlo code for modelling continuum radiative transfer in arbitrary geometry, and with arbitrary accuracy. Preliminary results have been presented in Stamatellos \\& Whitworth (2003b,c).  In this paper we develop continuum radiative transfer models of non-spherical cores. Since these models are intended to be exploratory, rather than definitive, we consider here only 2D models (i.e. we impose azimuthal symmetry so that in spherical polar co-ordinates $(r,\\theta,\\phi)$ there is no dependence on $\\phi$).   Since star formation is a chaotic process, it is not sensible to appeal to numerical simulations for the detailed density field in a prestellar core. From both observations (e.g. Kirk et al., 2004) and simulations (e.g. Goodwin et al., 2004), it is clear that each core has a unique distribution of gas in its outer envelope, and a unique radiation field incident on its boundary. Even if existing simulations are a good representation of the real dynamics of star formation, they are not presently able to reproduce particular sources, and therefore they can only be compared realistically with observations in a statistical sense.  We will be presenting SEDs and isophotal maps for prestellar cores formed in SPH simulations of star formation in turbulent molecular clouds in a subsequent paper (Stamatellos, Goodwin \\& Whitworth, in preparation). However, for interpreting observations of individual cores, it is more appropriate to generate SEDs and isophotal maps using simply parametrized models which capture generically the different features we might hope to detect.  In this regard, we have been guided by the observations, which indicate that prestellar cores have approximately uniform density in their central regions, and the density then falls off in the envelope. If the density in the envelope is fitted with a power law, $n(r) \\propto r^{-\\eta}\\,$, then $\\eta \\sim 2\\,-\\,4 \\,$. Here $\\eta \\sim 2$ is characteristic of more extended prestellar cores in dispersed star formation regions (e.g. L1544, L63 and L43), whereas $\\eta \\sim 4$ is characteristic of more compact cores in protoclusters (e.g. $\\rho$ Oph and NGC2068/2071).  These features are conveniently represented by a Plummer-like density profile (Plummer, 1915), \\begin{equation} n(r) = n_0\\,\\frac{1}{ \\left[ 1 + \\left( \\frac{r}{r_0} \\right)^2 \\right]^{\\eta/2} } \\,, \\end{equation} where $n_0$ is the density at the centre of the core, and $r_0$ is the extent of the region in which the density is approximately uniform.  The Plummer-like density profile is ad hoc, but given the transient, non-hydrostatic nature of prestellar cores, and the coarseness of the observational constraints, this is unavoidable. It has the advantage of being simple, with only three free parameters. Uniquely amongst analytic models, it predicts lifetimes, accretion rates, collapse velocity fields, SEDs and isophotal maps which agree well with observation (Whitworth \\& Ward-Thompson, 2001; Young et al., 2003). It also reproduces approximately the BE density profile, and the density profiles predicted by the ambipolar diffusion models of Ciolek \\& Mouschovias (1994) and Ciolek \\& Basu (2000). Tafalla et al. (2004) use a similar density profile to model the starless cores  L1498 and L1517B in Taurus-Auriga.  Furthermore, the Plummer-like profile can be modified easily \tto include azimuthally symmetric departures from spherical symmetry. We treat two types of asymmetry. In the first type (disk-like asymmetry), we construct flattened cores, using density profiles of the form \\begin{equation} n(r,\\theta) = n_0\\,\\frac{1 + A \\left( \\frac{r}{r_0} \\right)^2 {\\rm sin}^p(\\theta) }{ \\left[ 1 + \\left( \\frac{r}{r_0} \\right)^2 \\right]^{(\\eta+2)/2} } \\,. \\end{equation} The parameter $A$ determines the equatorial-to-polar optical depth ratio $e\\,$, i.e. the maximum optical depth from the centre to the surface of the core (which occurs at $\\theta = 90\\degr$), divided by the minimum optical depth from the centre to the surface of the core (which occurs at $\\theta = 0\\degr$ and $\\theta = 180\\degr$). The parameter $p$ determines how rapidly the optical depth from the centre to the surface rises with increasing $\\theta$, i.e. going from the north pole at $\\theta = 0\\degr$ to the equator at $\\theta = 90\\degr$.  In the second type (axial asymmetry), we construct cores which are denser towards the south pole ($\\theta = 180\\degr$) than the north pole ($\\theta = 0\\degr$), using density profiles of the form \\begin{equation} n(r,\\theta) = n_0\\,\\frac{1 + A \\left( \\frac{r}{r_0} \\right)^2 {\\rm sin}^p(\\theta/2) }{ \\left[ 1 + \\left( \\frac{r}{r_0} \\right)^2 \\right]^{(\\eta+2)/2} } \\,. \\end{equation} Here the parameter $A$ determines the south-to-north pole optical depth ratio $e\\,$, i.e. the maximum optical depth from the centre to the surface of the core (which now occurs at $\\theta = 180\\degr$), divided by the minimum optical depth from the centre to the surface of the core (which still occurs at $\\theta = 0\\degr$). The parameter $p$ again determines how rapidly the optical depth from the centre to the surface rises with increasing $\\theta$, but now going from the north pole at $\\theta = 0\\degr$ all the way round to the south pole at $\\theta = 180\\degr$.  In all models the core has a spherical boundary at radius $R_{\\rm core}$.  For the purpose of this paper, and in order to isolate a manageable parameter space, we fix $n_0 = 10^6\\, {\\rm cm}^{-3}$, $r_0 = 2 \\times 10^3\\,{\\rm AU}$, $\\eta = 2$, and $R_{\\rm core} = 2 \\times 10^4\\,{\\rm AU}$. These are typical values for isolated cores. We can then explore the effect of varying $A$ and $p$, or equivalently $e$ and $p$.  In Section 2, we outline the basic principles underlying our Monte Carlo radiative transfer code. In Section 3, we present results obtained for cores having disk-like asymmetry; we treat both non-embedded cores and cores embedded in molecular clouds. In Section 4, we present results for cores having axial asymmetry. In Section 5, we summarize our results. ",
        "conclusions": "\\label{sec:asym.discussion}  We have performed accurate two dimensional continuum radiative transfer calculations for non-spherical prestellar cores. We argue that such non-spherical models are needed because observed cores are clearly not spherically symmetric, and are not expected to be spherically symmetric. Our models illustrate the characteristic features on isophotal maps which can help to constrain the intrinsic density and temperature fields within observed non-spherical cores. They demonstrate the importance of observing cores at wavelengths around the peak of the SED and with high resolution.  Our main results are:  \\begin{itemize}  \\item For the cores treated here, which are optically thin at the long wavelengths where most of the emission occurs, the SED is essentially the same at any viewing angle. For example, in the case of cores having disk-like asymmetry, the SED does not distinguish between cores viewed from different angles.  \\item Isophotal maps at submm wavelengths (e.g. at 850~$\\micron$) are essentially column density tracers, whereas maps at far infrared wavelengths around the peak of the core emission (e.g. 200~$\\micron$) reflect both the column density and the temperature field along the line of sight. Therefore, sensitive,  high-resolution observations at 170-250~$\\micron$ ({\\it Herschel}) combined with long-wavelength observations (e.g. 850~$\\micron$ or $1.3~{\\rm mm}$) can in principle be used to constrain the orientation of a core and the temperature field within it.  \\item If we assume a universal ISRF, then cores embedded in ambient molecular clouds are colder than cores directly exposed to the ISRF and have lower temperature gradients within them. As a result, the characteristic features on 200~$\\micron$ isophotal maps are weaker for cores embedded in ambient clouds with visual extinction less than $A_{\\rm V}~\\sim~10$, and may even disappear completely for more deeply embedded cores ($A_{\\rm V}>20$).  \\item If the ISRF incident on the ambient molecular cloud is enhanced in the UV region (for example, by nearby luminous stars) but still isotropic, then the embedded cores are hotter and the temperature gradients inside them may be sufficient to produce detectable characteristic features, even in deeply embedded cores.  \\item The shapes of asymmetric cores depend strongly on the observer's viewing angle. For example, cores with disk-like asymmetry appear more flattened when viewed edge-on. Our models also indicate that such cores should be easier to detect when viewed from near the equatorial plane. This may introduce a selection effect that should be taken into account when studying the statistics of the shapes of cores, using solely optically thin continuum, or optically thin molecular-line, observations.  \\item If the ambient ISRF is isotropic, the characteristic features on 200~$\\micron$ maps are symmetric with respect to two axes for cores with disk-like asymmetry, and with respect to one axis for cores with axial asymmetry. Thus, just the symmetry of these features, could be indicative of the core's internal density structure. Lack of symmetry in the features could indicate triaxiality, but it could also simply indicate that the radiation field incident on the core is anisotropic, due to discrete local sources.  \\end{itemize}  Recently Gon\\c{c}alves, Galli, \\& Walmsley  (2004) have presented radiative transfer models of axisymmetric and non-axisymmetric toroidal cores, and also models of cores heated by an additional external stellar source. In their approach, they use a  similar but independent Monte Carlo radiative transfer code. Their models focus on the effect of different core density profiles and anisotropic heating of cores, whereas we focus on the effect of observing mildly asymmetric cores at different viewing angles. Both studies are in general agreement for similar core models, and are helpful for the study of non-spherical realistic prestellar cores.     We are now extending our study by treating the effect of an anisotropic illuminating radiation field. We are also studying the triaxial molecular cores that result from 3-dimensional hydrodynamic simulations."
    },
    "0404/astro-ph0404325_arXiv.txt": {
        "abstract": "{\\hh\\ is an extremely hot hydrogen-deficient white dwarf with an effective temperature close to 200\\,000\\,K. We present new FUV and soft X-ray spectra obtained with FUSE and Chandra, which confirm that \\hh\\ has an atmosphere primarily composed of carbon and oxygen. The Chandra LETG spectrum (60\\,--\\,160\\AA) shows a wealth of photospheric absorption lines from highly ionized oxygen, neon, and -- for the first time identified in this star -- magnesium and suggests relatively high Ne and Mg abundances. This corroborates an earlier suggestion that \\hh\\ represents a naked C/O stellar core or even the C/O envelope of an O-Ne-Mg white dwarf. ",
        "introduction": "\\hh\\ is a faint blue star that has been identified as the counterpart of a bright soft X-ray source (Nousek \\etal 1986)  discovered by an early X-ray survey (Nugent \\etal 1983). Spectroscopically, the star is a member of the PG1159 class which comprises hot hydrogen-deficient (pre-) white dwarf stars ($T_{\\rm eff}$=75\\,000\\,K--180\\,000\\,K, $\\log g$=5.5--8 [cgs]; Werner 2001). The PG1159 stars are probably the outcome of a late helium-shell flash, a phenomenon that drives the currently observed fast evolutionary rates of three well-known objects (FG~Sge, Sakurai's object, V605 Aql). A late helium-shell flash may occur in a post-AGB star or a white dwarf. Flash induced envelope mixing generates a H-deficient surface layer. The surface chemistry then essentially reflects that of the region between the H-and He-burning shells of the precursor AGB star. The He-shell flash transforms the star back to an 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.  Within the PG1159 group \\hh\\ is an extraordinary object, as it has been shown that it is not only hydrogen-deficient but also helium-deficient.  From optical spectra it was concluded that the atmosphere is primarily composed of carbon and oxygen, by equal amounts (Werner 1991, {\\small W91}). Strong neon lines were detected in soft X-ray spectra taken with the EUVE satellite and in an optical-UV Keck spectrum and an abundance of Ne=2--5\\% (all abundances in this paper are given as mass fractions) was derived (Werner \\& Wolff 1999, {\\small WW99}). The origin of this exotic surface chemistry is unclear.  Here we present new results of observations performed with FUSE (Far Ultraviolet Spectroscopic Explorer) and the Chandra X-ray Observatory, whose spectroscopic resolution is more than an order of magnitude better than that of previous FUV and EUV missions (HUT and EUVE). We will show that the Chandra spectrum contains a wealth of photospheric absorption lines from highly ionized metals. Ionization balances of O, Ne, and Mg lines provide new constraints on the effective temperature and, for the first time, allow an estimation of the Mg abundance. We also analyze the spectra with respect to other metals, namely Al, Na, and the iron group.  In the following sections we will first describe the model atmosphere calculations. Then we will present in detail the compilation of atomic data and design of the model atoms. After a coarse characterization of the FUSE and Chandra spectra we turn to the detailed comparison between synthetic and observed spectra. Finally we discuss the results and their implications on the evolutionary history of \\hh.  \\begin{figure}[tbp] \\resizebox{\\hsize}{!}{\\includegraphics{154fig01.ps}} \\caption[]{ Chandra observation (thin line) compared to two model spectra (thick lines). From one of the models we have omitted the bound-free cross-section of the level \\ion{Mg}{v} 2p$^5$\\,$^3$P$^{\\rm o}$ in order to demonstrate that a strong autoionization feature in the cross-section (dashed line) can result in a line-like absorption feature in the calculated spectrum (marked by the arrow, \\Teff=175\\,000\\,K). } \\label{autoionplot} \\end{figure} ",
        "conclusions": "Let us summarize the properties of \\hh: \\begin{eqnarray*} \\Teff &=& 200\\,000\\,{\\rm K} \\pm 20\\,000\\,{\\rm K}\\\\ \\logg &=& 8.0 \\pm 0.5 {\\rm \\ \\ [cgs]} \\end{eqnarray*} Element abundances in \\% mass fraction: \\begin{eqnarray*} {\\rm C}&=&48 \\\\ {\\rm O}&=&48 \\\\ {\\rm Ne}&=&2 \\\\ {\\rm Mg}&=&2 \\\\ {\\rm Fe-group}&=&0.14\\ \\ {\\rm (solar)} \\\\ {\\rm He}&<&1 \\\\ {\\rm Na}&<&0.1 \\\\ {\\rm Al}&<&0.1 \\end{eqnarray*} The value of \\logg\\ and abundances for C, O, and Ne were taken from previous work ({\\small W91, WW99}). Estimated errors for abundances are: $\\pm$20\\% for mass fraction of C and O, and a factor of 3 for Ne, Mg, and Fe-group.  Stellar mass and luminosity can be derived by comparing the position of \\hh\\ in the $g$-\\Teff\\ diagram with theoretical evolutionary tracks. We use the post-AGB tracks of Bl\\\"ocker (1995) and derive: \\begin{eqnarray*} M/{\\rm M}_\\odot&=&0.836^{+0.13}_{-0.10}\\\\ \\log L/{\\rm L}_\\odot&=&2.45^{+0.6}_{-0.4}\\\\ d/{\\rm kpc}&=&0.67^{+0.3}_{-0.53} \\end{eqnarray*} Note that the mass of \\hh\\ is considerably higher than the mean mass of the PG1159 stars (0.6\\,M$_\\odot$). The spectroscopic distance was obtained by comparing the measured visual flux (V=16.24, Nousek \\etal 1986) with the flux of the final model (\\Teff=200\\,000\\,K, \\logg\\,=\\,8): H$_\\nu$[5400\\AA]=$3.42 \\cdot 10^{-3}$ erg/cm$^2$/s/Hz). Interstellar reddening was neglected for this determination, because it is very low. In fact, the best model fit to the continuum shape of the UV/FUV spectrum taken with HUT provided E(B-V)=0 (Kruk \\& Werner 1998).  \\begin{figure*}[tbp] \\resizebox{\\hsize}{!}{\\includegraphics{154fig10.ps}} \\caption[]{ Synthetic spectra (\\Teff=200\\,000\\,K) with \\ion{Fe}{ix} lines (top panel) and \\ion{Ni}{ix} lines (bottom panel) compared to the Chandra spectrum. The POS line lists are used here to enable line identifications. Identifying individual lines is impossible, probably due to strong line blending by other heavy metal lines, as demonstrated in Fig.\\,\\ref{chandra_fe_group_a}. Observed count spectrum and model spectra (relative flux) are convolved with Gaussians with FWHM=0.02\\AA\\ and 0.03\\AA, respectively. } \\label{chandra_fe_group_b} \\end{figure*}    We have analyzed new FUV and soft X-ray spectra of the unique object \\hh. We confirm its exotic chemical composition, which is dominated by C and O. We confirm the high Ne abundance and find a similarly high abundance of Mg. This chemistry either reflects that of the core of a C/O white dwarf or the C/O envelope of a white dwarf with a O-Ne-Mg core. It therefore remains unclear if \\hh\\ has gone beyond 3$\\alpha$ burning through a subsequent C burning phase or not. In any case, the origin of the He-deficiency remains obscure. \\hh\\ could be an extreme PG1159 star which -- in contrast to the other stars of this group -- has for some unknown reason burned up its helium completely. Alternatively, \\hh\\ could have burned carbon, now being a O-Ne-Mg white dwarf. Some unidentified mechanism (C shell flashes?) may be responsible for the loss of helium by ingestion and burning in deep hot layers. Interestingly, Iben and collaborators have predicted that such C-burning stars in the super-AGB phase could loose their H- and He-rich envelopes by a radiation driven superwind (Ritossa \\etal 1996). \\hh\\ might resemble the result of such a scenario.  A surprisingly rich photospheric absorption line spectrum in the soft X-ray regime has been revealed by our Chandra observation. Although the overall flux distribution cannot be explained by a single model with a particular temperature, the ionization equilibria of O, Ne, and Mg suggest that \\Teff\\ is slightly higher than determined in previous analyses (200\\,000\\,K$\\pm$20\\,000\\,K). This makes \\hh\\ the hottest known post-AGB star and white dwarf ever analyzed in detail with model atmosphere techniques."
    },
    "0404/astro-ph0404113_arXiv.txt": {
        "abstract": "We report on continued monitoring of the Anomalous X-ray pulsar (AXP) \\tfe\\ using the \\textit{Rossi X-ray Timing Explorer}.  We confirm that this pulsar has exhibited significant pulsed flux variability.  The principal features of this variability are two pulsed X-ray flares.  Both flares lasted several months and had well-resolved few-week--long rises. The long rise times of the flares are a phenomenon not previously reported for this class of object. The epochs of the flare peaks were MJD $52,218.8 \\pm 4.5$ and $52,444.4 \\pm 7.0$.  Both flares had shorter rise than fall times.  The flares had peak fluxes of  $2.21 \\pm 0.16$ and $3.00 \\pm 0.13$ times the quiescent value.  We estimate a total 2--10~keV energy release of  $\\sim 2.7\\times 10^{40}$~ergs and $\\sim 2.8\\times 10^{41}$ ergs for the flares, assuming a distance of 5 kpc. We also report large (factor of $\\sim$12) changes to the pulsar's spin-down rate on time scales of weeks to months, shorter than has been reported previously.  We find marginal evidence for correlation between the flux and spin-down rate variability, with probability of nonrandom correlation 6\\%.  We discuss the implications of our findings for AXP models. ",
        "introduction": "\\label{sec:intro}  Anomalous X-ray pulsars (AXPs) are an exotic manifestation of young neutron stars.  AXPs are known for their steady, soft X-ray pulsations in the period range of 6--12~s.  The detection of X-ray bursts from two AXPs has confirmed the common nature of these objects with that of soft gamma repeaters \\citep{SGRs; gkw02,kgw+03}, another exotic type of young neutron star.  Both classes of objects are believed to be magnetars, i.e., powered by the decay of an ultrahigh magnetic field that has a magnitude of $10^{14}$--$10^{15}$ G on the stellar surface.  For recent AXP reviews, see \\citet{kg04a} and \\citet{kas04}.  One issue in AXP research has been flux stability.  Historically, two AXPs have been reported to be highly flux variable.  \\citet{opmi98} collected all published flux measurements for AXP \\tfe\\ and concluded that its total flux varies by as much as a factor of 10 between observations spaced by typically 1--2 yr over $\\sim$20 yr.  Those data were from a diverse set of instruments, including imaging and nonimaging telescopes.  Similarly, flux variability by a factor of greater than 4 was reported for AXP \\tfn\\ by \\citet{bs96}, using data also from a variety of instruments.  However, long-term {\\it Rossi X-Ray Timing Explorer (RXTE)} monitoring of the pulsed flux  of \\tfe\\ by \\citet{kgc+01} and of \\tfn\\ by \\citet{gk02} using a single instrument and set of analysis software showed no evidence to support such large variability.\\footnote{Total flux measurements with \\rxte\\ were difficult given the large field of view of the PCA and  the low count rates for the AXPs relative to the background.} Also, \\citet{tgsm02}, following a short {\\it XMM-Newton} observation of \\tfe, compared the observed flux with those measured by two other imaging instruments, {\\it ASCA} and {\\it BeppoSAX}.  They found that in the three observations, the total flux was steady to within $\\sim$30\\%--50\\%.  They argued that the nonimaging detections included in the \\citet{opmi98} analysis may have been contaminated by other sources in the instruments' fields of view; in particular, the bright and variable X-ray source $\\eta$ Carina lies only 38$'$ away.  A possible solution to this puzzle came with the discovery of a large (greater than 10 times) long-lived flux enhancement from \\tfn\\ at the time of a major outburst in 2002 June 18. This event was accompanied by many other radiative changes as well as by a large rotational spin-up \\citep{kgw+03,wkt+04}.  This suggests that past flux variability reported in AXPs could be attributed to similar outbursts that went undetected.  We report here, using data from our continuing {\\it RXTE} monitoring program, the discovery of significant pulsed flux variability in \\tfe.  This variability is mainly characterized by two long-lived pulsed flux flares, having well-resolved rises a few weeks long. These are unlike any previously seen flux enhancements in AXPs and SGRs and thus likely represent a distinct physical phenomenon.  We find no evidence for any major associated bursting behavior.  We also report large variations in the spin-down torque on timescales of a few weeks/months.  We find only a marginal correlation between the flux and torque variations.  We argue that this poses another significant challenge to any disk-accretion model for AXPs, but is not inconsistent with the magnetar model. ",
        "conclusions": "\\label{sec:discussion}  The long-lived flux enhancements with well-resolved rises that we have observed in \\tfe\\ are very different from previously detected X-ray flux variations in AXPs and SGRs, which show very abrupt rises associated with major outbursts \\citep[e.g.][]{kgw+03,wkt+04}. The long-lived flux decay in those sources has been attributed to burst afterglow, which is a cooling of the crust following an impulsive heat injection from magnetospheric bursts \\citep{let02}.  The much more gradual flux rises that we have observed in \\tfe\\ comprise a new phenomenon not yet observed in any other AXP, despite several years of careful and frequent {\\it RXTE} monitoring. These flux variations may provide a new diagnostic of the physical origin of the persistent nonthermal emission in SGRs and AXPs, since they are not contaminated by burst afterglow. Also interesting are the large variations in spin-down rate or torque.  Torque variations by nearly a factor of 5 were already reported from {\\it RXTE} observations \\citep{kgc+01}, on timescales of years. Here we have shown that the torque can change by at least a factor of $\\sim$2 more, and on much shorter timescales, namely, a  few weeks to months.  In considering the observed pulsed flux and torque variations, whether they are correlated is an important issue.  Our weekly monitoring of the source unfortunately commenced only after most of the first flare decayed.  Prior to that, the monthly observations, taken in the form of brief snapshots, did not allow anything about the rotational behavior of the source to be determined when phase-coherent timing was not possible.  This was the case during the first flare.  During the second flare, the spin frequency was, interestingly, {\\it most} stable during the rise and peak of the flare.  Furthermore, the stable spin-down rate was at a lower magnitude than the long-term average. Subsequently, $\\sim 60$ days after the flux began to decay, the rate of spin-down began to increase.  Given timing observations during only one flare, it is unclear whether these features are coincidences or not.  However, there is no strong evidence to support otherwise; similar torque variations were seen in the past and were not accompanied by any flaring (see Fig.~\\ref{fig:spin flux spectra}). Significant torque variations  unaccompanied by severe flux variability have been noted for \\tfe\\  prior to our \\rxte\\ monitoring \\citep[e.g.][]{pkdn00}. Nevertheless, statistically, the probability that they are uncorrelated is only 4\\%; studying Figure~\\ref{fig:fit} suggests that if anything, slope transitions are correlated, if not the slopes between transitions.  Continued \\rxte\\ monitoring will help identify any true correlations, particularly if the source exhibits more variability.  Can the magnetar model explain such behavior?  The persistent emission in magnetars has a spectrum that is well described by a two-component model, consisting of a blackbody plus a hard power-law tail. The thermal component is thought to arise from heat resulting from the active decay of a high internal magnetic field \\citep{td96a}; however, thermal X-ray flux changes are not expected on as short a time scale as we have measured in the absence of major bursts. \\citet{tlk02} put forth a model in which the nonthermal component arises from resonant Compton scattering of thermal photons by currents in the magnetosphere.  In magnetars, these currents are maintained by magnetic stresses acting deep inside its highly  conducting interior, where it is assumed that the magnetic field lines are  highly twisted.  These magnetospheric currents in turn twist the external dipolar field in the lesser conducting magnetosphere. These magnetic stresses can lead to sudden outbursts or more gradual plastic deformations of the rigid crust, thereby twisting the footpoints of the external magnetic field and inducing X-ray luminosity changes.  The persistent  non-thermal emission  of AXPs is explained in this  model  as being  generated by these currents through magnetospheric Comptonization and surface back-heating \\citep{td96a,tlk02}.  Changes in X-ray luminosity, spectral hardness, and torque have a common physical origin in this model and some correlations are expected.  Larger twists correspond to harder persistent X-ray spectra, as is observed, at least when comparing the harder SGR spectra to those of the softer AXPs.  As noted by \\citet{kgc+01}, \\tfe's hard photon spectral index ($\\Gamma = 2.9$) suggests that it is a transition object between the AXPs ($\\Gamma \\simeq$3--4) and the SGRs ($\\Gamma = $2.2--2.4). Hence, if during the flares \\tfe's magnetosphere was twisted to the SGR regime, we expect spectral index variations of $\\sim 0.5$.  Spectral measurements of such precision are not feasible with our short \\rxte\\ monitoring observations.  Decoupling between the torque and the luminosity can be accounted for in the magnetar model.  According to \\citet{tlk02}  the torque is most sensitive to the current flowing on a relatively narrow bundle of field lines that are anchored close to the magnetic pole, and so only a broad correlation in spin-down rate and X-ray luminosity is predicted, and in fact is observed for the combined population of SGRs and AXPs \\citep{mw01,tlk02}.  However, for a single source, whether an X-ray luminosity change will be accompanied by a torque change depends on where in relation to the magnetic pole the source of the enhanced X-rays sits.  Similarly, large torque variations, as we have observed, may occur in the absence of luminosity changes if the former are a result of changes in the currents flowing only in the small polar cap region.  Note that energetically, the total release in these flares is comparable to, although somewhat less than, that in the afterglows seen in SGRs and in AXP \\tfn\\ \\citep[see][for a summary]{wkt+04}.  It easily can  be accounted for given the inferred magnetic energy of the star.  Although the magnetar model for AXPs has been spectacularly successful in explaining their most important phenomenology, the anomalous behavior noted for \\tfe\\ raises the possibility that perhaps it has a physical nature different from other AXPs.  It has also been suggested that AXPs might be powered by accretion from fossil disks \\citep{chn00, alp01}.  An increase in luminosity $L_X$ can easily be explained in accretion models by an increase in the mass accretion rate $\\dot{M}$, given that $L_X \\propto \\dot{M}$.  Transient changes in $\\dot{M}$ are perhaps not unreasonable to expect in fossil disk models, given the huge variations seen in $\\dot{M}$ of conventional accreting sources.  However, in an accretion scenario, we expect correlations between luminosity and torque.  In conventional disk-fed accreting pulsars undergoing spin-up, one expects $\\dot{\\nu} \\propto L_X^{6/7}$. Such a correlation is seen approximately in accreting pulsars, with discrepancies possibly attributable to changed beaming or improper measurement of bolometric luminosities, the former due to pulse profile changes, and the latter due to finite bandpasses \\citep{bcc+97}.  As discussed by \\citet{kgc+01}, for a source undergoing regular {\\it spin-down} as in \\tfe, the prediction is less clear; the form of the correlation depends on the unknown functional form of the torque.  For the propeller torque prescription of \\citet{chn00}, we find that $L_X\\propto \\dot{\\nu}^{7/3}$, a much stronger correlation than in the conventional spin-up sources.  For a change in $L_X$ by a factor of $\\sim$3 as we have seen in the rise of the second flare, we would expect a simultaneous change in $\\dot{\\nu}$ by greater than $50\\%$, clearly ruled out by our data.  Conversely, for the abrupt change of $\\dot{\\nu}$ by a factor of $\\sim$2 (near MJD 52740), we expect a change in $L_X$ by a factor of $\\sim$5, definitely not seen.  This appears to pose a significant challenge to fossil-disk accretion models for \\tfe.  Two infrared  observations taken  on MJD 52,324 \\citep{ics+02} and MJD 52372 \\citep{wc02} have shown that the IR counterpart of this source is variable. However, the pulsed X-ray flux at both those epochs was consistent with the quiescent value. Furthermore, even though the X-ray flux has not yet returned to its quiescent value, recent observations show that the source's proposed IR counterpart is consistent with the fainter of the two previous observations \\citep{dk04}.  This decoupling  between the IR and the X-ray flux contrasts with what was observed in AXP \\tfn, whose IR flux increased then decayed in concert with the X-ray flux at the time of its 2002 outburst \\citep[][Tam et al. 2004, in preparation]{kgw+03}.  This is puzzling and suggestive of more than one mechanism for producing IR emission in AXPs."
    },
    "0404/astro-ph0404439_arXiv.txt": {
        "abstract": "{We present results from an ISOCAM survey in the two broad band filters LW2 (5-8.5 $\\mu$m) and LW3 (12-18 $\\mu$m) of a 0.13 square degree coverage of the Serpens Main Cloud Core. A total of 392 sources were detected in the 6.7 $\\mu$m band and 139 in the 14.3 $\\mu$m band to a limiting sensitivity of $\\sim$ 2 mJy. We identified 58 Young Stellar Objects (YSOs) with mid-IR excess from the single colour index $[14.3/6.7]$, and 8 additional YSOs from the $H-K/K-m_{6.7}$ diagram. Only 32 of these 66 sources were previously known to be YSO candidates. Only about 50\\% of the mid-IR excess sources show excesses in the near-IR $J-H/H-K$ diagram. In the 48 square arc minute field covering the central Cloud Core the Class\\,I/Class\\,II number ratio is 19/18, i.e. about 10 times larger than in other young embedded clusters such as $\\rho$ Ophiuchi or Chamaeleon. The mid-IR fluxes of the Class\\,I and flat-spectrum sources are found to be on the average larger than those of Class\\,II sources. Stellar luminosities are estimated for the Class\\,II sample, and its luminosity function is compatible with a coeval population of about 2 Myr which follows a three segment power-law IMF. For this age about 20\\% of the Class\\,IIs are found to be young brown dwarf candidates. The YSOs are in general strongly clustered, the Class\\,I sources more than the Class\\,II sources, and there is an indication of sub-clustering. The sub-clustering of the protostar candidates has a spatial scale of 0.12 pc. These sub-clusters are found along the NW-SE oriented ridge and in very good agreement with the location of dense cores traced by millimeter data. The smallest clustering scale for the Class\\,II sources is about 0.25 pc, similar to what was found for $\\rho$ Ophiuchi. Our data show evidence that star formation in Serpens has proceeded in several phases, and that a ``microburst'' of star formation has taken place very recently, probably within the last 10$^5$ yrs. ",
        "introduction": "\\label{intro}  The youngest stellar clusters are found deeply embedded in the molecular clouds from which they form. There are several reasons why very young clusters are particularly interesting for statistical studies such as mass functions and spatial distributions. Because mass segregation and loss of low mass members due to dynamical evolution has not had time to develop significantly for ages $\\la 10^8$ yrs \\citep{sca98}, the stellar IMF can in principle be found for the complete sample, at least for sufficiently rich clusters. For ages $\\la 10^5$ yrs the spatial distribution should in gross reflect the distribution at birth, which gives important input to the studies of cloud fragmentation and cluster formation. Only in the youngest regions of low mass star formation do we find the co-existence of newly born stars and pre-stellar clumps, which allows one to compare the mass functions of the different evolutionary stages. Low mass stars are more luminous when they are young, being either in their protostellar phase or contracting down the Hayashi track, which permits probing lower limiting masses. Severe cloud extinction, however, requires sensitive IR mapping at high spatial resolution to sample the stellar population of embedded clusters.  ISOCAM, the camera aboard the ISO satellite \\citep{kes96}, provided sensitivity and relatively high spatial resolution in the mid-IR \\citep{ces96}. The two broad band filters LW2 (5-8.5 $\\mu$m) and LW3 (12-18 $\\mu$m), designed to avoid the silicate features at 10 and 20 $\\mu$m, were selected to sample the mid-IR Spectral Energy Distribution (SED) of Young Stellar Objects (YSOs) in different evolutionary phases. According to the current empirical picture for the early evolution of low mass stars \\citep{ada87,lad87,and93,and94}, newborn YSOs can be observationally classified into 4 main evolutionary classes. Class\\,0 objects are in the deeply embedded main accretion phase ($\\ga$ 10$^4$ yrs), and have measured circumstellar envelope masses larger than their estimated central stellar masses, with overall SEDs resembling cold blackbodies and peaking in the far-IR. Class\\,I sources ($\\sim 10^5$ yrs) are observationally characterised by a broad SED with a rising spectral index\\footnote{The spectral index is defined as $\\alpha_{\\rm IR} = d \\log (\\lambda F_{\\lambda})/(d \\log \\lambda)$ and is usually calculated between 2.2 $\\mu$m and 10 or 25 $\\mu$m.} towards longer wavelengths ($\\alpha_{\\rm IR} > 0$) in the mid-IR.  The Class\\,II sources spend some 10$^6$ yrs in a phase where most of the circumstellar matter is distributed in an optically thick disk, displaying broad SEDs with $-1.6 < \\alpha_{\\rm IR} < 0$. At $\\alpha_{\\rm IR} \\approx -1.6$ the disk turns optically thin, and the sources evolve into the ($\\sim 10^7$ yrs) Class\\,III stage where the mid-IR imprints of a disk eventually disappear. A normal stellar photosphere has $\\alpha_{\\rm IR} = -3$. Thus, while Class\\,0 objects are not favourably traced by mid-IR photometry, they are expected to be rare. At the other extreme, Class\\,III sources cannot generally be distinguished using mid-IR photometry since most of them have SEDs similar to normal stellar photospheres. But mid-IR photometry from two broad bands, as obtained in this study with ISOCAM, is highly efficient when it comes to detection and classification of Class\\,I and Class\\,II sources. Thus, considering the fact that these latter objects constitute the major fraction of the youngest YSOs, the ISOCAM surveys provide a better defined sample for statistical studies than e.g. near-IR surveys for regions with very recent star formation \\citep[see][]{pru99}.  This paper presents the results from an ISOCAM survey of $\\sim $ 0.13 square degrees around the Serpens Cloud Core in two broad bands centred at 6.7 and 14.3 $\\mu$m. This cloud, located at $b^{\\rm II} = 5^{\\circ}$ and $l^{\\rm II} = 32^{\\circ}$ at a distance of $259 \\pm 37$ pc \\citep{str96,fes98}, comprises a deeply embedded, very young cluster with large and spatially inhomogeneous cloud extinction, exceeding 50 magnitudes of visual extinction. Only a few sources are detected in the optical \\citep{har85,gom88,gio98}. Serpens contains one of the richest known collection of Class\\,0 objects \\citep{cas93,hur96,wol98,dav99}, an indication that this cluster is young and active. On-going star formation is also evident from the presence of several molecular outflows \\citep{bal83,whi95,hua97,her97,dav99}, pre-stellar condensations seen as sub-mm sources \\citep{cas93,mcm94,tes98,wil00}, a far-IR source (FIRS1) possibly associated with a non-thermal triple radio continuum source \\citep{rod89,eir89,cur93}, a FUor-like object \\citep{hod96}, and jets and knots in the 2.1 $\\mu$m H$_2$ line \\citep{eir97,her97}. Investigations of the stellar content have been made with near-IR surveys \\citep{str76,chu86,eir92,sog97,gio98,kaa99a}, identifying YSOs using different criteria, such as e.g. near-IR excesses, association with nebulosities, and variability.  In this paper we identify new cluster members, characterize the YSOs into Class\\,I, flat-spectrum, and Class\\,II sources, estimate a stellar luminosity function for the Class\\,II sample and search for a compatible IMF and age, and finally describe the spatial distribution of both the protostars and the pre-main sequence population in this cluster.  {\\small \\begin{table*} \\caption[]{Observational parameters, detection statistics and photometric results for each of the 6 ISOCAM rasters CE, CW, CS, D1, D2, D3. See Fig.~\\ref{figmap} and text. \\label{tbl-1}}  \\[ \\begin{tabular}{lcccrrcrrrrrr} \\hline \\noalign{\\smallskip}  & $\\alpha$(2000) & $\\delta$(2000) & Size$^1$ & T$_{\\rm int}$ & pfov & $<$n$_{\\rm ro}>$ & N$_{\\rm det}$ & N$_{\\rm 6.7}$ & N$_{\\rm 14.3}$ & N$_{\\rm both}$ & 1$\\sigma_{6.7}$ & 1$\\sigma_{14.3}$  \\\\ &   &  & (\\arcmin $\\times$ \\arcmin) & (sec) & (\\arcsec) &  &  & & & & (mJy) & (mJy) \\\\ (1) & (2) & (3) & (4) & (5) & (6) & (7) & (8) & (9) & (10) & (11) & (12) & (13) \\\\ \\hline CE    & 18 29 48.3 & 01 16 04.5 &13$\\times$13  & 0.28& 3/6 &4$\\times$13& 133 & 113 & 56 & 51 & 2 &3 \\\\ CW    & 18 29 06.3 & 01 16 01.4 &13$\\times$13  & 2.1 & 6   &7          & 152 & 139 & 43 & 38 &0.8&2  \\\\ CS    & 18 29 52.5 & 01 02 37.9 &13$\\times$16.5& 0.28& 6   &4$\\times$13& 165 & 152 & 41 & 34 &1.2&4  \\\\ D1    & 18 29 48.7 & 01 15 20.5 &1.8$\\times$4.6& 2.1 & 3   &44         &  20 &  15 &  9 &  9 & 1 &2  \\\\ D2    & 18 29 52.3 & 01 15 20.8 &1.8$\\times$4.6& 2.1 & 3   &44         &  24 &  19 & 11 & 10 & 1 &4.5  \\\\ D3    & 18 29 57.8 & 01 12 55.4 &4.6$\\times$1.8& 2.1 & 3   &44         &  30 &  21 & 15 & 13 & 1 &3 \\\\ Tot$^2$ &            &            &              &     &     &           & 421 & 392 &140 &124 &   &  \\\\   \\noalign{\\smallskip} \\hline \\end{tabular} \\] $^1$ Approximate size since each raster has tagged edges. \\\\ $^2$ The total number is corrected for multiply observed sources in overlapped regions. \\\\  \\end{table*} } ",
        "conclusions": "We have used ISOCAM to survey 0.13 square degrees of the Serpens Cloud Core in two broad bands centred at 6.7 and 14.3 $\\mu$m. In combination with our ground based deep $JHK$ imaging of the 48 square arcminute central region as well as additional $K$ band imaging of about 30 square arcminutes to the NW, and a 1.3 mm IRAM map of the central dense filament, we have investigated the mid-IR properties of the young stellar population. The following results are found:  \\begin{itemize} \\item The number of point sources with reliable flux measurements are 392 at 6.7 $\\mu$m and 139 at 14.3 $\\mu$m. Of these, 124 are detected in both bands.  \\item On the basis of one single colour index, $[14.3/6.7]$, we found that 53 of the 124 objects possess strong mid-IR excesses. Only 28 of these were previously suggested YSO candidates. The large scale spatial distribution of these mid-IR excess sources is strongly concentrated towards the Cloud Core, where it is elongated along NW-SE.  \\item The near-IR $J-H/H-K$ diagram is found to have an efficiency of less than 50\\% in detecting IR excess sources. This efficieny is comparable to that one found in star formation regions in general from ISOCAM data \\citep{kaa00} and $\\rho$ Ophiuchi in particular \\citep{bon01} and should be kept in mind when interpreting $JHK$ based data in terms of disk fractions.  \\item The $H-K/K-m_7$ diagram separates well intrinsic IR excess from the effects of reddening. From this diagram we were able to increase the number of mid-IR excess sources to 70. This means a fractional increase in the investigated region by 25\\%.  \\item Combination of near-IR and mid-IR photometry for reddened stars without IR excesses enables us to estimate the extinction at 6.7 and 14.3 $\\mu$m relative to that in the $K$ band in the Serpens direction. We find A$_{7}$ = 0.41 A$_K$ and A$_{14}$ = 0.36 A$_K$. Our results agree with the extinction law measured towards the Galactic centre by \\citet{lut96}, as well as with the results from the ISOGAL survey \\citep{jia03}.  \\item Classification of the Serpens YSOs in terms of the SED indices gives 20 Class\\,I sources, 13 flat-spectrum sources, and 43 Class\\,II sources. The number of Class\\,I sources appears to be exceptionally large along the NW-SE oriented ridge, and the number fraction Class\\,I/Class\\,II is almost 10 times higher than normal, an indication that this part of the cluster is extremely young and active.  \\item The mid-IR luminosities of the Class\\,I sources are on the average larger than those of the Class\\,II sources. Since Class\\,II sources are expected to have SEDs which peak approximately in the near and mid-IR, in contrast to Class\\,Is which radiate most of their luminosity in the far-IR, we conclude that there is a weak indication of luminosity evolution in the Serpens Cloud Core.  \\item We have estimated extinction and stellar luminosities for the 43 Class\\,II sources found in our survey. The Class\\,II luminosity function is found to be compatible with co-eval formation about 2 Myrs ago and an underlying IMF of the three-segment power-law type \\citep{kro93,sca98}, similar to the mass function found in $\\rho$ Ophiuchi \\citep{bon01}. With this assumption on age, every fifth Class\\,II is a young brown dwarf.  \\item Except for one case, the Class\\,I sources are exclusively found in sub-clusters of sizes $\\sim$ 0.12 pc distributed along the NW-SE oriented ridge. The sub-clusters also contain several flat-spectrum sources. In total, each core has formed between 6 and 12 protostars (lower limit) within a very short time. The spatial distribution of the Class\\,II sources, on the other hand, is in general much more dispersed.  \\item On the assumption that the protostar candidates follow roughly the same IMF as the Class\\,IIs, we derive a SFR of 6.1 $\\times$ 10$^{-5}$ M$_{\\odot}$/yr and a local SFE of 9\\% in the recent microburst of star formation forming the dense sub-clusters of protostars.  \\end{itemize}  The results presented in this study show evidence that the sub-clusters in the central part of the Serpens Cloud Core were formed by a recent microburst of star formation. The extreme youth of this burst, deduced from the compact clusters of protostar candidates, is supported by independent investigations, such as the rich collection of Class\\,0 sources found by \\citet{cas93} and \\citet{hur96} in the same regions. In addition to the clustered protostar population, we also find a more distributed population of Class\\,II sources, for which we have deduced an age more or less typical of Class\\,IIs found in other regions (2 Myr). In addition to these YSO generations, there is probably also a population of Class\\,III sources, which is undistinguishable from the field star population in our study, but which should be looked for with proper search tools, such as e.g. X-ray mapping \\citep{gro00}.  \\citet{cas93} suggested that their submm sources without near-IR counterparts represented a second phase of active star formation in Serpens. Here we show the co-existence of sources in the various evolutionary stages from Class\\,II and flat-spectrum to Class\\,I and Class\\,0 sources. While we find an age estimate of 2 Myr for the Class\\,II population, it is highly unlikely that the clustered Class\\,I and flat-spectrum sources are older than a few 10$^5$ yrs. Our results therefore support the conclusions of \\citet{cas93} that star formation has proceeded in several phases in Serpens."
    },
    "0404/astro-ph0404263_arXiv.txt": {
        "abstract": "We present a \\xmm~observation of Markarian 304, a Seyfert 1 galaxy at $z$ = 0.066. The \\epic~data show that MKN~304 is affected by heavy (\\nh~$\\approx$ 10$^{23}$ \\cm2) obscuration due to ionized gas.   A two--phase warm absorber provides an adequate parameterization of this gas.  The  ionization parameter of the two components is $\\xi \\approx$ 6 \\cgs~and $\\xi \\approx$ 90 \\cgs, respectively. The observed continuum photon index ($\\Gamma \\approx$ 1.9) is typical  for Seyfert 1 galaxies. Two significant emission lines are detected at 0.57 keV and 6.4 keV, respectively.  The former is mostly likely due to He--like oxygen triplet emission arising from an ionized plasma (maybe the warm absorber itself). The latter is due to fluorescent emission of K--shell iron in a low--ionization state (FeI--XV). The upper limit for the line width of $\\sigma_{K\\alpha} <$ 0.18 keV most likely rules out an origin in the inner parts of the accretion disk. Interestingly, the strength of such line is consistent with the possibility that the emission is produced in the warm absorber itself. However, a  substantial contribution from the torus is plausible too. We have also found a weak (4\\% of the primary continuum) soft excess emission component.  The presence of this excess could be explained by either emission/scattering from a warm gas or partial covering, or a combination of them. ",
        "introduction": "The bulk of the X--ray emission in active galactic nuclei (AGNs) is produced via Comptonization of UV photons in the inner parts of the accretion flow close to the central supermassive black hole.  The resulting power law spectrum is modified by a number absorption and emission spectral features due to the reprocessing of this primary continuum (Mushotzky, Done \\& Pounds 1993). They are clear signatures of the presence  of  large amounts of gas characterized by different physical properties in the circumnuclear region.  One of the most common spectral features is the so--called {\\it warm} (i.e. partially ionized) {\\it absorber}  (Halpern 1984; Pan et al. 1990, Turner et al. 1993). Such a component has been found to contribute significantly to the opacity of the X--ray primary emission in $\\approx$ 50\\% of the Seyfert 1 galaxies with column densities up to \\simgt~10$^{23}$ \\cm2 (Reynolds 1997).  In particular, OVII (0.739 keV) and OVIII (0.871 keV) absorption edges resulted to be the most prominent  absorption signatures of the ionized gas in low--resolution {\\em ASCA} observations.  The advent of  grating spectrometers on--board \\xmm~and \\chandra~has dramatically enriched our knowledge in this field. High--resolution soft X--ray spectra of bright Type 1 AGNs present a wealth of absorption lines (e.g. Collinge et al. 2001; Kaspi et al. 2001; Kaastra et al. 2002),  which allow to infer a few remarkable",
        "conclusions": "The spectrum of MKN~304 unveiled by the \\xmm~observation discussed in this paper is complex both in the soft  and in the hard energy band (Fig.~1).  Our analysis fully confirms previous results on narrower energy ranges which reported a very flat continuum for this source, and  suggested the existence of heavy obscuration. However, due to the limited bandpass and the poor sensitivity of the previous observations, it was impossible to accurately investigate the properties of the absorption, and discriminate between a cold (neutral) or warm (ionized) medium.    This \\xmm~observation of MKN~304 has clearly revealed that the strong obscuration (\\nh~$\\approx$ 10$^{23}$ \\cm2) occurs in an ionized gas.  We have indeed found that the best fit model to the \\epic~data consists of a two--phase warm absorber with a high-- and a low--ionized component, a faint soft excess component and two emission lines at $\\sim$ 0.57 keV and $\\sim$ 6.4 keV,  respectively (see Model E in Table~1).  \\subsection{Warm absorber} \\label{sect:warm}  We fixed {\\it a priori} the temperature values at T = 1.5 $\\times$ 10$^{5}$ K and T = 3.0 $\\times$ 10$^{4}$ K for the two ``phases'' of the ionized plasma since they represent typical values observed in sources with a multi--phase warm absorber (Netzer et al. 2003; Kaastra et al. 2002).  From the spectral analysis we derived $\\xi^{T5}$ = 89.3$^{+13.9}_{-12.0}$ \\cgs~and $\\xi^{T4}$ = 5.9$^{+2.4}_{-0.9}$ \\cgs~for the hot and the cold component, respectively (see model E in Table 1).  Our findings for MKN~304 appear to be in agreement with  a ``multi--zone'' structure for the ionized absorber in Seyfert galaxies as discovered by high resolution spectroscopic observations recently performed  with grating spectrometers on--board of \\xmm~and \\chandra. They indeed resolved the complex structure of the AGN circumnuclear absorbing material into a multi--phase plasma with  several  orders of magnitude spread in $\\xi$  and systematic blueshifts (Blustin et al. 2003; Kaastra et al. 2002).   This observational evidence leads to interpret the warm absorber phenomenon essentially in terms of a photoionized  outflowing wind consisting of different inhomogeneous co--existing regions (or shells) characterized by a broad distribution in $\\xi$, temperature and density (Krolik 2002; Behar et al. 2003; Elvis 2000).  Interestingly, while our best--fit values for ionization parameters of the two absorbing components are similar to those found in other Seyfert 1 galaxies (Kaastra et al. 2002; Blustin et al. 2003),  the corresponding column densities (\\nhh~$\\approx$ 9 $\\times$ 10$^{22}$ \\cm2~and \\nhc~$\\approx$ 2 $\\times$ 10$^{22}$ \\cm2) are amongst the highest  seen by \\chandra~and \\xmm~so far.  Given these high values of \\nh, photoionization edges are expected to dominate the absorption features in MKN~304.  In fact, at high column densities the UTA absorption lines saturate. Edges (alongside with $2p-nd$ (with $n >3$) absorption lines)  provide the major contribution  to the absorption spectrum (Behar, Sako \\& Kahn 2001). Furthermore as shown in Fig.~\\ref{fig:ext2soft}, the absorption feature in MKN~304 is remarkably broad and deep and, hence, hard to reconcile with what simply expected  from UTA\\footnote{Nonetheless, it is also worth noting that the colder ionized component has values  of \\nh, $\\xi$ and $T$ approximately consistent with the production of UTA features around 0.75 keV. Unfortunately, the lack of an adequate energy resolution does not allow us to further test this hypothesis.}.  Besides absorption, also emission lines are expected to arise from the warm medium.  Due to the underlying illuminating bright continuum and the forest of absorption lines such emission features are however difficult to detect in low--resolution data. In particular He--like triplets of O  and Ne are usually expected.  The OVII line detected at 0.57 keV in the spectrum of MKN~304 is therefore consistent with reprocessing in the warm gas.  However, in sources showing both absorption and emission lines the former usually appear to have different  velocities respect to the latter (and in many cases consistent with the galaxy rest--frame, e.g. Collinge et al. 2001), so it is not clear if emitter and absorber systems are intrinsically the same.  For instance, Kaastra et al. (2002) found the optical Narrow Line Region as the most likely origin of OVII and Ne IX emission  lines observed in NGC~5548. Unfortunately, \\epic~data do not have enough resolution to shed light on this issue.  We used CLOUDY (Ferland 2001) to further test the idea of a possible origin of the oxygen emission line from the same {\\it warm} medium responsible of the absorption in MKN~304. We employed the table ``AGN''  ionizing continuum, with an X--ray photon index fixed to 1.88, a hydrogen column density set to 9 $\\times$ 10$^{22}$ \\cm2, and an ionization parameter equivalent to the highest one obtained  from XSPEC. For this high ionization state of the gas, the transmitted X-ray spectrum is only marginally dependent on the gas density, the EUV continuum and the UV/X-ray flux ratio. Following the results of the XSPEC fits, the total output spectrum from CLOUDY has been computed as the sum of the incident and transmitted spectra with relative contributions of 5 and 95\\%, respectively. In this composite spectrum the most prominent emission feature is due to OVII, whose EW ($\\sim$ 110 eV) is in excellent agreement with that inferred by our fit, i.e. EW = 105$\\pm$18 eV. Other emission features in the model (i.e. NeIX, NVI, MgXI) are several times weaker, except for the  Fe K$\\alpha$ line, whose EW is similar to that of OVII (see Sect. 4.3.). Therefore, we conclude that although this result does not ultimately proof the origin of the He--like oxygen line in the warm absorber, it clearly supports such a scenario.   Finally, the combination of an unreddened continuum in the optical band (H$\\alpha/$H$\\beta \\sim$ 3; e.g. Osterbrock 1977) and a column density of  \\nh~$\\sim$ 10$^{23}$ \\cm2~suggests a dust--free X--ray absorbing medium in MKN~304, maybe located  within the dust sublimation radius (see Crenshaw, Kramer \\& George 2003 and references therein).  \\subsection{Soft excess}  The soft excess component (whose normalization is $\\approx$ 4\\% of the primary emission) detected in MKN~304  appears similar to that observed in many X--ray obscured objects (Turner et al. 1997).  There are two possibilities to explain the origin of such excess, i.e: scattered$/$reflected emission from the photoionized gas  in the warm absorber and ``partial--covering''.   We discuss in turn these hypothesis in the following.  \\xmm~and \\chandra~results on Seyfert 2 galaxies (Sako et al. 2002 and reference therein) suggest that the  soft X--ray emission in these objects is the result of primary continuum scattering underlying a blend of strong emission lines, mostly from H-- and He--ions of C, N, O, Ne, and Fe.  Since warm absorber regions show physical properties (i.e. $\\xi$, column densities, temperatures)  similar to those observed in the X--ray  photoionized emitting plasma in Seyfert 2 galaxies, it is likely that they are different manifestations of the same phenomenon (Kinkhabwala et al. 2002) simply seen at different inclination angles.  Within this scenario the soft excess found in MKN~304 would be therefore  due to scattered$/$reflected emission from the ionized outflowing plasma.  Interestingly, Ogle et al. (2000), Bianchi et al. (2003) and Iwasawa et al. (2003) detected indeed an extended soft X--ray emission dominated by narrow emission lines overlapping the optical Narrow Line Region in NGC~4151, NGC~5506 and NGC~4388,  respectively.  According to this hypothesis, a significant contribution from this extended soft X--ray component to the OVII emission line would be also expected (see Sect.~4.1).  An alternative explanation of this ``soft excess'' could also be some fraction of the primary continuum leaking  through the clumpy obscuring material (i.e. along the line of sight the absorber does not completely cover the nuclear source):  the  so--called ``partial-covering'' scenario. Furthermore, as mentioned in Sect.~\\ref{sect:warm}, clumpiness in the ionized absorbing gas appears to be very likely.  Unfortunately, the present data do not allow to discriminate these hypothesis.  The unabsorbed power law in model E (with a slope identical to the primary continuum one) can be therefore considered as a good approximation of the real (more complex) spectrum. However, given the energy resolution of \\epic~CCD cameras and the relatively faintness of MKN~304, a  description of this soft X--ray component in terms of many discrete emission features, as observed in the high resolution spectra of some bright AGNs (see Kinkhabwala et al. 2002), cannot be ruled out.  Finally we rule out the hypothesis of an origin of the soft excess from reprocessing (i.e. reflection) by an ionized disk since this interpretation is inconsistent with the observed centroid of the iron line. In fact, assuming that the line arises from such an ionized disk (as expected if a strong relection component is present; Ross, Fabian \\& Young 1999),  it should be emitted at energy higher than 6.4 keV from iron at high ionization states.  On the contrary, we found an energy of the  Fe K$\\alpha$ emission line  strictly consistent only with ``neutral iron'' (see below).  \\subsection{Iron line} We report the first clear evidence of an iron fluorescence emission line around 6.4 keV in the spectrum of MKN~304.  The detection of a Fe K$\\alpha$ line at E = 6.34$\\pm$0.06 keV constrains the ionization parameter of the emitting medium to be $\\xi$ \\simlt~100 \\cgs, higher values of $\\xi$ would indeed produce a feature peaked at higher energies (Kallman \\& McCray 1982) excluded by the present data. Interestingly, both cold and hot warm absorber components that resulted in our best fit (model E) to the \\epic~spectrum have values of $\\xi$ which match well with this constraint.  This suggests that a significant component of the iron emission line may be originated in the warm medium.  We have run the same CLOUDY model as in Sect.~4.1 to predict the strength  of the Fe K$\\alpha$ emission line.  The expected value ($\\approx$ 110 eV) falls in the range we measured from the \\xmm~spectrum (i.e. EW = 79$\\pm$35 eV).  So, interestingly, the iron line detected in MKN~304 is consistent with an origin in the warm absorber observed in this source.  We have also tried to fit the data with a model including a neutral reflection Compton component ({\\bf pexrav} model in XSPEC) in order to test a possible origin of the iron line from the  optically--thick accretion disk.  This spectral parametrization yielded a good fit with a \\xnu~= 335(357): however, there is no statistical improvement respect to Model E.  Although the detection of cold reflection results therefore only marginal, we cannot rule out its presence. The observed spectral range (0.3--12 keV) of this \\epic~observation is not adequate to constrain the strength of this ``bump--like'' reprocessed emission feature  peaked at $\\approx$ 30--40 keV (Lightman \\& White 1988).  The resulting upper limit for the covering factor of the material irradiated by the X--ray source is $R\\equiv$$\\Omega/2\\pi <$ 0.7.  Both the value of the equivalent width of the line and $R$ suggest that the reflector -- if existing -- subtends a solid angle of  \\simlt~$\\pi$ steradian to the X--ray source, i.e. intercepts \\simlt~50\\% of the primary continuum once a slab geometry for the reflecting medium is assumed (George \\& Fabian 1991).  Furthermore the narrow profile of the Fe K$\\alpha$ line ($\\sigma_{K\\alpha} <$0.18 keV corresponding to $\\Delta$$v$(FWHM) \\simlt~8000 km/s for the line velocity width) implies that even if the reprocessing medium is the accretion disk, the fluorescence does not take place  from its innermost (relativistic) regions.  A substantial contribution to the Fe K$\\alpha$ emission could be provided from the putative molecular  torus located at 1--10 pc from the central X--ray source as invoked in the Unification model of AGNs (Antonucci 1993). In this case the material  is assumed to be Compton--thick (i.e. with \\nh~\\simgt~1.5 $\\times$  10$^{24}$  \\cm2, see Matt, Guainazzi \\& Maiolino 2003), and the emission is due to the reflection of the primary continuum off the inner walls of the torus into our line--of--sight (Ghisellini et al. 1994).  The observed properties of the Fe K$\\alpha$ line in MKN~304 appear to be very similar to the results obtained by the spectral analysis of the brightest Seyfert 1 galaxies using \\xmm~(Pounds \\& Reeves 2002) and  \\chandra~(Padmanabhan \\& Yaqoob 2002) data.  These works indeed shows the ``ubiquitous'' presence of a narrow iron line component centered at 6.4 keV with an EW  $\\approx$ 50--100 eV and $\\sigma_{K\\alpha} <$120 eV in the X--ray spectra of the low--luminosity (i.e. with \\lum~$<$ 10$^{45}$ \\ergs) AGNs."
    },
    "0404/astro-ph0404055_arXiv.txt": {
        "abstract": "In a 1996 JRO Fellowship Research Proposal (Los Alamos), the author suggested that neutrino oscillations may provide a powerful indirect energy transport mechanism to supernovae explosions.  The principal aim of this addendum is to present the relevant \\textit{unedited} text of Section 1 of that proposal. We then briefly remind, (a) of an early suggestion of Mazurek on vacuum neutrino oscillations and their relevance to supernovae explosion, and (b) Wolfenstein's result on suppression of the effect by matter effects. We conclude that whether or not neutrino oscillations play a significant role in supernovae explosions shall depend if there are shells/regions of space in stellar collapse where matter effects play no essential role. Should such regions exist in actual astrophysical situations, the final outcome of neutrino oscillations on supernovae explosions shall depend, in part, on whether or not the LNSD signal is confirmed. Importantly, the reader is reminded that neutrino oscillations form a set of flavor-oscillation clocks and these clock suffer gravitational redshift which can be as large as 20 percent. This effect must be incorporated fully into any calculation of supernova explosion. ",
        "introduction": "The following is an unedited text of Section 1 of author's 1996 J Robert Oppenheimer  Fellowship Research Proposal:\\footnote{ The reader is reminded that such proposals are written for a very broad readership, as the evaluators come not only from physics, but fields as far as biology, and other sciences. The proposal was submitted by M. B. Johnson in July 1996 in his ``Nomination of Dharam Ahluwalia for Oppenheimer Fellowship.'' The addendum title coincides with the title of Section 1 of the proposal.}  Neutrinos were introduced in physics by Pauli to save conservation of energy and momentum in the $\\beta$-decay: Neutron $\\rightarrow$ Proton + Electron + Anti-electron Neutrino. All the planets and  galaxies are embedded in a sea of neutrinos with a number density of roughly 100 neutrinos/cm$^3$. Our own Sun shines via thermonuclear processes that emit neutrinos in enormous number. Because of their weak interactions, neutrinos, unlike photons, can pass through extremely dense matter very efficiently. This fact makes neutrinos primary agents for energy transport in the dense matter associated with supernovae and neutron stars.  Since their initial experimental observation by Frederick Reines and C. L. Cowan, neutrinos are now known to exist in three types. These types are called ``electron,'' ``muon,'' and ``tau'' and are generically written as $\\nu_e$, $\\nu_\\mu$, and $\\nu_\\tau$. A series of empirical anomalies indicates that  neutrinos may not have a definite mass but, instead, be in a linear superposition of three different mass eigenstates. The mass differences in the underlying mass eigenstates would cause a neutrino of one type to ``oscillate'' to a neutrino of another type as may have been seen recently at the LSND neutrino oscillation experiment at LANL. The phenomenon of neutrino oscillations, if experimentally confirmed, will have profound consequences not only for nuclear and particle physics but also for astrophysics and cosmology.   I have already noted the neutrinos to be prime drivers of supernova explosions. The phenomenon of neutrino oscillations will alter the evolution of supernova explosion. The basic problem that still stands unsolved is a robust theory of supernova explosions. In the context of supernova explosions, and the problem of obtaining successful explosions, I now follow Colgate \\textit{et al.} [S. A. Colgate, M. Herant, and W. Benz, Phys. Rep. {\\bf 227}, 157 (1993)] and assume that the matter next to the neutron star is heated by neutrinos from the cooling neutron star. They note that in some models ``this result in strong, large scale convective flows in the gravitational field of the neutron star that can drive successful, albeit weak, explosions.'' I emphasize that all authors find that without ``fine tuning'' the explosions are weak and lack about five percent of the energy needed for an explosion. Qualitatively, this missing energy needed for a robust model of explosion may be provided if the length scales over which neutrino oscillations take place are of the same order of magnitude as the spatial extent of a neutron star and neutrino-sphere, because while  \\begin{quote} the energy flux in each of the electron neutrinos and antineutrinos is about $L_{\\nu_e}\\approx   L_{\\overline{\\nu}_e}\\approx$ few$\\times$  $10^{52}$   ergs s$^{-1}$, with comparable fluxes of $\\nu_\\mu$, $\\overline{\\nu}_\\mu$, $\\nu_\\tau$, and  $\\overline{\\nu}_\\tau$, \\end{quote} the \\begin{quote} average energy of $\\nu_e$ is about 10 MeV, the average energy of other neutrinos is \\textit{higher} by a factor of 2 for $\\nu_\\mu$    and   $\\overline{\\nu}_\\mu$, and by a factor of 3 for    $\\nu_\\tau$, and  $\\overline{\\nu}_\\tau$. \\end{quote} Any oscillation between neutrinos of different flavors is, therefore, an indirect energy transport mechanism towards the actively interacting $\\nu_e$    and   $\\overline{\\nu}_e$. Qualitatively this contributes in the direction of the robustness of the explosion. My resent work, with C. Burgard, on the solution of terrestrial neutrino anomalies provides precisely the neutrino oscillation parameters that yield the oscillation length scales of just the right order of magnitude for supernova physics (and in addition predict the observed solar neutrino deficit).   In order to make these qualitative arguments quantitative two additional physical processes affecting the above indicated \\textit{vacuum} neutrino oscillations must be incorporated: (a) The presence of large electron densities in astrophysical environment makes it necessary that relevant matter induced effects, suggested by Mikheyev, Smirnov and Wolfenstein, be considered, and (b) My work, with C. Burgard, on gravitationally induced neutrino oscillation phases also indicates  that strong gravitational fields associated with neutron stars may introduce important modifications to neutrino oscillations, and hence to the suggested energy transport mechanism via neutrino oscillations. As part of my JRO studies I propose to implement the above outlined program quantitatively. My quantitative and qualitative studies so far give reasons to claim that there is every physical reason to believe that the ``missing energy'' in the non-robust models for supernova explosion, the anomaly in the observed deficit in the solar neutrino flux, the excess $\\overline{\\nu}_e$ events seen at LSND at Los Alamos, and the anomaly associated with atmospheric neutrinos, \\textit{all} arise from the same underlying new physics \\textemdash the phenomenon of neutrino oscillations from one type to another. It is of profound physical importance to place these suspected physical connections on firm quantitative foundations. ",
        "conclusions": ""
    },
    "0404/hep-ph0404083_arXiv.txt": {
        "abstract": "{ In the 3-flavour framework we derive a simple approximate analytic expression for the day-night difference of the flux of solar $\\nu_e$ at terrestrial detectors which is valid for an arbitrary Earth density profile. Our formula has the accuracy of a few per cent and reproduces all the known analytic expressions for the Earth matter effects on the solar neutrino oscillations obtained under simplifying assumptions about the Earth's density profile (matter of constant density, 3 layers of constant densities, and adiabatic approximation).  It can also be used for studying the Earth matter effects on the oscillations of supernova neutrinos.  We also discuss the possibility of probing the leptonic mixing angle $\\theta_{13}$ through day--night asymmetry measurements at future water Cherenkov solar neutrino detectors. We show that, depending on the measured value of the asymmetry, the current upper bound on $\\theta_{13}$ may be improved, or even a lower bound on this mixing parameter may be obtained. } ",
        "introduction": "Solar neutrinos coming to terrestrial detectors at night travel some distances inside the Earth and so their oscillations are affected by the Earth's matter. This leads to a difference between the day-time and night-time solar neutrino signals -- the day-night, or ``regeneration'' effect \\cite{MS,Bouchez:1986kb,Cribier:1986ak,Carlson:1986ui,Baltz:1987hn,Dar:1987pj, Gonzalez-Garcia:2000dj}. Solar neutrino oscillations~\\footnote{The presence of non-oscillation phenomena is now strongly constrained by a combination of KamLAND and solar neutrino data, see, e.~g.  Ref.~\\cite{pakvasa:2003zv} and references therein.} depend mainly on two parameters, the mixing angle $\\theta_{12}$ and the mass squared difference $\\Delta m_{21}^2$. For example, the recent three--neutrino global fit of solar, atmospheric, reactor and accelerator neutrino data of Ref.~\\cite{Maltoni:2003da} gives for the solar neutrino oscillation parameters the $3\\sigma$ allowed ranges $\\theta_{12}=(28.6\\div 38.6)^\\circ$ and $\\Delta m_{21}^2=(5.4\\div 9.5)\\times 10^{-5}$ eV$^2$, with the best-fit values $\\theta_{12}=33.2^\\circ$ and $\\Delta m_{21}^2=6.9\\times 10^{-5}$ eV$^2$.  In the three-flavour framework there is an additional dependence on the mixing angle $\\theta_{13}$ which determines the component of the electron neutrino in the third mass eigenstate $\\nu_3$ separated by a large mass gap from the other two mass eigenstates. This mixing parameter is mainly constrained by the data of the CHOOZ reactor experiment \\cite{apollonio:1999ae} and is known to be small, \\begin{equation} \\theta_{13} \\lesssim 9.8^\\circ~(13.4)^\\circ\\,, \\quad\\mbox{or}\\quad \\sin \\theta_{13} \\lesssim 0.17~(0.23) \\label{eq:th13bound} \\end{equation} at 90\\% C.L. (3$\\sigma$) \\cite{Maltoni:2003da}. The smallness of $\\theta_{13}$ implies that its effect on the solar neutrino oscillations is rather mild; however, with increasing accuracy of the data even relatively small effects have to be taken into account.  Three-flavour analyses of the solar neutrino data, including the day-night effect, have been performed in a number of recent studies \\cite{Maltoni:2003da,Bandyopadhyay:2003pk,deHolanda:2003nj, Gonzalez-Garcia:2003qf,Fogli:2002au,Blennow:2003xw}. In particular, in Ref.~\\cite{Blennow:2003xw} a simplified analytic treatment of the day-night effect in the three--flavour framework was given, and it was shown that the day-night $\\nu_e$ flux difference scales essentially as $\\cos^6\\theta_{13}$ rather than $\\cos^4\\theta_{13}$, which is the expected scaling law for the day-time flux.  The authors of \\cite{Blennow:2003xw} considered the Earth matter effect on the solar neutrino oscillations under the assumption of constant matter density of the Earth. The purpose of the present paper is to extend their analysis beyond the constant density approximation, and to scrutinize to what extent day--night asymmetry measurements may become a significant way to probe for the mixing angle $\\theta_{13}$. Using the relative smallness of the matter-induced potential of neutrinos inside the Earth, we derive a simple and accurate analytic formula for the day-night flux difference in the case of an arbitrary Earth's density profile. We compare our result with the known analytic formulas obtained under simplifying assumptions about the density profile of the Earth as well as with the results of the exact numerical calculations.  The paper is organized as follows. In sec. \\ref{sec:day-night-effect} a general consideration of the day-night effect in the 3-flavour framework is given. In sec. \\ref{sec:an-analyt-expr} we present our main result -- the derivation of an approximate 3-flavour analytic formula for the day-night $\\nu_e$ flux difference in the case of an arbitrary Earth density profile.  We also discuss a number of special cases for which the day-night flux difference can be found analytically.  These include approximating the matter density profile of the Earth by one layer or three layers of constant densities, and also the adiabatic approximation to neutrino evolution inside the Earth.  In sec.  \\ref{sec:comp-with-numer} the comparison of our formulas with the results of the exact numerical calculations is performed. In sec.~\\ref{sec:future} we make a quantitative study of the simple correlation between the day--night asymmetry in the solar neutrino flux and the magnitude of the mixing angle $\\theta_{13}$ characterizing leptonic CP violation in neutrino oscillations~\\footnote{There are additional CP violating phases but they do not affect neutrino oscillations, only lepton-number changing processes, like neutrino-less double beta decay~\\cite{schechter:1980gr,schechter:1981gk,bilenky:1980cx,doi:1981yb}.}. We use this correlation in order to forecast the degree to which the mixing angle $\\theta_{13}$ can be probed at future solar neutrino measurements of the day--night asymmetry. We discuss our results in sec.~\\ref{sec:disc-outlook}. In the Appendix, we present the derivation of an improved perturbation-theoretic formula for the Earth regeneration factor. ",
        "conclusions": "\\label{sec:disc-outlook}  We have derived a simple and accurate analytic expression for the day-night difference of the flux of solar $\\nu_e$ coming to terrestrial detectors, with 3-flavour effects fully taken into account. Our approach was based on a simple perturbation theory in the matter-induced potential of neutrinos inside the Earth, without any assumptions regarding the Earth's density profile. Our results are therefore valid for an arbitrary density profile.  We have checked our analytic formula by comparing it with the results of the exact numerical calculations with the 3-flavour evolution equation, and found that the accuracy of the analytic approach is typically about a few per cent when the neutrino path length inside the Earth is small, or when the integration over the zenith angles or averaging over neutrino energies are performed. On the other hand, when both the neutrino energy and zenith angle are fixed, the formula in Eq.~(\\ref{P2e7}) does not provide a sufficient accuracy. For this case we have derived an improved perturbation-theoretic formula [Eq.~(\\ref{eq:adiab})]. Compared to the expression in Eq.~(\\ref{P2e7}) it contains the adiabatic oscillation phase in the integrand instead of the in-vacuum one. The improved formula gives a perfect agreement with the exact numerical results even when no averaging over the zenith angles or neutrino energy (i.e., over the oscillation phase) is performed.  We have studied the dependence of our results on the leptonic mixing angle $\\theta_{13}$ which is of great interest for many reasons, most importantly because it governs CP violation in neutrino oscillations. We have found that for an arbitrary Earth density profile the day-night difference of the solar neutrino flux at the detector scales mainly as $c_{13}^6$, as previously found for the constant density profile in \\cite{Blennow:2003xw}.  The remaining (milder) dependence stems primarily from the $\\theta_{13}$-dependence of the neutrino mixing angle in matter at the production point inside the Sun.  Although the smallness of $\\theta_{13}$ implies that the solar neutrino data, including the day-night asymmetry, should depend rather weakly on this parameter, this dependence may not be negligible. For example, the current $3\\sigma$ limit $\\sin^2\\theta_{13}<0.054$ \\cite{Maltoni:2003da} means that the day-night difference of the solar neutrino signal can be suppressed by up to 15\\% compared to the 2-flavour (i.e. $\\theta_{13}=0$) case. If one fixes $\\Delta m_{31}^2=2\\times 10^{-3}$ eV$^2$ (which is the best-fit value coming from the analysis of the atmospheric neutrino oscillations performed by the \\mbox{Super-K} collaboration~\\cite{SKatm04}), then the $3\\sigma$ upper bound is $\\sin^2\\theta_{13}<0.066$ \\cite{Maltoni:2003da}, and the day-night difference of the solar neutrino signal can be suppressed by as much as 20\\% compared to the 2-flavour case.  While the accuracy with which the solar neutrino day-night effect is measured by the current experiments is insufficient for probing the value of $\\theta_{13}$, future very large water Cherenkov detectors, such as UNO or Hyper-Kamiokande, may be able to provide a significant information on it. Depending on the measured value of $A_{ND}$, the current upper bound on $\\theta_{13}$ may be improved, or even a lower bound on this mixing parameter may be obtained. An advantage of the day-night asymmetry as opposed to the day-time signal is that $A_{ND}$ is independent of the overall normalization of the solar neutrino flux, which is currently not known with sufficient accuracy.  It should be noted that the Earth matter effects on the oscillations of the supernova neutrinos inside the Earth are governed by the same quantity, $P_{2e}^\\oplus-P_{2e}^{(0)}$, which determines the Earth matter effects on the solar neutrino oscillations (see, e.g., \\cite{Dighe:1999bi}).  Therefore our results, Eqs.~(\\ref{P2e7}) and (\\ref{eq:adiab}), can also be used for studying the supernova neutrino oscillations inside the Earth.  Note, however, that the typical energies of supernova neutrinos $(E\\sim 30$ MeV) are higher than those of solar neutrinos, and so the expected accuracy of our approximation for supernova neutrinos is $\\sim$ (10 -- 15)\\%."
    },
    "0404/astro-ph0404041_arXiv.txt": {
        "abstract": "In many nucleated dwarf elliptical galaxies (dE,Ns), the nucleus is offset by a significant fraction of the scale radius with respect to the center of the outer isophotes.  Using a high-resolution $N$-body simulation, we demonstrate that the nucleus can be driven off-center by the $m=1$ counter-streaming instability, which is strong in flattened stellar systems with zero rotation. The model develops a nuclear offset of the order of one third of the exponential scale-length.  We compare our numerical results with photometry and kinematics of FCC046, a Fornax cluster dE,N with a nucleus offset by 1\\farcs2; we find good agreement between the model and FCC046. We also discuss mechanisms that may cause counter-rotation in dE,Ns and conclude that the destruction of box orbits in an initially triaxial galaxy is the most promising. ",
        "introduction": "\\label{sec:intro}  Dwarf elliptical galaxies (dEs), faint, low-surface brightness galaxies with smooth elliptical isophotes, are the most common type of galaxy in clusters and groups of galaxies (see \\citet{bf94}). In hierarchical models of cosmological structure formation, dwarf galaxies form first and subsequently merge to form larger galaxies. Understanding their formation and evolution is therefore vital to a complete picture of structure formation. There is evidence that dEs constitute a heterogenous class of objects, kinematically (rotating versus unrotating, \\citet{ggv03}), chemically \\citep{mddzh03} and morphologically \\citep{bbj02}. About half of the dEs harbor a central bright nucleus and are called nucleated dEs (dE,Ns). dE,Ns appear to be significantly older than normal ellipticals and non-nucleated dEs \\citep{rs03}; upon tidal stripping, their nuclei have been suggested as sources of both the recently discovered ultra compact dwarfs \\citep{p01} and massive globular clusters like $\\omega$Cen \\citep{g02}. In $\\sim 20\\%$ of Virgo dE,Ns, the nucleus is significantly displaced with respect to the center of the outer isophotes, typically by $\\sim 1''$ ($\\sim 100$~pc) \\citep{bbj00}. There is a tendency for the displacement to increase with decreasing surface brightness but no relation between nuclear displacement and any other structural or environmental parameter has been found.  Various models have been proposed to explain these offset nuclei. In this Letter, we investigate whether the lopsided ($m=1$) counter-streaming instability can reproduce a dE,N with an off-center nucleus. This instability has been known since \\citet{zh78}, and been studied analytically \\citep{s88,pp90} and with $N$-body simulations \\citep{ms90,lds90,sm94,sv97}. \\citet{sv97} do not detect it in systems rounder than E6, which is mainly due to their rounder models being stabilized by a higher radial pressure. On the other hand, \\citet{ms90} find lopsidedness developing in systems as round as E1 but with negligible radial pressure. Partial rotation only introduces a pattern speed in an otherwise purely growing instability; simulations find that the lopsidedness produced by the instability is long-lived.  We use a realistic multi-component $N$-body model to generate a lopsided system which we then compare with observations. In Section \\ref{sec:theo}, we present the $N$-body model and in Section \\ref{sec:obs} we compare it with photometry and kinematics of FCC046, an example of such a dE,N with an offset nucleus. In Section \\ref{sec:disc}, we discuss ways in which counter-rotation may arise in dE,Ns. ",
        "conclusions": "\\label{sec:disc}  If the counter-streaming instability is indeed working in a sizable fraction of the dE,Ns to drive nuclei off-center, it is natural to ask why these dEs have counter-rotation in the first place. The first detections of counter-rotation in galaxy disks \\citep{4550,7217} lead to the hypothesis that the capture of gas with retrograde rotation and subsequent star-formation produced two counter-rotating stellar disks. This hypothesis predicts stellar populations of different ages and metallicities. Generally, the two disks will not have identical masses, so some net rotation is expected; this is, therefore, an unlikely explanation for FCC046.  \\citet{tr00} suggested that counter-rotation is produced when a triaxial halo with an initially retrograde pattern speed slowly changes to a prograde pattern speed.  However, this situation probably does not occur very often. Moreover, this mechanism requires evolution on long timescales $\\sim 10^{10}$~yr and hence requires dE,Ns to be old and unperturbed.  A more likely mechanism is the destruction of box-orbits, which are the backbone of any slowly rotating triaxial mass-distribution, by a growing nucleus. Stars on box orbits come very close to the nucleus and thus can be scattered into loop-orbits; conservation of angular momentum then requires that direct and retrograde loop orbits be equally populated \\citep{ec94,mq98}.  If the nucleus has a mass larger than $2\\%$ of the total mass of the galaxy, evolution towards an axisymmetric shape requires a few crossing times ($\\approx 10^7 - 10^8$~yr), while our simulations show that the counter-streaming instability develops on similar timescales. Late infall of gas \\citep{co03} (perhaps as a result of harassment \\citep{mo96,ma01}) or globular clusters \\citep{lo01} may be responsible for growing the nucleus. If this explanation holds, then the two counter-rotating stellar populations should have equal scale-lengths and chemical properties and nearly equal mass.  Additionally, the nucleus may retain some of its initial angular momentum; we note that the nucleus of FCC046 is indeed rotating.  To summarize:~using $N$-body simulations, we have shown that the counter-streaming instability is a viable explanation for at least the dE,N galaxy FCC046. The models can produce offsets as large as observed and successfully reproduce distinctive features in the surface photometry of FCC046, such as the oscillation of isophote centers.  Although at the limit of what can reliably be extracted from the spectra, we find a tantalizing hint of counter-streaming in the major-axis spectra of this galaxy, which, however, needs to be corroborated at higher spectral resolution."
    },
    "0404/astro-ph0404088_arXiv.txt": {
        "abstract": "{Recent Chandra and XMM X-ray observations of rich clusters of galaxies have shown that the amount of hot gas which is cooling below $\\sim$~1 keV is generally more modest than previous estimates. Yet, the real level of the {\\it cooling flows}, if any, remains to be clarified by making observations sensitive to different temperature ranges. As a follow-up of the {\\sl FUSE} observations reporting a positive detection of the O{\\sc vi} doublet at 1032, 1038\\AA\\ in the cluster of galaxies Abell~2597, which provided the first direct evidence for $\\sim 3\\ 10^5$~K gas in a cluster of galaxies, we have carried out sensitive spectroscopy of two rich clusters, Abell~2029 and Abell~3112 ($z \\sim 0.07$) located behind low HI columns. In neither of these clusters could we detect the O{\\sc vi} doublet, yielding fairly stringent limits of $\\sim$27~M$_\\odot$yr$^{-1}$ (Abell~2029) and $\\sim$25~M$_\\odot$yr$^{-1}$ (Abell~3112) to the cooling flow rates using the 10$^5$--10$^6$~K gas as a tracer. The non-detections support the emerging picture that the cooling-flow rates are much more modest than deduced from earlier X-ray observations. ",
        "introduction": "The concept of cooling flows, arising from simple physical considerations applied to the X-ray data, has provided an important input to the models of galaxy formation ({\\it e.g.,} Fabian 1994; Mathews \\& Brighenti 2003).  On the other hand, the expected outcome, often amounting to $\\geq 10^{10}$\\,M$_\\odot$ of cooled gas inside the cores of many rich clusters, has been rather elusive, despite sensitive searches made in multiple wavebands ({\\it e.g.,} Donahue \\& Voit, 2003 and references therein). Evidence does exist for some such cooler gas in the form of highly extended optical nebulosities seen in the cores of several cooling-flow clusters ({\\it e.g.,} Hu et al. 1985; Heckman 1985), or CO emission (Edge \\& Frayer 2003; Salom\\'e \\& Combes 2003).  However, the masses involved are tiny in comparison with the simple predictions of the cooling-flow rates.  The most recent addition to this intrigue comes from the high resolution spectroscopic observations of several cooling flows with XMM-Newton and Chandra which have revealed a clear deficit of gas at $\\leq$1 keV in comparison with the prediction of simple cooling-flow models ({\\it e.g.,} Peterson et al. 2001, 2003; Kaastra et al. 2001; Tamura et al. 2001).  In this context, spectroscopic observations to search for the UV resonance lines of O\\,{\\sc vi} 1032,1038\\,\\AA\\ emission doublet, a reliable tracer of gas in the temperature range $10^5 - 10^6$K, can play an important role. In fact, using the Far Ultraviolet Spectroscopic Explorer ({\\sl FUSE}), Oegerle et al. (2001) have reported a convincing detection of O\\,{\\sc vi} $\\lambda$1032\\AA\\ line emission from the core of the rich cluster Abell~2597, and have thus estimated a rate of $\\sim$40\\,M$_\\odot$yr$^{-1}$ of intra-cluster medium (ICM) in this cluster cooling through $\\sim$10$^6$\\,K. This rate is only a third of the value inferred from the analysis of ROSAT data (Sarazin et al. 1995). For the other strong cooling-flow candidate, Abell~1795, the ({\\sl FUSE}) observations by Oegerle et al. (2001) placed an upper limit of 28\\,M$_\\odot$yr$^{-1}$ within the central 28~kpc region. It may be noted that the cD galaxies in the cores of both these clusters host twin-jet radio sources extending on a 100~kpc scale and having radio luminosities typical of cluster radio sources in the nearby universe. Also, no O\\,{\\sc vi} has been detected by {\\sl FUSE} in the Virgo and Coma clusters (Dixon et al. 2001a).  Motivated by the above, partly successful attempt using {\\sl FUSE}, we have carried out fairly sensitive {\\sl FUSE} spectroscopy of another two nearby (z$\\sim$0.07) BM type~I clusters of richness class~2, namely, Abell~2029 and Abell~3112. For both clusters, the then available estimates of the cooling-flow rates were mostly in excess of 300\\,M$_\\odot$yr$^{-1}$, based on ROSAT and/or ASCA data (Peres et al. 1998; Sarazin et al. 1998).  A crucial advantage with these two clusters is the exceptionally low H\\,{\\sc i} foreground columns (Table~1), which is an essential pre-requisite for undertaking a sensitive search in the far ultraviolet. In the meanwhile, based on progressively new X-ray observations, the derived properties of these clusters (like many others), such as cooling-flow rates and the radial profile of the ICM temperature, have undergone a sharp ``evolution'' (Sect.~2), making it desirable to employ independent observational probes for addressing this issue.  We report here one such attempt, based on {\\sl FUSE} observations of the two clusters. ",
        "conclusions": "The UV lines of the O\\,{\\sc vi} $\\lambda\\lambda$1032-1038\\AA\\ doublet and/or of the C\\,{\\sc iii} $\\lambda$977\\AA\\ line offer a new tool and independent method to place constraints on the cooling-flow rates in clusters of galaxies. The upper limits for the cooling flow rates in Abell 2029 and Abell 3112, derived here from the UV resonance lines of O\\,{\\sc vi} $\\lambda\\lambda$1032-1038\\AA\\ are 10 to 20 times lower than those estimated earlier from ROSAT observations (Peres et al. 1998). The discrepancy may actually be somewhat smaller in view of the possibility that the cooling flow regions may have been covered only partially within the FUSE aperture. Moreover, the O\\,{\\sc vi} lines may well be weakened due to scattering by dust within the cluster cores (see Oegerle et al. 2001). In any case, our upper limits are fully consistent with those derived recently for the mass deposition rates of the keV gas, using the Chandra and XMM-Newton observations of these clusters (Lewis et al. 2002; Takizawa et al. 2003; see Sect. 2). Thus, the present results indicate a general consistency for two separate temperature regimes of the cooling ICM in these clusters. Begelman \\& Fabian (1990) have discussed the possibility that the $(1-3)\\times 10^5$\\,K gas could be rapidly generated by turbulent mixing of the hot ($10^7$\\,K) and warm ($10^4$\\,K) phases. Thus, provided the intracluster magnetic fields do not effectively suppress the mixing process, the cooling may proceed too rapidly for a strong emission of the Fe\\,{\\sc xvii} X-ray line. UV and X-ray observations therefore provide complementary information on the radiative processes within the intracluster medium.  It may also be noted that the upper limits deduced here for the O\\,{\\sc vi} $\\lambda$1032\\AA\\ line intensities from the two clusters are nearly half the intensity reported for Abell~2597 (Oegerle et al. 1991). Since the redshift of Abell~2597 is slightly larger than those of Abell~2029 and Abell~3112, and Galactic extinction is similar, it is clear that the O\\,{\\sc vi} emission from these two clusters and their corresponding cooling rates through $\\sim 3\\times 10^5$\\,K are intrinsically lower as compared to Abell~2597.    The non-detection of the O\\,{\\sc vi} $\\lambda\\lambda$1032-1038\\AA\\ and C\\,{\\sc iii} $\\lambda$977\\AA\\ lines from the clusters Abell~2029 and Abell~3112, previously believed to have strong cooling-flow rates of several hundred M$_\\odot$yr$^{-1}$ adds to the growing independent evidence that the magnitude of cooling flows in the cores of clusters has been overestimated by at least an order of magnitude.  On a cautionary note we highlight the apparent detection of the O\\,{\\sc vi} doublet at about zero radial velocity in the spectrum taken towards Abell~3112. The apparent emission almost perfectly mimics emission from the Galactic warm gas, similar to that already detected towards a few directions (e.g., Shelton et al. 2001). However, we argue that the O\\,{\\sc vi} emission seen in the spectrum of Abell~3112, in fact, arises from Solar contamination."
    },
    "0404/astro-ph0404277_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "\\label{sec-conclusion} This review of astrophysics at the bottom of the world has encompassed almost a century of human endeavour, beginning with the discovery of the Adelie Land Meteorite in 1912, during the `heroic' age of Antarctic exploration, a period when humans first began to appreciate the vastness and remoteness of the last continent.  It was to be nearly fifty years before the next advances in astronomy were made there, and it is little more than two decades % since Antarctic astrophysics began in earnest.  The 1990's have seen an avalanche of activities, spanning the continent and in particular focussing on the high inland plateau, where the conditions are particularly appealing for a wide range of astronomical observations.  This review has paid particular attention to the development of ``photon'' astronomy in Antarctica, but equally ambitious have been some of the developments in high energy astrophysics, seeking to measure the incident particle fluxes from space.  The scope of this article does not allow for a description of all the projects that have taken place, and a selection of which to portray has had to be made by the authors.  A comprehensive compilation of astronomical experiments that have been conducted in Antarctica is, however, given in Table 1.  This provides some further detail on experiments described here, as well as information on a number of experiments not mentioned in the text.  Antarctica has already produced a range of stunning astronomical results despite the relative infancy of most of the facilities that have been operating.  This is particularly so for the cosmic microwave background experiments, where the stable conditions have facilitated producing results that compare very favourably to those later obtained by satellites, as well as extending the range over which the anisotropies can be probed.  There is no doubt that the summits of the plateau provide superlative conditions for the conduct of a wide range of observational astronomy.  It will be a fascinating story to follow the development of new facilities on the high plateau in the third millennium, and to wonder about the scientific problems they will be able to tackle."
    },
    "0404/astro-ph0404107_arXiv.txt": {
        "abstract": "Using measured X-ray luminosities to 17 Gamma-Ray Bursts (GRBs) during the afterglow phase and accounting for radiative losses, we calculate the kinetic energy of these bursts and investigate its relation to other GRB properties.  We then use the observed radiated energy during the prompt phase to determine the radiative efficiency of these bursts, and explore how the efficiency relates to other GRB observables.  We find that the kinetic energy in the afterglow phase is directly correlated with the radiated energy, total energy as well as possibly the jet opening angle and spectral peak energy. More importantly, we find the intriguing fact that the efficiency is correlated with the radiated energy, and mildly with the total energy, jet opening angle and spectral peak energy. XRF020903 also seems to follow the trends we find for our GRB sample. We discuss the implications of these results for the GRB  radiation and jet models. ",
        "introduction": "In the past several years, there has been a significant increase in our understanding of the energetics of Gamma-Ray Bursts (GRBs).  By measuring the fluence during the prompt phase (usually when most of the energy is emitted), and determining the burst redshift and degree to which the outflow is beamed from afterglow, one can make an estimate of the radiated energy of the burst.  Frail et al. (2001) showed that in fact most GRBs with measured redshifts have  radiated energies distributed very narrowly around the value $5 \\times 10^{50}$ erg. This conclusion is confirmed by Bloom, Frail \\& Kulkarni (2003) with a larger jet data sample. Understanding this energy output is of course crucial in helping constrain progenitor models. However, an often overlooked point in discussing GRB energetics is that the radiated energy is not necessarily the {\\em total} energy of the burst. The radiated energy is some unknown fraction of the initial kinetic energy of the outflow - i.e. the kinetic energy is converted to radiation via some shock (or other) interaction(s) in the burst ejecta.  It is this kinetic energy or, more precisely, the {\\em efficiency} of converting kinetic energy to radiation that must be considered when discussing GRB energetics. This efficiency may help us better understand by what mechanism the kinetic energy becomes radiation. It may also play an important role in understanding so-called ``X-ray Flashes'' (XRFs; GRBs with lower spectral peak energy and flux) - i.e.  XRFs may well be GRBs with much smaller emission efficiency.  To assess the kinetic energy, the most adequate method is through broadband afterglow fits (e.g. Panaitescu \\& Kumar 2001), while another convenient method is to study the X-ray afterglow data alone (Freedman \\& Waxman 2001). Using both methods, it has been found that the kinetic energy - corrected for the degree of beaming of the outflow -  is also narrowly distributed (Panaitescu \\& Kumar 2001; Piran et al. 2001, under the assumption that there is a correlation between X-ray luminosity and beaming corrections ). The result (as well as the validity of the Piran assumption) is confirmed by Berger, Kulkarni \\& Frail (2003) with a larger X-ray afterglow data sample. { Although these works are the first to present calculations of burst kinetic energies, to our knowledge there are  no detailed studies of correlations between kinetic energy or GRB efficiency with other burst properties.}  With these points in mind, we  derive the absolute values of GRB kinetic energy for a large data sample, and investigate the behavior of the kinetic energy and efficiency of GRBs, in relation to other burst properties (such as radiated energy, total energy, characteristic jet angle, and spectral peak energy).  The paper is organized as follows:  In \\S 2, we describe how to compute the kinetic energy and efficiency of GRBs from the observed { X-ray} afterglow data.  We also discuss the basic assumptions and caveats with our method. In \\S 3, we present our results, including the correlation of kinetic energy and efficiency with other burst properties. Our main results shown that both kinetic energy and efficiency are correlated with radiated energy and to some extent, total energy.  There are also suggestions of relations with the jet opening (or viewing) angle and the spectral peak energy, although additional data are needed to confirm this. In \\S 4, we discuss the physical implications of these results, and in \\S 5, present our summary and conclusions. ",
        "conclusions": "We have computed the kinetic energy and radiative efficiency of a sample of 17 GRBs { and 1 XRF, accounting for radiative losses during the initial afterglow phase.}  We have found that both of these GRB properties are related to a number of other GRB observables such as radiated energy, total energy, jet characteristic angle, and spectral peak energy.  We have shown that kinetic energy is directly proportional to the radiated energy of the burst, as well as the total energy and spectral peak energy.  The kinetic energy also decreases with increasing jet characteristic (opening or viewing) angle. { More importantly,} we have also found that the efficiency is correlated with the radiated energy of the burst, which appears to be consistent with an internal shock picture for the GRB prompt phase. The efficiency also appears to increase with increasing total energy and spectral peak energy, and decreases with increasing jet angle. XRF 020903 (the only XRF for which we are able to compute the kinetic energy and efficiency) seems to follow all of the trends we find for the 17 GRBs.  The relation between kinetic energy and radiated energy, efficiency and radiated energy, combined with the previously known Frail and Amati relations can explain the additional correlations we find in our sample. { The results are compatiable with the quasi-universal model, although a large data sample for the GRB kinetic energy derived at an earlier epoch (attainable from the Swift obseratory) is essential to fully test this hypothesis. The results also strongly disfavor a ``narrow beam'' interpretation of GRBs.}  The relationship between efficiency and other GRB parameters is an important one in elucidating GRB physics - we can further explore these and other relations (such as efficiency and variability) and potentially distinguish between various GRB { radiation and jet} models with the launch of the Swift satellite in September, 2004."
    },
    "0404/astro-ph0404331_arXiv.txt": {
        "abstract": "s{We have mapped the AGN luminosity function and its evolution between $z=1$ and $z=5$ down to apparent magnitudes of $R<24$. Within the GEMS project we have analysed HST-ACS images of many AGN in the Extended Chandra Deep Field South, enabling us to assess the evolution of AGN host galaxy properties with cosmic time. } ",
        "introduction": "This article is a report on recent progress in the study of optically faint Active Galactic Nuclei (AGN). Most of the content has recently been published elsewhere; on the following pages we provide a concise summary and present some of the key figures. ",
        "conclusions": ""
    },
    "0404/astro-ph0404457_arXiv.txt": {
        "abstract": "We have revised the calculation of the flux of atmospheric neutrinos based on a 3-dimensional scheme with the realistic IGRF geomagnetic model. The primary flux model has been revised, based on the AMS and BESS observations, and the interaction model updated to DPMJET-III. With a fast simulation code and computer system, the statistical errors in the Monte Carlo study are negligible. We estimate the total uncertainty of the atmospheric neutrino flux prediction is reduced to $\\lesssim$~10~\\% below 10~GeV. The `3-dimensional effects' are found to be almost the same as the study with the dipole magnetic field, but the muon curvature effect remains up to a few tens of GeV for horizontal directions. The uncertainty of the absolute normalization of the atmospheric neutrino is still large above 10~GeV due to the uncertainty of the primary cosmic ray flux above 100~GeV. However, the zenith angle variation is not affected by these uncertainties. ",
        "introduction": "Introduction}  The discovery of the neutrino oscillation from the study of atmospheric neutrinos is a one of the most important results in recent physical research~\\cite{sk} (see also Refs.~\\cite{fukuda,imb,soudan2,macro}, and Ref.~\\cite{KT} for a review). The study is carried out by the comparison of theoretical calculation of the atmospheric neutrino flux and experimental data. Therefore, it is desirable that both theoretical and experimental studies are improved. The SuperKamiokande is improving the statistics and the accuracy steadily for experimental data. It is important to improve the theoretical prediction of the atmospheric neutrino flux also.  There have been some improvements in the theoretical prediction of the atmospheric neutrino flux\\cite{hkkm95,gaisser-new,fluka-battis, engel-hamburg, hkkm-hamburg, hkkm-tsukuba} (see Ref.~\\cite{gaisser-honda} for a review). These studies were useful to determine the flux ratios between different types of neutrinos and the variation over zenith angle with good accuracy, and establish neutrino oscillations and the existence of neutrino masses. We now wish to improve the accuracy of the absolute normalization as well as the ratio and directionality of the atmospheric neutrino fluxes for further studies.  In the time since our last comprehensive study of the atmospheric neutrino flux~\\cite{hkkm95}, knowledge of the primary cosmic ray has been improved by observations such as BESS~\\cite{bess} and AMS~\\cite{ams} below 100~GeV. There have also been theoretical developments in hadronic interaction models such as Fritiof 7.02~\\cite{fritiof7.02}, FLUKA97~\\cite{fluka} and DPMJET-III~\\cite{dpmjet3}. Here, we adopt these revised primary flux and hadronic interaction models.   It has also been pointed out that the atmospheric neutrino flux calculated in a 3-dimensional scheme is significantly different from that calculated in a 1-dimensional scheme at low energies for near horizontal directions~\\cite{fluka-battis,lipari-ge,lipari-ew,hkkm-dipole,wentz,liu,bartol-3d}. The 1-dimensional approximation has been widely used in the past, and was used in our previous calculation and others~\\cite{hkkm95,gaisser-new}. This approximation is justified by the nature of hadronic interactions for calculations of high energy ($\\gtrsim$~10~GeV) atmospheric neutrino fluxes, but not at lower energies. With the computer resources then available, however, it was difficult to complete the calculation of atmospheric neutrino fluxes in a full 3-dimensional scheme within a tolerable length of the time. Some of 3-dimensional calculations employ approximations based on symmetry to circumvent the impact of limited computer resources. In Ref.~\\cite{fluka-battis}, spherical symmetry is assumed, ignoring the magnetic field in the atmosphere, and in our previous 3-dimensional calculation~\\cite{hkkm-dipole} we assumed an axial symmetry and used a dipole geomagnetic field model. Thus, a detailed calculation in a full 3-dimensional scheme without symmetry remains a challenging job.  We have developed a new and fast simulation code for the propagation of cosmic rays in the atmosphere to calculate the atmospheric neutrino flux in a full 3-dimensional scheme without having to assume symmetry. This fast simulation code and a fast computation system allow us to calculate the atmospheric neutrino flux with good accuracy over a wide energy region from 0.1 to a few tens of GeV, as is shown in this paper. The differences between 3-dimensional and 1-dimensional calculation schemes are similar to that we found in the study with a dipole geomagnetic field~\\cite{hkkm-dipole}, and are small above a few GeV. The neutrino flux calculated in the 3-dimensional scheme is smoothly connected to the one calculated in the 1-dimensional scheme at a few tens of GeV. We are therefore able to discuss the atmospheric neutrino flux up to 10~TeV in this paper.  Although progress in our theoretical study of the atmospheric neutrinos flux has been reported partly elsewhere~\\cite{hkkm-hamburg,hkkm-tsukuba}, this is the first comprehensive report since 1995~\\cite{hkkm95}. ",
        "conclusions": "\\end{table}  \\subsection{\\label{sec:general}Zenith angle dependence of the neutrino flux}  The most prominent difference between 3-dimensional and 1-dimensional atmospheric neutrino flux calculations is the horizontal enhancement at low energies. (For the origin of the horizontal enhancement, see Refs.\\cite{gaisser-honda,lipari-ge,hkkm-dipole}) We compare the zenith angle dependences of atmospheric neutrino fluxes calculated in the 3D, 1D, FLUKA, and DIPOLE cases for Kamioka (Fig.~\\ref{fig:kam-zdep}) and North America (Fig.~\\ref{fig:sno-zdep}), integrating over several energy bins and averaging over azimuth angles.  \\begin{figure}[tbh] \\includegraphics[width=15cm]{kam-zdep.eps}% \\caption{\\label{fig:kam-zdep} Zenith angle dependence of the atmospheric neutrino flux at Kamioka in 3 energy bins. For the azimuthal directions, averages are taken. The thick solid lines are for 3D, dotted lines for 1D, dashed lines for FLUKA, and thin solid lines for DIPOLE. } \\end{figure}  \\begin{figure}[tbh] \\includegraphics[width=15cm]{sno-zdep.eps}% \\caption{\\label{fig:sno-zdep} Zenith angle dependence of the atmospheric neutrino flux at North America in 3 energy bins. For the azimuthal directions, averages are taken. The notation for lines is the same as Fig.~\\ref{fig:kam-zdep}. } \\end{figure}  In these figures, we see the horizontal enhancements in 0.1--0.32~GeV and 0.32--1~GeV energy bins for all the 3-dimensional calculations. On the other hand, the flux near horizontal directions is rather smaller than neighboring directions in the 1D case due to the high cutoff rigidities. In the 1--3.2~GeV energy bin, the differences between the calculations are small. To study the difference of zenith angle dependences of neutrino fluxes due to the calculation scheme, we normalize each flux by the omni-directional flux average and depict the ratio to the 3D case in Figs~\\ref{fig:kam-zratio} and \\ref{fig:sno-zratio} as a function of zenith angle.  \\begin{figure}[tbh] \\includegraphics[width=15cm]{kam-zratio.eps}% \\caption{\\label{fig:kam-zratio} The normalized ratio of each flux to the 3D case as a function of zenith angle, for Kamioka. Dotted lines are for 1D, dashed lines for FLUKA, and thin solid lines for DIPOLE. } \\end{figure}  \\begin{figure}[tbh] \\includegraphics[width=15cm]{sno-zratio.eps}% \\caption{\\label{fig:sno-zratio} The ratio of atmospheric neutrino fluxes calculated by 1D, FLUKA, and DIPOLE for North America to that from the 3D case. The notations are the same as for Fig.~\\ref{fig:kam-zratio}. } \\end{figure}  In these figures, we find that the horizontal enhancement still exists in the 1--3.2~GeV energy bin, but that it decreases rapidly with neutrino energy. The difference at near horizontal directions is more than 50~\\% in 0.1--0.32~GeV bin, but it reduces to $\\lesssim$~10\\% in 1--3.2~GeV bin, for all kinds of neutrino.  The differences among the 3-dimensional calculations (3D, FLUKA, DIPOLE) are small, especially that between 3D and DIPOLE. However, the amplitude of the horizontal enhancement in DIPOLE is clearly slightly smaller than that in the 3D calculation. This is thought to be due to the difference of interaction model, especially to the average transverse momentum of secondary mesons. Note, the $<P_t>$ of pions is 0.289~GeV$/$c in DPMJET-III, while it is 0.256~GeV$/$c in Fritiof~1.6 for P~+~Air interactions at $E_{lab}=$10~GeV. The difference between 3D and FLUKA is asymmetric below and above the horizontal direction ($\\cos\\theta = 0$). It is difficult to deduce differences in the interaction model in this comparison.  \\begin{figure}[tbh] \\includegraphics[width=10cm]{kam-2dir.eps}% \\caption{\\label{fig:kam2dir} Atmospheric neutrino fluxes at Kamioka for 3D and 1D, averaged over azimuth angles. The left panel shows the comparison for vertical directions, and the right panel for horizontal directions. The solid lines show the 3D results and the dashed lines the 1D results. } \\end{figure}  \\begin{figure}[tbh] \\includegraphics[width=10cm]{sno-2dir.eps}% \\caption{\\label{fig:sno2dir} Atmospheric neutrino fluxes at North America for 3D and 1D, averaged over azimuth angles. The left panel shows the comparison for vertical directions, and the right panel for horizontal directions. The solid lines show the 3D results and the dashed lines the 1D results. } \\end{figure}  To see the energy dependence of the horizontal enhancement more clearly, we compared the 3D and 1D energy spectra for vertical and horizontal directions averaging over azimuth angles in Figs.~\\ref{fig:kam2dir} (Kamioka) and \\ref{fig:sno2dir} (North America). We find the differences disappear at $\\sim$1~GeV for vertical directions, and $\\sim$3~GeV for horizontal directions, for all neutrinos.  Moreover, we find the fluxes averaged over all directions for 3D and 1D cases are very close to each other, even at low  energies. Averaging the neutrino fluxes in Figs ~\\ref{fig:kam-zdep} and \\ref{fig:sno-zdep} over zenith angles, the 3D/1D ratios are tabulated in table~\\ref{tab:alldirratio}. They agree with each other to within a few \\% in all cases.  \\begin{table}[!htb] \\caption{3D/1D ratio for the all-direction average of atmospheric neutrino flux.} \\vspace*{2truemm} \\begin{tabular}{|c| l l l l |}% \\toprule \\hspace{7mm}E$_\\nu$(GeV)\\hspace{7mm} &\\hspace{7mm}$\\nu_\\mu$\\hspace{7mm}&\\hspace{7mm}$\\bar\\nu_\\mu$\\hspace{7mm}&\\hspace{7mm}$\\nu_e$\\hspace{7mm}&\\hspace{7mm}$\\bar\\nu_e$\\\\ % \\hline \\multicolumn{5}{|c|}{Kamioka} \\\\ \\hline 0.1 -- .32 & 0.979 & 0.980 & 0.970 & 0.978 \\\\ .32 -- 1.0 & 0.997 & 1.000 & 0.999 & 0.992 \\\\ 1.0 -- 3.2 & 0.983 & 0.984 & 0.982 & 0.975 \\\\ \\hline \\multicolumn{5}{|c|}{North America} \\\\ \\hline 0.1 -- .32 & 1.036 & 1.035 & 1.028 & 1.025 \\\\ .32 -- 1.0 & 1.019 & 1.020 & 1.021 & 1.014 \\\\ 1.0 -- 3.2 & 0.992 & 0.990 & 0.989 & 0.985 \\\\ \\botrule \\end{tabular} \\label{tab:alldirratio} \\end{table}  \\subsection{\\label{sec:eastwest}East--West effect}  Here, we use the 3D and 1D fluxes only, since they are calculated under the same conditions (except for the 1- or 3-dimensional calculation scheme). Contrary to the quantitative agreements between 3D and 1D above a few GeV in the azimuthally averaged fluxes, they are quite different when the azimuth angles are limited to East or West directions, even at higher energies ($\\sim$~10~GeV).  We depict the 3D and 1D  atmospheric neutrino fluxes arriving horizontally ($0. < \\cos\\theta < 0.1$) from the East ($60^\\circ <\\phi <120^\\circ$), and West ($240^\\circ <\\phi <300^\\circ$) for Kamioka (Fig.~\\ref{fig:kam-horizontal}) and North America (Fig.~\\ref{fig:sno-horizontal}), where we measure the azimuth angle from South ($\\phi=0$), and $\\phi =90^\\circ, 180^\\circ, 270^\\circ$ are East, North, and West directions respectively.  \\begin{figure}[tbh] \\includegraphics[width=10cm]{kam-horizontal.eps}% \\caption{\\label{fig:kam-horizontal} The flux of horizontal ($0. < \\cos\\theta < 0.1$) atmospheric neutrinos arriving from the East ($60^\\circ <\\phi<120^\\circ$, Left), and West ($240^\\circ <\\phi<300^\\circ$, Right) directions at Kamioka. The solid lines show the result of the 3D calculation and the dashed line the 1D case. } \\end{figure}  \\begin{figure}[tbh] \\includegraphics[width=10cm]{sno-horizontal.eps}% \\caption{\\label{fig:sno-horizontal} The flux of horizontal ($0. < \\cos\\theta < 0.1$) atmospheric neutrinos arriving from the East ($60^\\circ <\\phi<120^\\circ$, Left), and West ($240^\\circ <\\phi<300^\\circ$, Right) directions at North America. The solid lines show the result of the 3D calculation and the dashed line the 1D case. } \\end{figure}  The differences in the fluxes from the 3D and 1D calculations for different kinds of neutrinos may be classified into two groups, $\\nu_\\mu$ and $\\bar\\nu_e$, and $\\bar\\nu_\\mu$ and $\\nu_e$. The former group are the decay products of $\\mu^-$, and the latter are the decay products of $\\mu^+$. The geomagnetic field deflects $\\mu^+$'s toward the same direction as for primary cosmic rays. Therefore, it enhances the East and West differences of the $\\bar\\nu_\\mu$ and $\\nu_e$ fluxes caused by the rigidity cutoff. On the other hand, the geomagnetic field works on the $\\mu^-$'s in the opposite direction to that for primary cosmic rays. Therefore, it reduces the East and West differences for the $\\nu_\\mu$ and $\\bar\\nu_e$ fluxes caused by the rigidity cutoff~\\cite{lipari-ew}. In the Figs.~\\ref{fig:kam-horizontal} and \\ref{fig:sno-horizontal}, we mainly see the horizontal enhancement in the neutrino energies $\\lesssim$~1~GeV in the difference of 3D and 1D. For $\\gtrsim$ 1~GeV, however, the muon curvature in the geomagnetic field is a larger effect than the horizontal enhancement, and this extends to several tens of GeV for near horizontal directions.  \\begin{figure}[tbh] \\includegraphics[width=15cm]{kam-adep.eps}% \\caption{\\label{fig:kam-adep} Azimuth angle dependence of the atmospheric neutrino flux at Kamioka, averaged over the zenith angle range $-0.5 < \\cos\\theta < 0.5$. The solid lines show the 3D result and the dashed lines the 1D result. } \\end{figure}  \\begin{figure}[tbh] \\includegraphics[width=15cm]{sno-adep.eps}% \\caption{\\label{fig:sno-adep} Azimuth angle dependence of the atmospheric neutrino flux at North America, averaged over the zenith angle range $-0.5 < \\cos\\theta < 0.5$. The solid lines show the 3D result and the dashed lines the 1D result. } \\end{figure}  \\begin{table}[!htb] \\caption{Max/Min ratio in Figs.~\\ref{fig:kam-adep} and ~\\ref{fig:sno-adep} as the amplitude of the azimuthal angle variation.} \\vspace*{2truemm} \\begin{tabular}{|c| l l l l | l l l l |}% \\toprule E$_\\nu$(GeV) &$\\nu_\\mu$&$\\bar\\nu_\\mu$&$\\nu_e$&$\\bar\\nu_e$ &$\\nu_\\mu$&$\\bar\\nu_\\mu$&$\\nu_e$&$\\bar\\nu_e$ \\\\ % \\hline &\\multicolumn{4}{c|}{Kamioka, 3D}&\\multicolumn{4}{c|}{Kamioka, 1D} \\\\ \\hline 0.1 -- .32 & 2.27 & 2.51 & 2.75 & 2.03 & 4.13 & 4.13 & 4.36 & 3.85\\\\ .32 -- 1.0 & 2.27 & 2.82 & 3.22 & 1.98 & 2.73 & 2.76 & 2.98 & 2.65\\\\ 1.0 -- 3.2 & 1.58 & 2.00 & 2.42 & 1.42 & 1.61 & 1.63 & 1.80 & 1.61\\\\ \\hline &\\multicolumn{4}{c|}{North America, 3D}&\\multicolumn{4}{c|}{North America, 1D} \\\\ \\hline 0.1 -- .32 & 1.26 & 1.33 & 1.40 & 1.19 & 1.47 & 1.48 & 1.54 & 1.41\\\\ .32 -- 1.0 & 1.20 & 1.39 & 1.49 & 1.11 & 1.27 & 1.28 & 1.31 & 1.26\\\\ 1.0 -- 3.2 & 1.07 & 1.25 & 1.35 & 1.07 & 1.08 & 1.09 & 1.12 & 1.09\\\\ \\botrule \\end{tabular} \\label{tab:adep-amplitude} \\end{table}  The muon curvature effect should be seen in the azimuthal variation of the atmospheric neutrino fluxes. We show the integrated azimuthal angle dependence in the same energy bins as in Sec.~\\ref{sec:general} for Kamioka (Fig.~\\ref{fig:kam-adep}) and North America (Fig.~\\ref{fig:sno-adep}). Also we tabulated the ratio of maximum to minimum fluxes in the figures in Table~\\ref{tab:adep-amplitude}, to see the variation amplitude. In the 1D case, the amplitudes are similar to each other for all kinds of neutrinos in each energy bin. This is because the 1D azimuthal variation  is caused by the rigidity cutoff of the primary cosmic rays. In the 3D case, the amplitudes are different among the different kinds of neutrino even in the same energy bin. The amplitudes of $\\nu_\\mu$ and $\\bar\\nu_e$ are suppressed, while those of $\\bar\\nu_\\mu$ and $\\nu_e$ are enhanced, except for the lowest energy bin of 0.1--0.32~GeV. In the lowest energy bins, smearing suppresses the 3D azimuth angle dependence.  Note that $\\nu_e$ has the largest amplitude among all $\\nu$'s in 3D. This is because about $1/2$ of the $\\bar\\nu_\\mu$'s are created by pion decay directly, while all the $\\nu_e$'s are created by $\\mu$-decay at these energies. It is noteworthy that the amplitude of $\\nu_e$ in the 1--3.2~GeV energy bin is still large. This is important for the experimental confirmation of the effect of muon curvature, because the determination of the arrival direction is better for higher energy neutrinos.  Generally speaking, the azimuthal angle dependence of atmospheric neutrinos at high magnetic latitude sites such as North America is smaller than that at low magnetic latitude sites such as Kamioka because the rigidity cutoff is too low. However, we find in Fig.~\\ref{fig:sno-adep} and Table~\\ref{tab:adep-amplitude} that the difference in the azimuthal angle dependence of atmospheric neutrinos between 3D and 1D due to the muon curvature is similar to that for the low magnetic latitude site.  Summary and discussion}  We have revised the calculation of the atmospheric neutrino flux, according to recent developments in primary cosmic ray observations and hadronic interaction models. We have also updated the calculations from a 1-dimensional scheme to a 3-dimensional one. For the interaction model, we compared the available interaction models with the secondary cosmic rays observed at balloon altitudes, and selected DPMJET-III as the preferred model for this study. We have constructed an inclusive hadronic interaction code based on DPMJET-III for speed. The computation speed is very important in the 3-dimensional calculation. We have processed 307,618,204,971 cosmic rays with $E_k/A > 1$~GeV for lower energy neutrino fluxes, and 415,711,823,606 cosmic rays with $E_k/A > 10$~GeV for neutrino fluxes above 10~GeV. Combining both simulations, the statistical error due to the Monte Carlo method is negligibly small for neutrino energies below a few tens of GeV.  With the primary fluxes accurately determined by BESS and AMS below 100~GeV, the DPMJET-III interaction model, and the fast 3-dimensional simulation code for the cosmic ray propagation, we consider we have reduced the uncertainty to $\\sim$~10~\\%  in the calculation of the atmospheric neutrino flux at energies below 10~GeV, since we could reproduce the muon fluxes at various altitude with good accuracy from 1 to a few 10~GeV~\\cite{abe2,sanuki-mu} in this calculation scheme. However, for the atmospheric neutrino flux above 10~GeV, the uncertainties in the atmospheric neutrino fluxes are still large due to the uncertainties of the primary cosmic ray flux and interaction model above 100~GeV.  The differences we find between the 1- and 3-dimensional calculation schemes are similar to those we found in the previous study with a dipole geomagnetic field. When we average the atmospheric neutrino flux over azimuthal angles, the fluxes calculated in the 3-dimensional scheme quickly converge with those calculated in the 1-dimensional scheme at a few GeV. With the larger statistics, however, it becomes clear that muon curvature affects the horizontal neutrino flux even at energies $\\gtrsim$~10~GeV, while most other `3-dimensional effects' disappear at a few GeV.  In comparison with other calculations of atmospheric neutrino flux, we find the zenith angle dependences of different calculations are very similar, although there are differences in the absolute values. The remaining differences of the zenith angle dependence at higher energies ($\\gtrsim$ 10~GeV) are consistent with the uncertainty of kaon production in the hadronic interaction~\\cite{noon2003}. Therefore, we may conclude that the main reason for the remaining difference of the zenith angle dependence is in the kaon production of hadronic interaction model used by different calculations. Note, there are large differences in the interaction model used by the different calculations, as is known from the large $\\nu_e/\\bar\\nu_e$ and $\\nu_\\mu/\\bar\\nu_mu$ differences at higher energies.  The production height distributions in the 1- and 3-dimensional calculation schemes are different depending on the arrival direction. When they are averaged over azimuthal directions, they agree with each other well above a few GeV, except for a small distortion at the lower tail for the vertical directions. This situation is similar to that of the flux value. There are azimuthal variations of the production height due to the muon curvature, however it is difficult to study them in detail with the statistics of this work. However, for the experimental study of atmospheric neutrinos for neutrino oscillations, the azimuthal variations are not important.  It is interesting that the effect of albedo particles observed by AMS at satellite altitudes is seen in the neutrino production time distribution. The contribution of such particles to the atmospheric neutrino flux is a little higher for the low magnetic latitude site (Kamioka) than the high magnetic latitude site (North America), but small ($ \\ll 1$~\\%) for both sites.  We expect that the validity of the calculation scheme and the effect of the muon curvature will be confirmed by the observation of the azimuthal variation of the neutrino events. Although the statistics are still insufficient, the SuperKamiokande experiment observed a larger amplitude of the azimuthal variation for the e-like events than that for the $\\mu$-like events~\\cite{futagami,moriyama} as is predicted in Sec.~\\ref{sec:eastwest}. We would like to note that the muon curvature effect has been confirmed by the azimuthal variation of the muon flux with an amplitude almost the same as the value predicted in our calculation scheme~\\cite{hanoi-mu}."
    },
    "0404/astro-ph0404511_arXiv.txt": {
        "abstract": "We have computed the higher-order Balmer absorption line indices \\Hg\\ and \\Hd\\ \\citep{WO97} for stellar population models with variable element ratios. The response of these indices to abundance ratio variations is taken from detailed line formation and model atmosphere calculations. We find that \\Hg\\ and \\Hd, unlike \\Hb, are very sensitive to \\aFe\\ ratio changes at super-solar metallicities. Both line indices increase significantly with increasing \\aFe\\ ratio. This effect cannot be neglected when these line indices are used to derive the ages of metal-rich, unresolved stellar populations like early-type galaxies. We re-analyze the elliptical galaxy sample of \\citet{Kun00}, and show that consistent age estimates from \\Hb\\ and \\Hg\\ are obtained, only if the effect of \\aFe\\ enhancement on \\Hg\\ is taken into account in the models. This result rectifies a problem currently present in the literature, namely that \\Hg\\ and \\Hd\\ have up to now led to significantly younger ages for early-type galaxies than \\Hb. Our work particularly impacts on the interpretation of intermediate to high-redshift data, for which only the higher-order Balmer lines are accessible observationally. ",
        "introduction": "\\label{sec:intro} Absorption line indices in the visual as defined by the Lick group (e.g., \\Hb, \\Mgtwo, Fe5270, Fe5335, etc., \\citealt{Buretal84,Fabetal85}) have proven to be a useful tool for the derivation of ages and metallicities of unresolved stellar populations. One of the largest merits of the Lick indices is to have signaled the presence of non-solar Mg/Fe abundance ratios in the stars of early-type galaxies \\citep*{WFG92,DSP93,CD94,FFI95,SB95,WPM95,Greggio97}.  This interpretation is confirmed empirically by the similarity between the Lick indices of early-type galaxies and those of Bulge globular clusters \\citep{Maretal03}, which are known from spectroscopy of individual stars to be Mg/Fe enhanced.  Indeed the indices \\Mgtwo/\\Mgb\\ and Fe5270/Fe5335 are dominated by absorption lines from the elements Mg and Fe, and model atmosphere calculations by \\citet{Bar94} and \\citet{TB95} have theoretically demonstrated that these indices are sensitive to Mg/Fe abundance ratios. Following an extension of the method introduced by \\citet{Traetal00a}, we have used the calculations of \\citet{TB95} to produce stellar population models with variable element abundance ratios (\\citealt*{TMB03a}, hereafter TMB03). These models are now able to match simultaneously the Mg and Fe line indices of globular clusters and early-type galaxies.  The Balmer line index \\Hb, which is used as age indicator because it measures the presence of warm A-type stars, is only moderately contaminated by metallic lines and therefore relatively insensitive to abundance ratio variations \\citep[][TMB03]{TB95,Traetal00a}. Does this convenient attribute of \\Hb\\ apply also to the higher-order Balmer line indices \\Hg\\ and \\Hd\\ defined in \\citet{WO97}?  The comparison of the predictions from solar-scaled stellar population models with galactic globular cluster data shows that the effect from \\aFe-enhancement must be small in a metallicity range up to solar \\citep{Maretal03}.  There are several good reasons for serious doubts, however, that this is the case also at super-solar metallicities.  1) From data in the literature (\\citealt{KD98}; \\citealt{Teretal99}; \\citealt{Pogetal01a}; \\citealt{Kunetal02a}; \\citealt*{TFD04}) it can be seen that the higher-order Balmer line indices and \\Hb\\ lead to very inconsistent age estimates.  2) The \\Hd\\ measurements of early-type galaxies in the Sloan Digital Sky Survey cannot be reproduced by current solar-scaled stellar population models \\citep{Eisetal03}.  3) Both \\Hg\\ and \\Hd\\ have very prominent Fe absorption features in their pseudo-continua.  \\citet{TB95} have not included \\Hg\\ and \\Hd, however, so that a quantitative assessment of the sensitivity of these line indices to abundance ratios in stellar population models was not possible until now.  Since these line indices are already being widely used and will certainly supersede \\Hb\\ as age indicators in future studies at intermediate and high-redshifts, it is urgent to clarify this issue. To this aim in Korn et al.\\ (in preparation) we have extended the \\citet{TB95} approach and computed model atmosphere calculations with variable abundance ratios including the wavelengths of the higher-order Balmer lines ($\\lambda\\sim 4000\\, {\\rm \\AA}$).  In this paper we present the resulting stellar population models of \\Hg\\ and \\Hd\\ with different \\aFe\\ ratios, which are computed following the recipe of TMB03.  The paper is organised as follows. In Section~2 we will briefly introduce the stellar population model, and calibrate it with galactic globular clusters. In Section~3 we shall confront the new models with data of elliptical galaxies focusing on the issues outlined above. We will conclude with Section~4. ",
        "conclusions": "We have extended our models of absorption line indices of stellar populations with variable abundance ratios (TMB03) to the higher-order Balmer line indices \\HgA, \\HgF, \\HdA, and \\HdF. We show that the model predictions are well consistent with data of galactic globular clusters.  The key result of this work is that, unlike \\Hb, all four indices show a marked dependence on the \\aFe\\ ratio. The significance of this effect increases with increasing metallicity, and therefore impacts on the age derivation for elliptical galaxies.  Re-analyzing the elliptical galaxy sample of \\citet{Kun00}, we show that \\Hb\\ and \\HgF\\ (and a bit less convincingly also \\HgA) yield consistent age estimates if (and only if) the effect of \\aFe\\ enhancement is taken into account.  \\medskip A posteriori, it is not surprising that blue absorption line indices are more affected by abundance ratio effects, because of the presence of more metallic lines in the bluer parts of the spectrum. This fact certainly diminishes the advantage of the higher-order Balmer line indices over \\Hb\\ as age indicators, in spite of their lower sensitivity to emission line filling.  The consequence for the age derivation by means of \\Hg\\ and \\Hd\\ is twofold: 1) The element ratio effect on the Balmer line indices needs to be taken into account in the models.  2) The additional knowledge of the \\aFe\\ ratio is necessary, which requires the simultaneous measurement of \\Mgb\\ and Fe5270/Fe5335. If these wavelengths are not accessible, as typically the case for intermediate and high-redshift data, \\aFe\\ ratios may be best estimated from the indices Fe4383 and \\CNone\\ or \\CNtwo.  \\medskip Model tables are listed in the Appendix and are available electronically under http://www.mpe.mpg.de/$\\sim$dthomas."
    },
    "0404/astro-ph0404393.txt": {
        "abstract": "We present an extensive new time-series of spectroscopic data of the peculiar SN~1999aa in NGC 2595. Our data set includes 25 optical spectra between $-$11 and +58 days with respect to B-band maximum light, providing an unusually complete time history. The early spectra resemble those of a SN~1991T-like object but with a relatively strong Ca H\\&K absorption feature. The first clear sign of Si~{\\sc ii}~$\\lambda6355$, characteristic of Type Ia supernovae, is found at day $-$7 and its velocity remains constant up to at least the first month after B-band maximum light. The transition to normal-looking spectra is found to occur earlier than in SN~1991T suggesting SN~1999aa as a possible link between SN~1991T-like and Branch-normal supernovae.  Comparing the observations with synthetic spectra, doubly ionized Fe, Si and Ni are identified at early epochs. These are characteristic of SN~1991T-like objects.  Furthermore, in the day $-$11 spectrum, evidence is found for an absorption feature which could be identified as high velocity C~{\\sc ii}~$\\lambda$6580 or H$\\alpha$. At the same epoch C~{\\sc iii}~$\\lambda$4648.8 at photospheric velocity is probably responsible for the absorption feature at 4500~\\AA.  High velocity Ca is found around maximum light together with Si~{\\sc ii} and Fe~{\\sc ii} confined in a narrow velocity window. Implied constraints on supernovae progenitor systems and explosion hydrodynamical models are briefly discussed. ",
        "introduction": "\\label{sec_intro} The observed homogeneity and brightness of Type Ia supernovae (SNe~Ia) make them excellent tools for distance estimates over extremely large distances and hence for measurements of cosmological parameters (see e.g. \\markcite{1998Natur.391...51P,1998ApJ...493L..53G,1998ApJ...507...46S,1998AJ....116.1009R,1999ApJ...517..565P,2003ApJ...598..102K,2003ApJ...594....1T}{Perlmutter} {et~al.} (1998); {Garnavich} {et~al.} (1998); {Schmidt} {et~al.} (1998); {Riess} {et~al.} (1998); {Perlmutter} {et~al.} (1999); {Knop} {et~al.} (2003); {Tonry} {et~al.} (2003)). However, cosmological results derived from supernovae rely on the evidence that distant explosions are similar to well-studied nearby ones and that they can be calibrated with the same techniques, e.g. the luminosity-light curve-timescale relation \\markcite{1993ApJ...413L.105P}({Phillips} 1993).  The Supernova Cosmology Project (SCP) coordinated an extensive campaign to study a large number of z $<$ 0.1 SNe~Ia in the spring of 1999 in order to better understand the intrinsic properties of SNe Ia and thereby improve cosmological measurements using them \\markcite{2000coex.conf...75A,2000sgrb.conf...47N}({Aldering} 2000; {Nugent}, {Aldering}, \\& {The Nearby  Campaign} 2000). The subject of this work, SN~1999aa, was one of the SCP targets in that campaign.  At this time, several fundamental questions about Type Ia supernova physics remain. The nature of the progenitor system is still poorly constrained as are the details of the explosion and thus the origin of many of the differences observed among Type Ia supernovae.  Normal SNe~Ia (sometimes called Branch-normal, \\markcite{1983ApJ...270..123B,1993AJ....106.2383B}{Branch} {et~al.} (1983); {Branch}, {Fisher}, \\& {Nugent} (1993)) present early spectra dominated by intermediate mass elements (IMEs), such as Mg, Si, S, and Ca, which are replaced by features due to iron-peak ions (such as Fe, Co, and Ni) as the spectrum evolves with time. However, peculiar events such as SN~1991T, SN~1997br and SN~2000cx \\markcite{1992ApJ...384L..15F,1992ApJ...397..304J,1992AJ....103.1632P,1992ApJ...387L..33R}({Filippenko} {et~al.} 1992; {Jeffery} {et~al.} 1992; {Phillips} {et~al.} 1992; {Ruiz-Lapuente} {et~al.} 1992) show different characteristics. Their early spectra have weak IME lines and strong doubly ionized iron lines while spectra after maximum light look almost completely normal. Moreover, their light curves are generally characterized by a slow post maximum decline rate, and thus, SN~1991T-like supernovae are sometimes called peculiar slow decliners. These characteristics have been regarded as possible signs of different classes of progenitors.  SN~1999aa exhibited spectral characteristics common both to Branch-normal and peculiar slow decliner SNe~Ia, and thus has been proposed as a key object which may help in understanding the physical origin of the observed diversity \\markcite{Branch:2000zm,Branch:2001jf,2001ApJ...546..734L}(Branch 2000; Branch, Baron, \\& Jeffery 2001; {Li} {et~al.} 2001b).  Hydrodynamical models predict atmospheric compositions of supernovae in good agreement with what is found in observed spectra. None of the currently available models, though, produces an exhaustive description of the whole spectral time evolution and none is able to reproduce the complete range of observed differences among SNe~Ia. For extensive reviews of the theoretical models and observations see \\markcite{1997ARA&A..35..309F,2000ARA&A..38..191H,2000A&ARv..10..179L,Livio:2000cx,Branch:2001jf}{Filippenko} (1997); {Hillebrandt} \\& {Niemeyer} (2000); {Leibundgut} (2000); Livio (2000); Branch {et~al.} (2001).  The most widely accepted model for Type Ia supernovae involves the thermonuclear disruption of a C+O white dwarf star (WD) accreting material from a companion star \\markcite{1973ApJ...186.1007W,1982ApJ...257..780N,1984ApJS...54..335I,1985cvlm.proc....1P}({Whelan} \\& {Iben} 1973; {Nomoto} 1982; {Iben} \\& {Tutukov} 1984; {Paczynski} 1985). From an observational point of view this conclusion is supported by the amount of energy released, the absence of hydrogen lines (but see \\markcite{2000ApJS..128..615M}{Marietta}, {Burrows}, \\&  {Fryxell} (2000)) and the occurrence of SNe~Ia in elliptical galaxies exclusive of other types.  A more controversial issue is whether SNe~Ia are the result of explosions at Chandrasekhar mass \\markcite{1969Ap&SS...5..180A,1969Ap&SS...3..464H,1976Ap&SS..39L..37N,1991A&A...245..114K,1994ApJ...423..371W}({Arnett} 1969; {Hansen} \\& {Wheeler} 1969; {Nomoto}, {Sugimoto}, \\&  {Neo} 1976; {Khokhlov} 1991; {Woosley} \\& {Weaver} 1994) or at sub-Chandrasekhar mass \\markcite{1990ApJ...354L..53L,1991ApJ...370..272L,1994ApJ...423..371W}({Livne} 1990; {Livne} \\& {Glasner} 1991; {Woosley} \\& {Weaver} 1994). In the former, thermonuclear burning of carbon occurs in proximity to the center of the star and the burning front proceeds outward. In the latter, helium accreting in the external layer of the supernova, ignites. A detonation then propagates outward through the He layer and another inward compressing the C+O nucleus that ignites off-center.  While Chandrasekhar mass models have been successful in reproducing many of the observed characteristics of Branch-normal and SN~1991T-like supernovae \\markcite{1997ApJ...485..812N,1999MNRAS.304...67F}({Nugent} {et~al.} 1997; {Fisher} {et~al.} 1999), sub-Chandrasekhar models are in good agreement with fainter explosions such as SN~1991bg-like objects \\markcite{1997ApJ...485..812N}({Nugent} {et~al.} 1997).  Pre-maximum spectra can discriminate among these two scenarios since sub-Chandrasekhar models have external layers dominated by He and Ni and do not leave any unburned carbon or produce IMEs at expansion velocities above 14000 km~s$^{-1}$. The identification of C, or strong Ca or Si lines in pre-maximum spectra would then rule out this possibility.  Most of the supernovae observed seem to show characteristics in agreement with  carbon ignition occurring at the center of the WD. However, this condition can be the result of both a single degenerate C+O WD accreting hydrogen from a companion or a merging double degenerate C+O WD, as sometimes suggested for SN~1991T-like supernovae \\markcite{1999MNRAS.304...67F}({Fisher} {et~al.} 1999).  The detection of narrow hydrogen lines in supernova spectra \\markcite{2003Natur.424..651H}({Hamuy} {et~al.} 2003) (see also a discussion in \\markcite{2003ApJ...594L..93L}{Livio} \\& {Riess} (2003)) would favor the single degenerate scenario but a non-detection cannot rule out these models. A substantial amount of hydrogen can be removed from the companion star and get mixed into the exploding WD \\markcite{2000ApJS..128..615M}({Marietta} {et~al.} 2000). In this case the detection of H lines is expected even in low resolution spectra \\markcite{2002ApJ...580..374L}({Lentz} {et~al.} 2002) and would then exclude a double degenerate progenitor system.  While single degenerate models are currently considered the most promising, other questions remain open. The hydrodynamics of the explosion and the details of the flame propagation pose many numerical and conceptual problems currently under investigation, e.g. computational resolution in 3D simulations, and details of flame instabilities, for extensive discussion see e.g. \\markcite{Blinnikov:2002gu,2000ARA&A..38..191H}{Hillebrandt} \\& {Niemeyer} (2000) and {Blinnikov} \\& {Sorokina} (2002).  Several models have been proposed to describe the explosion mechanism. Next, we highlight the possible spectroscopic observables that could help in constraining the parameters of the models.  Pure one-dimensional (1D) deflagration models (such as W7, \\markcite{1984ApJ...286..644N}{Nomoto}, {Thielemann}, \\&  {Yokoi} (1984)) have carbon present down to v~$>$~14900~km~s$^{-1}$ while 1D delayed-detonation (DD) models \\markcite{1991A&A...245..114K,1992ApJ...393L..55Y,1994ApJ...423..371W}({Khokhlov} 1991; {Yamaoka} {et~al.} 1992; {Woosley} \\& {Weaver} 1994) burn C almost completely up to v~$\\sim$~30000~km~s$^{-1}$. Thus, the identification of carbon lines at low velocity would disfavor published DD models. Delayed-detonation models can also produce an unusually high Doppler blue shift of IMEs by tuning the deflagration to detonation transition density \\markcite{2001ApJ...547..402L}({Lentz} {et~al.} 2001).  Thus, IME lines confined at higher than normal velocities could be more naturally described by DD models. Lines of stable Fe and Ni at high velocity in the spectra prior to maximum are also consistent with the prediction of a deflagration to detonation transition \\markcite{1999ApJS..125..439I}({Iwamoto} {et~al.} 1999).  The complexity of explosion models has advanced to the level where deflagration in 3D can be explored \\markcite{Khokhlov:2000gj,2003Sci...299...77G}(Khokhlov 2000; {Gamezo} {et~al.} 2003). These results motivated a few attempts to investigate the spectral outcome of these models by means of parametrized synthetic spectra codes \\markcite{2002ApJ...567.1037T,2003ApJ...588L..29B,2003ApJ...593..788K}({Thomas} {et~al.} 2002; {Baron}, {Lentz}, \\& {Hauschildt} 2003; {Kasen} {et~al.} 2003). In 3D deflagration, due to the highly convoluted turbulent flame propagation, heavy mixing of freshly synthesized and unburned material can occur \\markcite{2003Sci...299...77G}({Gamezo} {et~al.} 2003).  Evidence for C and O lines at low velocities, as well as clumps of burned material (e.g: IME or Fe) in the external layers (high velocity), would confirm the plausibility of, and the need for, such a degree of complexity.  The same mechanism that mixes carbon and oxygen with IME in 3D calculations could probably also work for hydrogen.  Polarization measurements in SNe~Ia spectra \\markcite{1996AAS...189.4510W,2001ApJ...556..302H,2003ApJ...593..788K,2003ApJ...591.1110W}({Wang}, {Wheeler}, \\&  {Hoeflich} 1996; {Howell} {et~al.} 2001; {Kasen} {et~al.} 2003; {Wang} {et~al.} 2003) have strengthened the conviction that some degree of asymmetry can be found in supernova atmospheres. In particular, the high velocity (HV) component of the Ca~{\\sc ii} IR triplet found in some supernovae (e.g: SN~2000cx, SN~2001el) required 3D simulations to be fully reproduced \\markcite{Thomas:2003ip,2003ApJ...593..788K}(Thomas {et~al.} 2003; {Kasen} {et~al.} 2003).  To explore the results of 3D explosion models, full three-dimensional radiative transfer calculations of supernova spectra would be required. However, these are not yet available and 1D parametrized radiative transfer calculations are still in their infancy. Thus, direct analysis by means of parametrized codes -- both in 1D and in 3D -- remains the fastest and most versatile way of testing models predictions and guiding new developments.  In the present work we present a new, comprehensive spectroscopic timeseries for SN~1999aa.  We use the direct analysis code SYNOW \\markcite{1990sjws.conf..149J,1997ApJ...481L..89F,1999MNRAS.304...67F}({Jeffery} \\& {Branch} 1990; {Fisher} {et~al.} 1997, 1999) as a tool to describe the data and to identify potentially interesting features. This, together with the measurements of the velocities inferred from the minima of the spectral features, is used to investigate the structure of the expanding atmosphere. In particular, we have looked for evidence of carbon, oxygen, and hydrogen lines in early spectra as well as the velocity ranges of iron, nickel and intermediate mass element lines, in order to try to answer some of the open questions outlined above.  This paper is organized as follows. The spectroscopic data of SN 1999aa, and a short description of the data reduction scheme are presented in section \\ref{sec_data}. The analysis methodology is introduced in section \\ref{modeling}. In section \\ref{sec_comparison} our SN~1999aa spectra are compared with those of spectroscopically peculiar and normal supernovae taken from the literature. The synthetic spectra for days $-$11, $-$1, +5, +14 and +40, produced with the highly parametrized SYNOW code, are discussed in the same section outlining the spectral peculiarities of this object. Velocities inferred from several spectral features are analyzed in section \\ref{sec_velocities}. A discussion about possible implications for supernova models is presented in section \\ref{sec_summary} and our conclusions are given in section \\ref{sec_conclusion}. ",
        "conclusions": "\\label{sec_conclusion} We have presented new high signal to noise optical spectroscopic data of SN~1999aa with good temporal coverage between $-$11 and +58 days with respect to B-band maximum light.  The overall evolution of the spectral features suggests an object with characteristics common to both SN~1991T-like and Branch-normal supernovae.  By means of SYNOW synthetic spectra we have attempted to identify the absorption features in the spectra at days $-$11, $-$1, +5, +14 and +40 with respect to B-band maximum light. Highlights of this modeling include the presence of \\begin{itemize} \\item C~{\\sc iii} at the photospheric velocity, \\item Possible C~{\\sc ii} or H at high velocity \\item Ca~{\\sc ii}~IR triplet at high velocity, \\item doubly ionized Si, Ni and Fe, \\item confined IMEs and Fe~{\\sc ii}, \\item probable iron peak core at 10000 km~s$^{-1}$. \\end{itemize} A schematic view of the resulting composition (in the velocity regime) is presented in Fig. \\ref{onion} .   The line identification for the earliest spectrum shows the presence of doubly ionized Si, Fe and Ni suggesting that the temperature in the outermost layer of SN~1999aa is higher than in normal SNe. After day 5, doubly ionized elements no longer form visible absorption features, highlighting the cooling of the supernova atmosphere.  The presence of IME above 14000 km~s$^{-1}$ excludes the possibility that SN~1999aa was the result of a sub-Chandrasekhar mass explosion.  The broad absorption feature around 6150~\\AA\\ in the spectrum at day $-$11 can be matched comparably well by a weak component of Si~{\\sc ii}~$\\lambda$6355 plus HV C~{\\sc ii}~$\\lambda$6580 or H$\\alpha$. The contribution of either C~{\\sc ii} or H$\\alpha$ helps in reproducing the velocity time evolution of the absorption feature at 6150~\\AA\\ that otherwise (i.e. considering only Si~{\\sc ii}~$\\lambda$6355) would remain unexplained. Definitive identification is not possible due to the lack of other absorption features produced either by C~{\\sc ii} or H at the input optical depths. However, both the possibilities impose constraints either on the nature of the progenitor system or the hydrodynamics of the explosion. A hydrogen component would imply that SN~1999aa is the result of a single degenerate WD explosion; a complete 3D calculation would be necessary to model it. Alternatively, a C~{\\sc ii} component can be reproduced by a pure 1D deflagration model, but would not be consistent with delayed-detonation models.  C~{\\sc iii}~$\\lambda$4648.8 at photosphere velocity was able to reproduce the absorption around 4500~\\AA, and no good alternatives were found for that feature. If this identification is confirmed by detailed radiative transfer models, the use of three dimensional simulations of the explosion would be necessary to describe the presence of unburned material down to the inner layers of the WD atmosphere.  At day $-$11 and at day $-$1, the profile of the absorption at 3800~\\AA\\ requires a HV component of Ca to be reproduced in addition to the photosphere component. The HV component at day $-$1 also reproduces the observed minima on the blue side of the usual Ca~IR triplet. For an accurate description of this feature, multi-dimensional simulations would be required.  The spectral evolution of SN~1999aa in the photospheric phase shows that Fe~{\\sc ii} is confined below 15000 km~s$^{-1}$. IMEs populate a narrow velocity window above 10000 km~s$^{-1}$.  Similar evidence is found in other well known supernovae (SN~1991T, SN~1997br and SN~2000cx) by studying the time evolution of the expansion velocity as computed from fits to the minima of Si~{\\sc ii}~$\\lambda$6355. The comparison with other SN~1991T-like objects suggests that the transition between IME to iron-peak dominant composition can occur at slightly different phases. Tuning of the deflagration to detonation transition density might reproduce this sequence.  The analysis of the spectrum of day +40 indicates that the iron-peak core limit should be set to 10000 km~s$^{-1}$, similar to that of SN~1991T \\markcite{1999MNRAS.304...67F}({Fisher} {et~al.} 1999).  The origin of the differences between normal supernovae and SN~1991T-like or SN~1999aa-like objects depends on several factors.  Our analysis of the optical spectra of SN~1999aa reveals that, among the present explosion models, none is able to reproduce each one of our findings. Higher temperature could account for some of the peculiarities \\markcite{1995ApJ...455L.147N}({Nugent} {et~al.} 1995), but the evidence for high velocity components and unburned material at all velocities probably requires the development of full NLTE 3D explosion models. SN~1999aa was spectroscopically less extreme than other genuine SN~1991T-like SNe, suggesting that perhaps a single model could eventually explain both normals and SN~1991T-like SNe as a continuous sequence."
    },
    "0404/astro-ph0404256_arXiv.txt": {
        "abstract": "We have modeled the time-variable profiles of the H$\\alpha$ emission line from the non-axisymmetric disk and debris tail created in the tidal disruption of a solar-type star by a $10^{6}\\Msun$ black hole. Two tidal disruption event simulations were carried out using a three dimensional relativistic smooth-particle hydrodynamic code, to describe the early evolution of the debris during the first fifty to ninety days. We have calculated the physical conditions and radiative processes in the debris using the photoionization code CLOUDY. We model the emission line profiles in the period immediately after the accretion rate onto the black hole became significant. We find that the line profiles at these very early stages of the evolution of the post-disruption debris do not resemble the double peaked profiles expected from a rotating disk since the debris has not yet settled into such a stable structure. As a result of the uneven distribution of the debris and the existence of a ``tidal tail'' (the stream of returning debris), the line profiles depend sensitively on the orientation of the tail relative to the line of sight. Moreover, the predicted line profiles vary on fairly short time scales (of order hours to days). Given the accretion rate onto the black hole we also model the H$\\alpha$ light curve from the debris and the evolution of the H$\\alpha$ line profiles in time. ",
        "introduction": "\\subsection{Tidal Disruption of a Star by a Black Hole and Related Issues}  A star in an orbit around a massive black hole can get tidally disrupted during its close passage by the black hole. After several orbital periods the debris from the disrupted star settles into an accretion disk and gradually falls into the black hole \\citep*{Rees,CLG,SC,LU}. As material gets swallowed by the black hole intense UV or soft-X ray radiation is expected to emerge from the innermost rings of the accretion disk \\citep{FR,LS,Frank,Phinney,SW,MT,SU}. For black hole masses $M_{bh} < 10^{7}\\Msun$, tidal disruption theory predicts flares with luminosities of the order of the Eddington luminosity with durations of the order of months, and with spectra that peak in the UV/X-ray domain band \\citep{Rees,EK,Ulmer,KPL,Gezari02}. High-energy flares from the central source illuminate the debris, the photons get absorbed, and some are re-emitted in the optical part of the spectrum (i.e. the light is ``reprocessed''). One of the spectral lines in which this phenomenon can be observed is the Balmer series H$\\alpha$ line ($\\lambda_{rest}=6563$~\\AA).  The disruption of a star begins when the star approaches the tidal radius, $r_{t}\\simeq r_{\\star}(M_{bh}/M_{\\star})^{1/3}$, the point where the surface gravity of the star equals the tidal acceleration from the black hole across the diameter of the star ($r_{\\star}$ and $M_{\\star}$ are the radius and mass of the star and $M_{bh}$ is the mass of the black hole). A $10^{6}\\Msun$ black hole is often used as a prototypical example in tidal disruption calculations. This choice is motivated by the criterion for a solar-type star to be disrupted before it crosses the black hole horizon (i.e. the Schwarzchild radius, $r_{s}$) in order for emission to be observable. For supermassive black holes with $M_{bh} > 10^{8}\\Msun$, $r_{s}>r_{t}$ and the star falls into the black hole before it gets disrupted.  The tidal disruption process has been the subject of many simulations \\citep*{CL82,CL83,BG,EK,KNP,Laguna,Frolov,MLB,Deiner,ALP,IN,ICN}. It has been shown that tidal processes in the vicinity of a massive black hole could lead to tidal capture, tidal heating and tidal spin-up of a star \\citep*{NPP,AK,AH,AM}, and in some cases ultimately to the explosion of the star. Such explosions, as well as accretion of post-disruption debris, should manifest themselves as luminous flares from the centers of galaxies \\citep{CL82,Rees}.  Stars close to a black hole may experience mixing or may eject some of their mass \\citep{Alexander}. As a consequence, stellar populations in nuclear clusters are expected to be somewhat unusual in comparison with populations whose evolution was not affected by a massive black hole \\citep{AL,DiStefano}. This has important implications for observations of stellar cluster in the center of our Galaxy \\citep*{Ghez,Schoedl,EOG,Figer,Gezari} where high concentration of otherwise rare blue He supergiants has been observed \\citep{Krabbe,Najarro}.  The tidal disruption and accretion of stars can fuel black holes in the centers of galaxies \\citep{Hills,DS,DDCa,DDCb,MCD,FB,Yu} and its contribution to nuclear activity in galaxies and the growth of the black hole mass depends on the rate of disruption events in a galaxy. The predicted tidal disruption rate in a typical inactive galaxy is $10^{-4}$--$10^{-5}~{\\rm yr}^{-1}$ \\citep{MT,Alexander}. This value is consistent with the rate of UV/X-ray outbursts observed with {\\it ROSAT} from inactive nuclei selected as tidal disruption candidates \\citep{Donley}. The rate of tidal disruptions in active and more luminous nuclei is estimated to be lower, with the lowest value of $10^{-9}\\,{\\rm yr}^{-1}$, for galaxies with the most massive black holes. This may occur partly because massive central black holes ($M_{bh} >10^{8}\\Msun$) swallow stars promptly, without disruption, and partly because stars are less centrally concentrated in these galaxies \\citep{MT}. The observed UV/X-ray flaring rate in these galaxies is about $9\\times10^{-4}\\,{\\rm galaxy}^{-1}\\,{\\rm yr}^{-1}$ \\citep{Donley} suggesting that in such nuclei, outbursts may be due to another mechanism, such as accretion-disk instabilities \\citep*{SCK,BKS}.  The tidal encounter of a compact star with a black hole can also result in emission of gravitational waves, which may be observable with upcoming instruments. More specifically, compact stars (helium stars, white dwarfs, neutron stars, and stellar-mass black holes) which can withstand large tidal forces without being disrupted, may get captured in relativistic orbits around a supermassive black hole. Due to the in-spiral and decay of the orbit those objects are expected to emit the peak of their gravitational wave power in the LISA frequency band \\citep{HB,SR,Freitag01}. It has been recently suggested by \\citet{Freitag03} that very-low mass main sequence stars (MSSs; $M\\ll 1~{\\rm M}_{\\odot}$) may contribute to events detected by LISA, which was not previously expected for capture of these objects by a supermassive black hole. Although MSSs produce a relatively weak gravitational signal during the in-spiral, compared to compact objects, their detection in the Galactic center is more likely because such stars have a predicted close-encounter rate that is an order of magnitude higher than that of white dwarfs (WDs), and about two or more orders of magnitude higher than that of neutron stars (NSs) and stellar-mass black holes (BHs). These compact MSSs are expected to produce a strong enough signal to allow for 0.5--2 detections from our Galactic center, with a signal-to-noise ratio of 10 or higher, for a LISA mission duration of one year \\citep{Freitag03}.  Moreover, MSSs are expected to be disrupted relatively early during the in-spiral, giving rise to possibly detectable electromagnetic flares. The sudden appearance of an electromagnetic counterpart to a transient gravitational wave source, expected in the case of MSSs and helium stars, could allow identification of the progenitor. In the case of a tidal disruption event in another galaxy, the coincidence of an electromagnetic flare and a gravitational wave signal would provide an indication that the event occurred at the very nucleus of the galaxy, and possibly allow the measurement of the redshift. More compact objects which can spiral in without being disrupted (such as WDs, NSs, and BHs) are expected to produce stronger gravitational wave signatures and weaker or no electromagnetic flares, with the exception of white dwarfs in which tidal interaction may trigger thermonuclear explosion \\citep*{GBW}.  \\subsection{Observational Motivation: Transient Emission Lines in Inactive Galaxies and LINERs}  In view of the above theoretical considerations, it is necessary to make predictions of the likely observational signatures of a tidal disruption event. The prompt UV/soft-X-ray flash that is expected to accompany the disruption does provide strong evidence for such an event and has in fact been detected in a number of galaxies with {\\it ROSAT} and {\\it HST} \\citep*{BPF,Grupea,Grupeb,Donley,BKD,KG,Greiner,GTL,Renzini,LNM}. However, the duration of this flash is short enough that it can easily be missed. Aftereffects with a longer duration, such as line emission from the debris, have a better chance of being detected. If the appearance of emission lines just after an X-ray flare could be detected from the same object it could be used to identify the tidal disruption in the early phase and would provide strong support for the overall picture, but such cases are rare \\citep*[see][]{Cappellari,Gezari}.  A set of tantalizing observations in the past decade show that several LINERs \\citep[low-ionization nuclear emission regions;][]{Heckman} have transient Balmer emission lines, which are often double-peaked; examples include NGC~1097 \\citep*{SBBW}, M81 \\citep{Bower}, NGC~4450 \\citep{Ho}, NGC~4203 \\citep{Shields} and NGC~3065 \\citep{EH2001}. Such line profiles are characteristic of rotating disks and resemble the persistent double-peaked Balmer lines found in about 10--20\\% of broad-line radio galaxies and about 3\\% of all active galaxies \\citep[e.g.,][]{EH1994,EH2003,strateva03}. Their abrupt appearance in LINERs led to suggestions that this transient event was related to the tidal disruption of a star by a supermassive, nuclear black hole \\citep{ELHS,SB95} or a change in the structure of the accretion disk associated with a change in accretion rate \\citep{SB97}.  To investigate the possibility of line emission from the post-disruption debris and to evaluate the suggestion that the transient double-peaked lines of LINERs are related to tidal disruption events, we have undertaken a calculation of the strength and profile of the H$\\alpha$ line emitted from the debris. In \\S2, We describe two SPH (Smoothed Particle Hydrodynamics) simulations of tidal disruption on which we base our further calculation of the line properties.  Our line profile calculation follows the method used for line profiles emitted by relativistic, Keplerian disks and is described in \\S3. In \\S4 we present the resulting line profiles and in \\S5 we discuss the physical conditions in the debris and the approximations used. In \\S6 we summarize our conclusions and consider future prospects. ",
        "conclusions": "We modeled the emission-line luminosity and profile from the debris released by the tidal disruption of a star by a black hole in the early phase of evolution. Our model predicts prompt optical evolution of post-disruption debris and profile shapes different from circular and elliptical disk model profiles. Since line profiles observed so far in LINERS look more disk-like and evolve slowly, the observations are likely to have caught the event at late times ($\\ge$ 6 months after the initial disruption), after the debris has settled into a quasi-stable configuration.  The line profiles can take a variety of shapes for different orientations of the debris tail relative to the observer. Due to the very diverse morphology of the debris, it is almost impossible to uniquely match the multi-peaked profile with the exact emission geometry. Nevertheless, the profile widths and shifts are strongly indicative of the velocity distribution and the location of matter emitting the bulk of the H$\\alpha$ light. Profile shapes do not depend sensitively on the shape of the light curve of the X-rays illuminating the debris. They strongly depend on the distance of the emitting material from the central ionizing source, which is a consequence of the finite propagation time of the ionization front and the redistribution of the debris in phase space. It may be possible to distinguish between the two effects observationally, based on their different characteristic time scales. The onset of the optically thick spheroidal halo should cause the disappearance of the broad H$\\alpha$ emission line on the time scale of months, and give rise to the emission of narrower, strong, blueshifted or redshifted emission line, arising from the portion of the tidal tail unobscured by the halo.  If X-ray flares and the predicted variable profiles could be observed from the same object they could be used to identify the tidal disruption event in its early phase. The X-ray flares can be promptly detected by all-sky synoptic X-ray surveys and high energy burst alert missions such as {\\it Swift}. The evolution of the tidal event may then be followed with optical telescopes from the ground on longer time scales and give an insight in the next stage of development of the debris. Thus, simulations of the tidal disruption process on longer time scales (of order several months to a few years) are sorely needed. Calculations of the long-term evolution of a tidal disruption event can predict the type of structure that the debris finally settles into and whether its emission-line signature resembles the transient double-peaked lines observed in LINERs. This study would provide an important insight into the evolution of LINERs.  Finally, the observed rate of tidally disrupted solar type stars can constrain the rate of captured compact objects (which are important gravitational wave sources), and the capture rate of main sequence stars in our Galaxy, which are expected to emit the peak of the gravitational radiation in the LISA frequency band and can be detected in the local universe."
    },
    "0404/astro-ph0404060_arXiv.txt": {
        "abstract": "The crisis of the standard cooling flow model brought about by {\\it Chandra} and {\\it XMM-Newton} observations of galaxy clusters, has led to the development of several models which explore different heating processes in order to assess if they can quench the cooling flow.  Among the most appealing mechanisms are thermal conduction and heating through buoyant gas deposited in the ICM by AGNs. We combine Virgo/M87 observations of three satellites ({\\it Chandra}, {\\it XMM-Newton} and {\\it Beppo-SAX}) to inspect the dynamics of the ICM in the center of the cluster.  Using the spectral deprojection technique, we derive the physical quantities describing the ICM and determine the extra-heating needed to balance the cooling flow assuming that thermal conduction operates at a fixed fraction of the Spitzer value. We assume that the extra-heating is due to buoyant gas and we fit the data using the model developed by Ruszkowski and Begelman (2002). We derive a scale radius for the model of $\\sim 5$ kpc, which is comparable with the M87 AGN jet extension, and a required luminosity of the AGN of a $few \\times 10^{42} {\\rm erg \\, s^{-1}}$, which is comparable to the observed AGN luminosity.  We discuss a scenario where the buoyant bubbles are filled of relativistic particles and magnetic field responsible for the radio emission in M87. The AGN is supposed to be intermittent and to inject populations of buoyant bubbles through a succession of outbursts.  We also study the X--ray cool component detected in the radio lobes and suggest that it is structured in blobs which are tied to the radio buoyant bubbles. ",
        "introduction": "The hot diffuse X-ray emitting gas (intracluster medium, ICM for short) provides a powerful tool to inspect the internal dynamics of galaxy clusters.  For the typical density and temperature of the intracluster gas, the main emission mechanism is the bremsstrahlung and, for a large amount of clusters, the radiative cooling time in the central regions is significantly shorter than the Hubble time.  As a consequence, if no additional heating mechanism is present, the gas cools and is expected to flow inwards, forming a {\\it cooling flow}. The standard model of cooling flows \\citep[see][for a review]{fa94} predicted the gas to be a multiphase medium in which there is a broad range of temperatures and densities present at all radii.  Mass deposition rates were estimated to be as large as hundreds of solar masses per year \\citep{Allen01}.  This model was strengthened by the general thought that in presence of magnetic fields the thermal conduction must be highly suppressed \\citep{BC81,fa94,CC98, Mal01}, which is a necessary condition for the multiphase cooling to operate. In fact, no heating exchange between the different phases must occur in order that they may coexist.  There is some observational evidence that modest magnetic fields are present throughout the intracluster medium.  The current measurements of intracluster magnetic fields are based on Faraday rotation measure (RM) in radio sources seen through clusters \\citep[e.g.][]{kim91,CKB,Fer99,Tay01}; direct evidence also comes from measurements of extended regions of radio synchrotron emission in clusters \\citep[see e.g.][]{GF00,fusco,OMV,F99}. Both the excess RM values and the radio halo data suggest modest magnetic fields, at a few microgauss levels, throughout the cluster.  Recent \\XMM\\, and {\\it Chandra} observations have shown that in the central regions, the temperature drops to about one third of its overall mean value and there is no evidence of temperatures smaller than $\\sim 1-2$ keV \\citep{Pet01,Kaa01,Tam01,Allen01}, suggesting that the gas does not cool below these cutoff temperatures.  Moreover, the new estimated mass deposition rates are significantly smaller than those evaluated by using previous X-ray satellites data \\citep{McN01,Pet01}.  Lastly, the new data show that clusters spectra are better represented by a single (or double) temperature model rather than the standard multiphase (multi-temperature) cooling flow model \\citep{MP01,Bohr01,Fab01,matsu}.  These new results clearly show that the standard cooling flow model is not a satisfying description of the internal dynamics of the ICM. Some source of heat which stops the cooling flow and balances radiative losses must be sought.  The nature of this source and the origin of the heat mechanism is still unclear.  One possible candidate is thermal conduction.  Recent works by \\citet{NM01} and \\citet{gruz02} show that in the presence of turbulent magnetic fields, the conductivity can be as large as a fraction of the Spitzer value and thus can play a significant role in balancing cooling flows. As a consequence, thermal conduction has been recently re-introduced as a possible heat source to balance the energy losses \\citep[see e.g.][]{Voigt,VF04,Fab02, Mal01,ZN02}.  However, as we will discuss more in detail in \\S \\ref{sub:cond}, thermal conduction fails in supplying the needed heat in the central regions \\citep[see also][]{Voigt,VF04, ZN02}.  Heating from a central active galactic nucleus is another possibility. The idea is supported by the fact that most of ``cooling clusters'' host a central active galactic nucleus with strong radio activity \\citep{Burns,BM02}. Several models in the literature explore AGN heating processes to assess if they can balance the radiative losses. One of the most appealing mechanisms involves buoyant gas bubbles, inflated by the AGN, that subsequently rise through the cluster ICM heating it up \\citep{ch02,Bohr02,BKa,BKb,Brug02}.  However, all the models predicting that radiative cooling is balanced only by energy input from the central AGN fail \\citep{McN02,ZN02,BM02}.  In particular, \\citet{BM02} analyzed several heating mechanisms induced by the central AGN and concluded that no simple mechanism is able to quench the cooling flow.  Moreover, the required mechanism needs a finely tuned heating source. Indeed, the heat source must provide sufficient energy to stop the cooling flow, but not enough to trigger strong convection or the metallicity gradients observed in all cooling flow clusters \\citep{DM01} would be destroyed. As a consequence, an AGN can be an efficient mechanism in the very center of the cluster but it is unlikely to be strong enough to provide energy to the outer parts of the ``cooling region''. So, it may be viewed as complementary to thermal conduction which fails in quenching the cooling flow in the innermost regions.  Recently, \\citet{RB02} and \\citet{ZN02} concluded that both thermal conduction and heating from a central AGN can play an important role in balancing the cooling.  In particular, \\citeauthor{RB02} (RB02 hereafter) developed a model where both thermal conduction and heating from a central AGN co-operate in balancing the radiative losses.  One of the main advantages of this model is that it reaches a stable final equilibrium state and it is able to reproduce the main observed quantities, such as the temperature profile (with a minimum temperature $T \\sim 1$ keV).  In this paper, we use M87/Virgo observations of three satellites (namely \\XMM, {\\it Chandra} and {\\it Beppo-SAX}) to test various heating models on this cluster. To this end, we apply to the M87 data the deprojection technique to recover some physical quantities of the ICM such as the gravitational mass, the entropy and the heating required to balance the cooling flow.  The paper is organized as follows: in \\S \\ref{sec:sp2d} we report details about the analysis of the three ({\\it Chandra}, \\XMM\\, and {\\it Beppo-SAX}) M87 datasets; in \\S \\ref{sec:depro}, we revise briefly the spectral deprojection technique that we adopt for our analysis; in \\S \\ref{sec:m87}, we deproject the M87 data, we test that the spherical symmetry hypothesis holds and we derive the gravitational mass for M87; in \\S \\ref{sec:heat}, we determine the amount of extra-heating needed to balance the cooling flow when thermal conductivity is assumed to operate at a fraction of the Spitzer value.  Lastly, in \\S \\ref{sec:concl} we summarize our results. ",
        "conclusions": "Starting from the results inferred in the previous section, we can try to draw a more general picture, using also informations coming from radio observations of M87.  We can suppose that the buoyant bubbles are radio bubbles filled with magnetic field and relativistic particles \\citep{GN73,ch00,ch01,BK01,ch02} responsible for the synchrotron emission in M87.  As already outlined by OEK, the radio structures are highly filamented.  This suggests that the dimensions of the bubbles are small.  \\citet{EH02} discussed the dynamics of the rise of buoyant light bubbles within the cluster atmosphere \\citep[see also][]{ch01,Kai03}. The buoyant bubble rapidly reaches a terminal velocity $v_b$. In the limit of small bubbles, $v_b$ can be estimated by balancing the buoyancy force with the ram pressure (drag force) of the cluster gas.  The buoyancy force is  \\begin{equation} \\label{eq:fbuoy} F_b = V g (\\rho  - \\rho_b)  \\, = V g \\rho \\Delta \\, , \\end{equation} \\noindent where $V = 4/3 \\pi r_b^3$ is the volume of the bubble, $r_b$ is the bubble radius, $g=G M(<r)/r^2$ is the local gravity acceleration at the radius $r$ ($M$ is the gravitational mass within the radius $r$); $\\rho$ and $\\rho_b$ are respectively the density of the ICM and of the bubble. $\\Delta =(\\rho-\\rho_b)/\\rho$ is the density contrast.  The drag force for subsonic motion can be approximated by  \\begin{equation} \\label{eq:drag} F_{drag} = C \\pi r_b^2 \\rho v_b^2 \\end{equation} \\noindent with the drag coefficient $C\\simeq 0.5$.  By equating eqs. \\eqn{eq:fbuoy} and \\eqn{eq:drag} the velocity of the bubble can be determined.  \\citet{EH02} derived $v_b$ under the assumption that the density of the ICM is well described by an isothermal $\\beta$-model and the density contrast $\\Delta \\sim 1$. In this case $v_b$ is a fraction ($\\propto \\sqrt{r_b/r_c}$; $r_c$ is the core radius of the cluster) of the sound velocity.  Correspondingly, for small bubbles, $v_b$ is subsonic which is consistent with the fact that no shocks are detected in M87.  We derived the rise velocity of the bubbles in M87 at a radius $r \\sim 10$ kpc, using the results from the deprojection for $M(<r)$ and $\\rho$.  The density contrast $\\Delta$ can be inferred considering that the bubbles filled with relativistic plasma are in pressure equilibrium with the ICM.  Indeed, the pressure equilibrium between radio bubbles and thermal plasma implies that $n_e kT \\sim \\beta n_{er} \\langle E_e\\rangle$ (where $\\beta$ is a factor of the order of the unity which accounts for the contribution to the pressure of the magnetic field which is almost in equipartition with the particle energy).  $\\langle E_e\\rangle$ is the mean energy of the relativistic particles and $n_{er}$ is the numerical density of the relativistic particles. The relativistic particles mean energy is much larger than that of the thermal electrons in the ICM, leading to a density contrast $\\Delta \\sim 1$.  The values of $v_b$ for the bubble radius $r_b$ varying in $[0.01-1.]$ kpc range are plotted in Fig. \\ref{fig:velb} and the motion is indeed subsonic.  It is worth considering that, in a recent analysis of these data \\citep[M02 hereafter]{ms02}, we have found that, cospatially with the radio lobes, there exists a thermal component with $T \\sim {1\\over2} T_{ICM}$. This component is likely structured in blobs with typical scales smaller than $\\sim 100$ pc.  The filling factor of these blobs has been estimated to be of the order of the percent.  Some informations about the filling factor of the radio bubbles can be obtained using recent results from radio observation of M87.  Using the standard minimum pressure analysis \\citep[e.g.][]{Pacho70,bor,OO}, OEK evaluate the magnetic field in the lobes and in the filaments visible in the radio map and evaluate the minimum pressure of the magnetic field and relativistic particles.  They assume that the proton-to-electron energy is $k=1$ and that $\\zeta\\phi=1$ where $\\phi$ is the filling factor of the relativistic particles and of the magnetic field and $\\zeta$ is the ratio of the true magnetic pressure (comprehensive of the tension) to the magnetic pressure $B^2/8\\pi$. For a tangled magnetic field $\\zeta=1/3$.  The estimations for the minimum pressure derived by OEK can be compared with the thermal pressure that we can infer from our deprojected $T$ and $n_e$ profiles.  In agreement with OEK, we find that the pressure of the relativistic plasma is smaller than the thermal pressure by about an order of magnitude.  By keeping the condition $k=1$ and assuming $\\zeta=1$, we can derive the filling factor which reconciles the minimum pressure $P_{min}$ with the thermal pressure of the plasma, considering that the minimum pressure scales with $\\phi$ according to:  \\begin{equation} {{P_{min}(\\phi) } \\over {P_{min}(\\phi=1)}} =\\left(\\phi\\right)^{-4/7} \\, . \\end{equation}  The discrepancy of a factor $\\sim 10$ between the two pressures can be eliminated by assuming a filling factor $\\sim 1 \\%$. By using the estimations of $P_{min}$ from OEK in different plasma of the radio lobes and filaments, we infer filling factors of few $\\%$. Assuming $\\zeta=1/3$ (tangled magnetic field) reduces $\\phi$ by a factor $\\sim 2$.  The filling factors of the radio bubbles and of the cold thermal blobs are of the same order.  The survival of the cold thermal blobs in the hotter ICM requires that thermal conduction be substantially suppressed; this may happen if these blobs are tied to the radio bubbles and magnetic fields shield them from collisions with ICM particles. Since the filling factors of blobs and bubbles are similar and the density of the blobs is about twice that of the ICM (M02), the mean density of each bubble+blob is about that of the surrounding ICM.  Thus, assuming that the blobs are tied to the bubbles and that they occupy similar volumes, their density contrast with the surrounding ICM will be small. In Fig. \\ref{fig:velb} the dotted line plots the bubble rise velocity $v_b$ for a $\\Delta=0.1$.  Starting from the rise velocity, it is straightforward to derive the rise timescale of the bubbles.  For a density $\\Delta \\sim 1$ and a bubble radius of $\\sim 100$ pc, the risetime is $t_{rise} \\sim 10^8$ yr at $r \\sim 10$ kpc and it becomes even larger for smaller bubbles and smaller density contrasts, holding $t_{rise} \\propto {(r_b \\Delta)}^{-1/2}$.  As outlined previously, this value is significantly larger than the duty cycle timescale $t_i$ of intermittency of the AGN, so that the ratio $t_{rise}/t_i \\simgt 10^3-10^4$ and the mechanism approaches the steady state.  It is worth noting that the above picture refers to radio bubbles and cool X-ray blobs located in the lobes. However, the heating process related to the rise of the bubbles must be isotropic throughout the cluster in order to balance the cooling flow.  Under the assumption that the radio galaxy is intermittent on a short timescale, the mechanism is expected to heat isotropically the ICM since it is likely that no direction is preferred for the AGN ejection; we also recall that when the AGN is turned off the bubbles continue their rise within the ICM heating it up; within this picture, by averaging on a cooling time, the radio bubble populations are likely to be isotropically distributed in the cluster.  Nevertheless, the cool X--ray blobs are detected only in the lobes regions.  So the picture emerging here is that of an AGN which injects radio bubbles, in all the directions, through a succession of outbursts.  In the lobes, the outburst is occurring at the present time and, here, also the cool X-ray blobs are present. They are tied to the radio bubbles and the magnetic fields shield them from collisions with ICM particles. The thermal conduction is inhibited allowing the cool blobs to survive in a hotter ambient medium.  In the halo, where the radio bubbles have been injected during a past activity of the AGN, only ``old'' populations of bubbles which are buoyantly rising are present.  It is likely that the cool blobs which were tied to the radio bubbles when the AGN was active there, have thermalized, the magnetic fields slow down but do not stop entirely thermalization so that only a single phase gas is detected in the halo regions."
    },
    "0404/astro-ph0404310_arXiv.txt": {
        "abstract": "{ We report results from a new analysis of a deep 95 {\\rm $\\mu$m} imaging survey with ISOPHOT on board the Infrared Space Observatory, over a $\\sim$1 deg$^2$ area within the Lockman Hole, which extends the statistics of our previous study (Rodighiero et al. 2003). Within the survey area we detect sixty-four sources with S/N$>3$ (roughly corresponding to a flux limit of 16 {\\rm mJy}). Extensive simulations indicate that the sample is almost complete at $S_{95 \\mu m} \\geq 100$ {\\rm mJy}, while the incompleteness can be quantified down to $\\sim$30 {\\rm mJy}. The 95 {\\rm $\\mu$m} galaxy counts reveal a steep slope at $S_{95 \\mu m} \\le 100$ {\\rm mJy} ($\\alpha\\simeq 1.6$), in excess of that expected for a non-evolving source population. In agreement with counts data from ISO at 15 and 175 {\\rm $\\mu$m}, this favours a model where the IR populations evolve both in number and luminosity densities. We finally comment on some differences found with other ISO results in this area. ",
        "introduction": "The infrared sky observed from space is providing tight constraints on the evolution of cosmological source populations. The IRAS extragalactic number counts (Oliver et al. 1992, Bertin et al. 1997) showed some marginally significant excess of galaxies compared to non-evolutionary predictions. More recently, ISOPHOT (Lemke et al. 1996) found some evidence of evolution for the IR population  at 170{\\rm $\\mu$m} (e.g. Puget et al. 1999, Dole et al. 2001), but source confusion at these wavelengths limits the observations to moderately faint fluxes ($S_{175\\mu m}>135$ {\\rm mJy}). At shorter wavelengths, ISOCAM (Cesarsky et al. 1996) detected a number of faint sources consistent with strong evolution in the mid-IR  (e.g. Elbaz et al. 1999), a factor $\\sim$10 in excess with respect to non-evolution predictions.  The ISO deep counts, together with those in the submillimeter (e.g. Scott et al. 2002) are used as constraints on models of galaxy formation and evolution (e.g. Lacey et al. 1994, Franceschini et al. 2001,  Lagache et al. 2003). These studies are also needed to estimate the confusion noise for future IR and submillimeter telescopes (Xu et al. 2001, Lagache et al. 2003).  In this paper we present a new analysis of ISOPHOT 95 $\\mu$m data in the direction of the Lockman Hole, extending our previous study (Rodighiero et al. 2003, hereafter Paper I) and significantly improving the source statistics.  The Spitzer Space Observatory (Fazio et al., 1999) will soon observe the same region of the Lockman Hole in complementary wavebands at 24, 70 and 160 {\\rm $\\mu$m} with MIPS and in the near-IR with IRAC. This will provide additional extensive information on the Spectral Energy Distributions (SEDs) of ISO sources.  The paper is organized as follows: in Section 2 we introduce the dataset and in Section 3 we summarize features of the adopted reduction procedure. The infrared maps and the new catalogue are presented in Section 4. The source counts and a discussion of the effects of confusion and cosmic variance are reported in Section 5. ",
        "conclusions": ""
    },
    "0404/astro-ph0404126_arXiv.txt": {
        "abstract": "On January 6$^{th}$ 2004, the IBAS burst alert system triggered the 8$^{th}$ gamma--ray burst (GRB) to be detected by the \\textit{INTEGRAL} satellite. The position was determined and publicly distributed within 12\\,s, enabling ESA's \\textit{XMM--Newton} to take advantage of a ToO observation just 5\\,hours later during which the X--ray afterglow was detected. Observations at optical wavelengths also revealed the existence of a fading optical source. The GRB is $\\sim$\\,52\\,s long with 2 distinct peaks separated by $\\sim$\\,24\\,s. At gamma--ray energies the burst was the weakest detected by \\textit{INTEGRAL} up to that time with a flux in the 20\\,-\\,200\\,keV band of 0.57\\,photons\\,cm$^{-2}$\\,s$^{-1}$. Nevertheless, it was possible to determine its position and extract spectra and fluxes. Here we present light curves and the results of imaging, spectral and temporal analyses of the prompt emission and the onset of the afterglow from \\textit{INTEGRAL} data. ",
        "introduction": "Gamma--ray bursts are an amazingly energetic phenomenon, capable of an isotropic output of order 10$^{52}$-10$^{54}$\\,erg in a few seconds. Although first detected in the late 1960s, significant progress has mostly been achieved in the last dozen years. That GRBs are extra--galactic in origin was suggested by the isotropic distribution of GRBs observed by BATSE on board the Compton Gamma--Ray Observatory \\citep{mee1992,fish1994}. The discovery by BeppoSAX of afterglows in the X--ray \\citep{costa:1997} and subsequent discoveries at optical \\citep{vanp:1997} and radio \\citep{frail:1997} wavelengths have led to redshift measurements \\citep{metz:1997} for $\\sim$\\,40 bursts ranging from $z = 0.168-4.5$. A theory of gamma--ray bursts must provide a mechanism capable of releasing enormous quantities of non--thermal energy by compact sources at cosmological distances.  Although not built as a GRB oriented mission, \\textit{INTEGRAL} has a burst alert system (IBAS) and the two main instruments on board have coded masks, a wide field of view (FoV), cover a wide energy range (15\\,keV\\,-\\,8\\,MeV) and offer high resolution capabilities in imaging (IBIS) and spectroscopy (SPI). IBAS carries out rapid localisations for GRBs incident on the IBIS detector with precision of a few arcminutes \\citep{vonk2003}. The public distribution of these co--ordinates enables multi--wavelength searches for afterglows at lower energies. \\textit{INTEGRAL} data of the prompt emission in combination with the early multi--wavelength studies offers the best currently available probe of the origin of these transient phenomena.  In \\S 2 observations and imaging analysis of GRB\\,040106 are presented. \\S 3 describes the spectral analysis methods utilised and the results obtained from analysis of SPI data. A brief account of the temporal analysis is presented in \\S 4.  \\begin{figure} \\centering \\includegraphics[width=0.9\\linewidth]{lightcurve1.eps} \\includegraphics[width=0.9\\linewidth]{lightcurve2.eps} \\caption{ISGRI light curves of GRB\\,040106 in 2 energy ranges (upper panel) 15\\,-\\,40\\,keV and (lower panel) 40\\,-\\,200\\,keV}\\label{ibis_lc} \\end{figure} ",
        "conclusions": "On December 19$^{th}$ 2003, SPI's detector \\#2 was confirmed dead after several attempts to revive it met with no success. As yet a new redistribution matrix has not been released to take account of this change in SPI's response. GRB\\,040106 was incident on detector \\#2 and the surrounding detectors, so an improvement in the fit is expected when the new matrix is issued.  The results of spectral analysis of GRB\\,040106 confirm that the second peak is harder than the first and that it is well fit by a single power--law model of photon index $\\alpha \\sim$ -1.3. It is unlikely that this is an X--ray rich GRB with a peak energy at or below the low end of the SPI detector sensitivity (i.e. $\\sim$\\,20\\,keV) since the spectral index would correspond to an unusually hard value for the high--energy index above the spectral turnover \\citep{pree2000}. It is more likely that the weakness of this GRB washes out evidence for a spectral break at more typical energies of a few hundred keV. There is no evidence in the IBIS/ISGRI light curve for soft extended or delayed emission such as that observed by, for example, SIGMA/GRANAT in GRB\\,920723 \\citep{buren1999} or by HETE--II in GRB\\,021211 \\citep{crew2003}. The temporal decay of the 2$^{\\rm nd}$ peak is consistent with a power--law of slope $\\beta = -0.3\\pm 0.7$ which may indicate the presence of a high--energy afterglow, due to external shocks, during the burst itself. However, the data are not sufficiently constrained to indicate fast or slow, or radiative or adiabatic, cooling in the synchrotron shock model \\citep{gib2002,piro2004}. Further analysis and comparison with XMM--Newton results are on--going \\citep{moran2004}."
    },
    "0404/astro-ph0404191.txt": {
        "abstract": "{Since the 1990's, protoplanetary disks and planetary disks have been intensively observed from the optical to the millimetre wavelength and many models have been developed to investigate their gas and dust properties and dynamics. These studies remain empirical and rely on poor statistics with only a few well known objects. However, the late phases of the stellar formation are among the most critical for the formation of planetary systems. Therefore, we believe it is timely to tentatively summarize {\\it the observed properties of circumstellar disks around young stars from the protoplanetary to the planetary phases}. Our main concern is to present the physical properties considered as observationally robust and to show their main physical differences associated to an evolutionary scheme. We also describe areas still poorly understood such as how protoplanetary disks disappear to lead to planetary disks and eventually planets.} ",
        "introduction": "% Before the 1980's, the existence of protoplanetary disks of gas and dust around stars similar to the young Sun (4.5 billion years ago) was inferred from the theory of stellar formation (e.g. Shakura \\& Suynaev 1973), the knowledge of our own planetary system and dedicated models of the Proto-Solar Nebula. The discovery of the first bipolar outflow in L1551 in 1980 drastically changed the view of the stellar formation. In the meantime, optical polarimetric observations by Elsasser \\& Staude (1978) revealed the existence of elongated and flattened circumstellar dust material around some Pre-Main-Sequence (PMS) stars such as the low-mass TTauri stars. The TTauri are understood to be analogs to the Sun when it was about $10^6$ years old. A few years later, observations from the InfraRed Astronomical Satellite (IRAS) found significant infrared (IR) excesses around many TTauri stars, showing the existence of cold circumstellar dust (Rucinski 1985). More surprisingly, IRAS also show the existence of weak IR excess around Main-Sequence (MS) stars such as Vega, $\\epsilon$ Eridani or $\\beta$ Pictoris (Aumann \\etal 1984). These exciting discoveries motivated several groups to model the Spectral Energy Distribution (SED) of TTauri stars (e.g. Adams \\etal 1987) and Vega-like stars (Harper \\etal 1984).  On one hand, for PMS stars, the emerging scenario was the confirmation of the existence circumstellar disks orbiting the TTauri stars, the gas and dust being residual from the molecular cloud which formed the central star (Shu \\etal 1987). Since such disks contain enough gas (H$_2$) to allow, in theory, formation of giant planets, they are often called ``protoplanetary'' disks. During this phase, the dust emission is optically thick in the Near-IR (NIR) and the central young star still accretes from its disk. Such disks are also naturally called accretion disks.  On the other hand, images of $\\beta$ Pictoris by Smith \\& Terrile (1984) demonstrated that Vega-type stars or old PMS stars can also be surrounded by optically thin dusty disks. These disks were called ``debris'' disks or later ``planetary'' disks because planetesimals should be present and indirect evidence of planets was found in some of them ($\\beta$~Pic).   In this chapter, we review the current observational knowledge of circumstellar disks from the domain of the UV to the millimeter (mm). We discuss in Sections \\ref{ttauri} and \\ref{herbigae} the properties of protoplanetary disks found around young low-mass (TTauri) and intermediate-mass (Herbig AeBe) stars. In section \\ref{transition}, we summarize the properties of {\\it transition disks} which have still some gas component but also have almost optically thin dust emission in the NIR, objects which are thought to be in the phase of dissipating their primary gas and dust. We present the properties of optically thin dust disks orbiting old PMS, Zero-Age-Main-Sequence (ZAMS) or Vega-type stars, such as the $\\beta$ Pic {\\it debris disk} in Section \\ref{planeto}. We conclude by reviewing future instruments and their interest for studying such objects. \\label{intro} %AD - 1 page % ",
        "conclusions": ""
    },
    "0404/astro-ph0404195_arXiv.txt": {
        "abstract": "We report the results of a cosmic shear survey using the 4.2m William Herschel Telescope on La Palma, to a depth of $R=25.8$ ($z\\approx0.8$), over 4 square degrees. The shear correlation functions are measured on scales from $1'$ to $15'$, and are used to constrain cosmological parameters. We ensure that our measurements are free from instrumental systematics \\com{by performing a series of tests, including a decomposition of the signal into $E$- and $B$-modes. We also reanalyse the data independently, using the shear measurement pipeline developed for the COMBO-17 survey. This confirms our results and also highlights various effects introduced by different implementations of the basic ``KSB'' shear measurement method.} We find that the normalisation of the matter power spectrum on 8 $h^{-1}$Mpc scales is $\\sigma_8=(1.02\\pm 0.15)(0.3/\\Omega_m)^{0.5}$, where the 68\\%CL error includes noise, sample variance, covariance between angular scales, systematic effects, redshift uncertainty and marginalisation over other parameters. We compare these results with other cosmic shear surveys and with recent constraints from the WMAP experiment. ",
        "introduction": "\\label{intro}  Weak gravitational lensing by large-scale structure, or ``cosmic shear'', has emerged as a powerful cosmological probe, as it is directly sensitive to foreground mass \\citep[for reviews, see][]{bs,bern99,melrev,refrev,witrev}. A measurement of cosmic shear is therefore closely tied to cosmological theories, which are principally concerned with the distribution of dark matter. In particular, the systematic biases of this technique are not limited by unknown physics such as biasing \\citep{oferbias,meg901,hoebias,smithbias,weinbias} or the mass-temperature relation for X-ray selected galaxy clusters \\citep{hutw,pier,vlnew}.  Cosmic shear surveys are rapidly growing in size and precision \\citep{bmer,browncs,hamanacs,hoecs,jarviscs,ref02,rrgmeas,vw02}. Cosmological parameter constraints from these surveys are now approaching the precision of other methods.  However, cosmic shear surveys can be subject to several systematic biases of their own. Imperfect telescope tracking, telescope flexure or optical misalignment within the camera, even at a level that is acceptable for most purposes, can artificially distort images and align the shapes of distant galaxies in a way that mimics cosmic shear.  The survey described in this paper represents a culmination of effort at the William Herschel Telescope. We have combined the experience of instrumentalists with detailed image simulations and careful data analysis to control the various sources of systematic error. Our first cosmic shear paper \\citep{bre} reported an initial detection of cosmic shear using a 0.5 square degree survey with the William Herschel Telescope (WHT). The second paper \\citep{bmer} compared the WHT shear signal with an independent measurement using the Keck~II telescope, and examined systematics from these two very different instruments. In this paper, we extend our WHT survey to cover 4 square degrees to constrain cosmological parameters, while paying great care in monitoring and correcting systematic effects. \\com{We also now test our entire pipeline for shear measurement and cosmological parameter estimation by comparing it to external code, developed independently for the COMBO-17 survey by \\citet{browncs}.}  This paper is organised as follows.  In \\S\\ref{observations} we describe our survey strategy and observational parameters. In \\S\\ref{results} we present our results and draw constraints upon cosmological parameters. In \\S\\ref{systematics} we test for the absence of any systematic errors. \\com{In \\S\\ref{mlb}, we present a second set of results obtined via an independent shear measurement pipeline}. We conclude in \\S\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  We have measured the weak lensing shear-shear correlation functions in four square degrees of deep $R$-band imaging data from the William Herschel Telescope. Our measurements constrain the amplitude of the mass power spectrum, $\\sigma_8(\\Omega_{m}/0.3)^{0.52}=1.02\\pm 0.15$, including all contributions to the total 68\\%CL error budget: statistical noise, sample variance,  an additional estimated error due to binning instabilities from non-Gaussian outliers, covariance between different angular scales,  systematic measurement and detection biases, source redshift uncertainty, and marginalisation with priors over other parameters. We have examined our data for contamination by systematic effects using a variety of tests including an $E$-$B$ decomposition. These demonstrate a well-understood and modest contribution to our uncertainties from systematic errors.  Using the pipeline developed for earlier WHT studies, we find our measurement of the normalization of the dark matter power spectrum lies at the relatively high end  of the distribution of published values. However, it is still consistent with those from  equivalently deep surveys by \\citet{ref02} and \\citep{jr_stis}. Our results are also consistent at the 1$\\sigma$ level  with CMB results from the Wilkinson Microwave Anisotropy Probe (WMAP) \\citep{wmap1yr}.  The wide distribution of $\\sigma_8$ constraints from recent cosmic shear surveys understandably casts some aspersion upon their precision. For example, it might be argued that the dispersion largely arises from unknown or poorly-understood systematic effects. For the first time we analyze our dataset with an independent pipeline -- that developed for the COMBO-17 data (\\citet{browncs}). This provides a valuable check on the extent to which dispersion in the published cosmological results, as well as our apparently high normalization, might arise from different techniques, both in constructing shear catalogues and in analysing the shear-shear correlation functions. Reassuringly, we find remarkable concordance between the two pipelines for the same dataset. However, the various differences we have explored, for example in the selection thresholds, seem to be insufficient to reconcile the spread in observed $\\sigma_8$ values, suggesting that significant differences remain at some level in the actual data themselves.  Uncertainties in the redshift distribution of source galaxies in deep data clearly contribute. It is difficult to determine the precise redshift distribution of galaxies  after excluding those smaller than a fixed apparent size. We have been conservative in this analysis and, as seen in equation~(\\ref{eq:con_full_errors}), source redshift uncertainty is already a major component of our total error budget. The resolution of such issues will require extensive spectroscopic follow-up and more complete image simulations. Such advances are essential if the potential of the next generation of cosmic shear surveys is to be fully realised.  Finally, we note that most recent cosmic shear results remain discrepant at the $3\\sigma$ level with measurements derived from the abundance of X-ray selected cluster samples based on an observational rather than theoretical mass-temperature relation \\citep{bor,selclus,rei,vlnew}. These suggest $\\sigma_8\\approx 0.75$. \\citet{adamgauss} concluded that even extreme non-Gaussianity in the mass distribution would be insufficient to explain this discrepancy, because the two techniques probe similar mass scales. Further studies are therefore needed in both the cluster method, to understand the difference between the observed mass-temperature relation and that found in numerical simulations; and in the weak lensing method, to construct more reliable and better calibrated shear measurement methods. Such consistency checks will represent a crucial verification of the standard $\\Lambda$CDM paradigm, so resolving this issue is of paramount importance."
    },
    "0404/astro-ph0404476_arXiv.txt": {
        "abstract": "Multipole expansion of spatial three-point statistics is introduced as a tool for investigating and displaying configuration dependence. The novel  parametrization renders the relation between bi-spectrum and three-point correlation function especially transparent as a set of two-dimensional Hankel transforms. It is expected on theoretical grounds, that three-point statistics can be described accurately with only a few multipoles. In particular, we show that in the weakly non-linear regime, the multipoles of the reduced bispectrum, $Q_l$, are significant only up to quadrupole. Moreover,  the non-linear bias in the weakly non-linear regime only affects the monopole order of these statistics. As a consequence, a simple, novel set of estimators can be constructed to constrain galaxy bias. In addition,  the quadrupole to dipole ratio is independent of the bias, thus it becomes a novel diagnostic of the underlying theoretical assumptions: weakly non-linear gravity and perturbative local bias. To illustrate the use of our approach, we present predictions based on both power law, and CDM models. We show that the presently favoured SDSS-WMAP concordance model displays strong ``baryon bumps'' in the $Q_l$'s. Finally, we sketch out three practical techniques estimate these novel quantities: they amount to new, and for the first time edge corrected, estimators for the bispectrum. ",
        "introduction": "Three-point statistics are on track to become main stream tools in astronomy. They have been used by several authors with considerable success to constrain the statistical bias between the distribution of galaxies and dark matter \\citep[e.g.][]{Fry1994,JingBoerner1998,FriemanGaztanaga1999, SzapudiEtal2000,ScoccimarroEtal2001,VerdeEtal2002}, as well as to constrain primordial non-Gaussianity of the Cosmic Microwave Background \\citep{KomatsuEtal2003}.  Nevertheless, the three-point correlation function and its Fourier-transform pair, the bispectrum, are complicated objects. In their most general form, they depend on three vectors, i.e. nine variables. Even after translational and  rotational  symmetries are taken into account, three-point functions depend on the size and shape of a triangle, i.e. still three parameters are needed. Exploration and visualization of the full configuration space becomes a surprizingly daunting task,  and most previous studies have been forced to restrict their investigation to a few hand-picked triangle shapes and sizes, which are not necessarily representative.  In the past essentially three distinct parametrizations have been proposed. \\cite{peebles1980} used the parameters $r,u,v$, where the three sides of the triangles are $r=r_1\\le r_2\\le r_3$, and $u=r_2/r_1$, and $v=(r_3-r_2)/r_1$. Alternatively, \\cite{SutoMatsubara1994,deeprange01} have used logarithmic bins for the three sides of a triangle. Then shapes can be uniquely described by a pair of integers $(b_1-b_3,b_2-b_3)$ formed from the bin numbers $b_1 \\le b_2 \\le b_3$. Finally, several authors parametrized triangles with two sides and the angle between then, $r_1,r_2,\\theta$ \\citep[e.g.][]{Fry1994}.  Systematic exploration of the full configuration space, and comparison of results are somewhat tedious using any of the above parametrizations. Moreover, the connection between  real and transform space representations (three-point function and bispectrum) is somewhat obscure, whichever description one uses. Hoping to alleviate  these shortcomings, we introduce the multipole expansion of three-point statistics, which arguably expresses the rotational symmetries in the most natural way. The next section presents the definition of the three-point multipoles and the relationship between the multipoles in real and Fourier space. Section 3 applies these results in the weakly non-linear regime, bias, and presents predictions for power law and CDM models. The final section contains summary and discussion of the results, including measurement techniques for the proposed quantities. ",
        "conclusions": ""
    },
    "0404/astro-ph0404530_arXiv.txt": {
        "abstract": "Recent findings indicate that the Monogem Ring and the associated pulsar PSR B0656+14 may be the `Single Source' responsible for the formation of the sharp knee in the cosmic ray energy spectrum at $\\sim$3PeV. The energy spectum of cosmic rays expected for the Monogem Ring supernova remnant (~SNR~) from our SNR acceleration model \\cite{EW9} has been published by us elsewhere \\cite{EW7}. In this paper we go on to estimate the contribution of the pulsar B0656+14 to the cosmic rays in the PeV region. We conclude that although the pulsar can contribute to the formation of the knee, it cannot be the dominant source of it and an SNR is still needed.  We also examine the possibility of the pulsar giving the peak of the extensive air shower (~EAS~) intensity observed from the region inside the Monogem Ring \\cite{Chil1}. The estimates of the gamma-ray flux produced by cosmic ray particles from this pulsar indicate that it can be the source of the observed peak, if the particles were confined within the SNR during a considerable fraction of its total age. The flux of gamma quanta at PeV energies has a high sensitivity to the duration of the confinement. The estimates of this time and of the following diffusion of cosmic rays from the confinement volume turn out to be in remarkable agreement with the time needed for these cosmic rays to propagate to the solar system and to form the observed knee in the cosmic ray energy spectrum.  Other possible mechanisms for the production of particles which could give rise to the observed narrow peak in the EAS intensity were also examined. Electrons scattered on the microwave background or on X-rays, emitted by SNR, can not be responsible for the gamma-quanta in the peak. Neutrons produced in PP - collisions or released from the disintegration of accelerated nuclei seem to be also unable to create the peak since they cannot give the observed flux.  If the experimental EAS results concerning a point-like source are confirmed, they can be important, since\\\\ (i) they will give evidence for the acceleration of protons or heavier nuclei by the pulsar; \\\\ (ii) they will give evidence for the existence of a confinement mechanism in SNR ;\\\\ (iii) they will confirm that cosmic rays produced by the Monogem Ring SNR and associated pulsar B0656+14 were released recently giving rise to the formation of the sharp knee and the observed narrow peak in the EAS intensity~;\\\\ (iv) they will give strong support for the Monogem Ring SNR and the associated pulsar B0656+14 being identified as the Single Source proposed in our Single Source Model of the knee.  A number of predictions of the examined mechanism are made. ",
        "introduction": "A few years ago we suggested the `Single Source Model' to explain the remarkable sharpness of the knee in the cosmic ray energy spectrum at $\\sim$3 PeV \\cite{EW1,EW2}, a feature noticed even in the first publication on this subject, 46 years ago \\cite{Kulik}. The model is based on the assumption that a single, relatively recent and nearby supernova remnant contributes significantly to the cosmic ray intensity at PeV energies. The sharpness is due to the cutoff in the energy spectrum of cosmic rays accelerated by SNR. According to the theoretical model of the SNR acceleration mechanism developed by Berezhko et al. \\cite{Berez} the cutoff for the acceleration in the hot and low density interstellar medium (~ISM~) is at a rigidity of $\\sim$0.4 PV. To match the position of the knee at an energy of $\\sim$3 PeV and to explain the second peak of the intensity at $\\sim$10-15 PeV, observed in most of the experiments, the model assumes that the Single Source emits predominantly medium (~oxygen~) and heavy (~iron~) nuclei with an admixture of sub-iron nuclei. This assumption is reasonable, in view of the ISM, which provides the nuclei, having been seeded by a previous supernova (~see the next paragraph~). Comparing the shape of the energy spectrum of cosmic rays from the Single Source and its total energy content with the model of SNR acceleration and the propagation of cosmic rays through the ISM we derived a likely interval of distance (~230-350 pc~) and age (~84-100 kyear~) for the Single Source \\cite{EW3}.  On the basis of our estimates of distance and age we calculated the possible flux of high energy gamma rays from the Single Source and found that it is unlikely to be observed at sub-GeV and TeV gamma rays with gamma telescopes of the present sensitivity \\cite{EW7}. The reason is that being nearby the Single Source is in our Local Superbubble with its low (~$\\sim 3\\cdot 10^{-3} cm^{-3}$~) density of target gas and the SNR should not be a discrete source , but extended with an angular radius of $\\sim$20$^\\circ$. It is difficult to detect an excess intensity from such an extended source since the estimates of the background are very unreliable. Among the sources which would satisfy these limits of distance and age we indicated the Monogem Ring and Loop I \\cite{EW4}. The same sources were also discussed in connection with their possible contribution to the flux of high energy electrons \\cite{Koba1,EW5,Koba2}  Recently, Thorsett et al. \\cite{Thors}, using the triangulation technique found the distance of the pulsar PSR 0656+14 associated with the SNR Monogem Ring. It is 288$\\pm$30 pc and its spin-down age is $\\sim$110 kyears, both of which are in remarkable agreement with our estimates for the Single Source \\cite{EW7}. Such determinations had not previously been possible from observations of the SNR itself. Thorsett et al. themselves claimed that the SNR Monogem Ring and its associated pulsar PSR 0656+14 can be the Single Source responsible for the formation of the knee.  Armenian physicists have studied the sky near the Monogem Ring in the sub-PeV and PeV range using the EAS technique and found a 6$\\sigma$ excess of the EAS intensity in one of their angular bins \\cite{Chil1}; we refer to this interesting result as the `Armenian peak'. Since their bin of $3^\\circ \\times 3^\\circ$ is narrower than the size of the SNR it is thought that the excess is not due to the extended source, but to a discrete source, viz. the pulsar.  Bhadra \\cite{Bhadr} analysed theoretically the possibility for a pulsar to be the Single Source, and concluded that the most likely pulsar candidates are Geminga and Vela.  In this paper we analyse the possibility of the pulsar PSR B0656+14, associated with the SNR Monogem Ring, being the Single Source responsible for the knee and also see to what extent the `Armenian peak' could come from the same object. Independently, we search for other evidence which might confirm the reasonableness of the peak. ",
        "conclusions": "We conclude that the pulsar B0656+14 can contribute to the intensity of cosmic rays in the knee region, but cannot be the dominant source responsible for its formation. Its contribution to the intensity of the Single Source, needed to form the sharp knee, appears not to exceed 15\\%. The SNR associated with the Monogem Ring, rather than the pulsar, still remains the most likely Single Source which gives the dominant contribution to the formation of the cosmic ray energy spectrum in the vicinity of the knee.  We have also examined the possibility of the pulsar B0656+14 giving the peak (~the`Armenian peak'~) of the EAS intensity, observed from the region inside the Monogem Ring. The estimates of the gamma-ray flux produced by cosmic ray protons from this pulsar evidence that it can be the source of the observed peak, if the protons were confined within the SNR during a considerable fraction (~$\\sim$90\\%~)of its total age. The flux of gamma quanta at PeV energies has a high sensitivity to the duration of the confinement. The estimate of this time and the following diffusion of cosmic rays from the confinement volume turns out to be in remarkable agreement with the time needed for these cosmic rays to propagate to the solar system and to form the observed knee in the cosmic ray energy spectrum.  Other possible mechanisms for the production of particles which could give rise to the Armenian peak were also examined. Electrons scattered on the microwave background or on X-rays, emitted by SNR, could not be responsible for the gamma-quanta in the peak. Neutrons produced in PP - collisions or released from the spallation of accelerated nuclei also seem to be improbable mechanisms since they cannot give the observed flux.  If the EAS results are confirmed, they will be important, since\\\\ (i) they give evidence for the possibility of the acceleration of protons by the pulsar; \\\\ (ii) they give evidence for the existence of a confinement mechanism in SNR ;\\\\ (iii) they confirm that cosmic rays produced by the Monogem Ring SNR and associated pulsar B0656+14 were released recently giving rise to the formation of the sharp knee and the observed narrow peak in the EAS intensity ;\\\\ (iv) they give strong support for the Monogem Ring being identified as the source proposed in our Single Source Model of the knee.  \\vspace{0.5cm}  {\\large{\\bf Acknowledgements}}  Authors thank A.A.Chilingarian for useful discussions and an unknown referee for the remarks. One of the authors (ADE) thanks The Royal Society for financial support."
    },
    "0404/astro-ph0404524_arXiv.txt": {
        "abstract": " ",
        "introduction": "Though X-rays dominate the energetic output from X-ray binaries, the optical emission is a potentially powerful diagnostic of the accretion flow, in particular in combination with the X-ray signal. Reprocessing of X-rays into optical light by the accretion disk provides information on the geometry (thickness as a function of distance) and state of ionization of the disk surface. Reprocessing signals of this kind have been seen in a few neutron star accreters, where the optical `echos' formed by reprocessing of X-rays from type I X-ray bursts on the neutron star surface by the disk have been observed (Matsuoka et al. 1984, Turner et al. 1985, van Paradijs et al. 1990, Kong et al. 2000).  Optical emission of an entirely different kind has been observed in three black hole candidates. In  a brief (90s) segment of simultaneous X-ray+optical data of GX 339-4 by Motch et al. (1981,1983) the optical emission was observed to show {\\it dips} some seconds before X-ray peaks. In simultaneous HST+XTE data on GRO 1655-40, Hynes et al. (1998) found the optical emission to correlate positively with the X-rays, with a delay of some seconds. The most detailed observations of this kind were obtained by Kanbach et al. (2001), on the black hole transient XTE J1118+480 (=KV UMa). These observations showed a strong correlation between X-rays and optical emission, with a cross correlation function showing both a dip preceding X-ray maximum and a following sharp peak. The shape of the cross correlation function was found to be highly variable, on time scales as short as 25s (Spruit and Kanbach 2003). One of the most curious properties of this variability is a variation of time scale: in different segments of data, the cross correlation function tends to have the same shape, except for expansion/contractraction of the time axis. The physical interpretation of these observations is still very uncertain.  The short time scales observed in the optical emission in KV UMa  (down to 30ms) and the large optical luminosity are both suggestive of an origin close to the black hole. The only radiation process with sufficient emissivity in the optical is then synchrotron or thermal (cyclo-) synchrotron emission, for which a strong magnetic field would then be required. Such strong magnetic fields, in turn, are the ingredient of choice for current models for the production of the relativistic outflows seen from several black hole transients (Mirabel and Rodr{\\'{\\i}}guez 1999). These outflows appear to occur in particular when the sources are in X-ray `hard' states, when the energetic output is dominated by photons around 100keV (Corbel et al. 2000).  In contrast with the `soft' state, understood as evidence of the theoretically predicted of accretion disk structure, the nature of the plasma producing the hard X-ray emission, and the geometry of the accretion flow in which it occurs is still unclear. In the few cases where fast correlated X-ray/optical emission has been found so far, the sources were in such hard states. Correlated X-ray/optical observations thus provide a potentially powerful new diagnostic on the uncertain accretion physics close to the hole.  A difficulty in exploiting this diagnostic is the fact that most black hole candidates are transients whose outbursts are infrequent and of relatively short duration. Other difficulties are the tendency for the sources to lie in optically obscured regions near the galactic center,  and the lack of suitably fast photometric instrumentation on most modern telescopes.  The sources selected for the observations reported here are Cir X-1 and the black hole transient XTE J1746-321 (=IGR 17464-3213 =H1743-322), which happened to be in outburst at the time of the observations. Cir X-1 is a microquasar (Stewart et al. 1993, Fender et al. 2004) in which the accreter is believed to be a neutron star (the other binary of this type being Sco X-1, Fomalont et al. 2001). Comparison of this object with black hole candidates would answer the question whether fast optical emission is connected with the presence of a black hole, or more generically with the presence of relativistic outflows. Since it is also a persistent X-ray source, it was the primary target of the observations.  The second planned target was GX 339-4, the black hole candidate where correlated X-ray/optical variability was detected for the first time (Motch et al. 1981,1983). Unlike most black hole candidates, GX 339-4 has frequent outbursts of varying strength, often several per year. It was active when our observations were planned, but it turned into an `off' state about 1 month before the observations (as it did in two previous attempts). However, another X-ray transient and possible black hole candidate, XTE J1746-321, happened to be active during the observations. Though rather obscured, it was still marginally bright enough to attempt optical observations. ",
        "conclusions": ""
    },
    "0404/astro-ph0404238_arXiv.txt": {
        "abstract": "We have modeled the time-variable profiles of the H$\\alpha$ emission line from the non-axisymmetric disk and debris tail created in the tidal disruption of a solar-type star by a $10^{6}\\,M_{\\odot}$ black hole. We find that the line profiles at these very early stages of the evolution of the post-disruption debris do not resemble the double peaked profiles expected from a rotating disk since the debris has not yet settled into such a stable structure. The predicted line profiles vary on fairly short time scales (of order hours to days). As a result of the uneven distribution of the debris and the existence of a ``tidal tail'' (the stream of returning debris), the line profiles depend sensitively on the orientation of the tail relative to the line of sight. Given the illuminating UV/X-ray light curve, we also model the H$\\alpha$ light curve from the debris. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404462_arXiv.txt": {
        "abstract": "{We present the first results of a project aiming to derive the physical properties of high-redshift lensed galaxies, intrinsically fainter than the Lyman break galaxies currently observed in the field. From FORS and ISAAC spectroscopy on the VLT, we use the full rest-frame UV-to-optical range to derive the physical properties (SFR, extinction, chemical abundances, dynamics, mass, etc) of low-luminosity $z \\sim 2$ star-forming galaxies. Although the sample is still too small for statistical studies, these results give an insight into the nature and evolutionary status of distant star-forming objects and their link with present-day galaxies. Such a project will serve as a basis for the scientific analysis of the EMIR/GOYA survey on the GTC.}    \\addkeyword{galaxies: evolution} \\addkeyword{galaxies: starburst} \\addkeyword{galaxies: abundances} \\addkeyword{galaxies: infrared} \\addkeyword{galaxies: kinematics and dynamics}  \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404181_arXiv.txt": {
        "abstract": "We discuss four recycled pulsars found in Arecibo drift-scan searches. PSR~J1944+0907 has a spin period of 5.2 ms and is isolated. The 5.8-ms pulsar J1453+19 may have a low-mass companion. The isolated 56-ms pulsar J0609+2130 is possibly the remnant of a disrupted double neutron star binary. The 41-ms pulsar J1829+2456 is in a relativistic orbit. Its companion is most likely another neutron star. ",
        "introduction": "Recycled pulsars are believed to be formed when an old neutron star is spun up through the accretion of matter from a binary companion (Bisnovatyi-Kogan \\& Komberg 1974). For those companions massive enough to explode as a supernova, the binary lifetime is relatively short ($\\sim 10^{6-7}$ yrs) and the most likely outcome is the disruption of the binary. Those systems which survive the explosion are the double neutron star (DNS) binaries.  For less massive companions, where the period of spin-up is longer ($\\sim 10^8$ yr), the collapse leaves a white dwarf star in orbit around a rapidly spinning millisecond pulsar (MSP).  Of the roughly 1700 known radio pulsars, only 100 are MSPs and less than 10 are in DNS systems.  In order to understand the population of recycled pulsars in detail, a larger sample is required.  Here we describe the implications for four recycled pulsars recently detected in Arecibo drift-scan searches. ",
        "conclusions": ""
    },
    "0404/astro-ph0404148_arXiv.txt": {
        "abstract": "We have obtained high resolution Echelle spectra (R = 30,000 - 50,000) of the Na D absorption doublet ($\\lambda\\lambda$5890, 5896) for six dwarf starburst galaxies and two more luminous starbursts:  M82 and NGC 1614.  The absorption features were separated into multiple components and separated into stellar and interstellar parts based on kinematics. We find that three of the dwarfs show outflows, with an average blueshift of 27 km s$^{-1}$.  This is small compared to the highest velocity components in NGC 1614 and M82 (blueshifted by 150 km s$^{-1}$ and 91 km s$^{-1}$, respectively); these two brighter galaxies also show more complex absorption profiles than the dwarfs. None of the outflow speeds clearly exceed the escape velocity of the host galaxy.  Sightlines in NGC 2363 and NGC4214 apparently intersect expanding shells.  We compare the shocked gas velocity (v$_{NaD}$) to the ionized gas velocity (v$_{H\\alpha}$) and interpret the velocity difference as either a trapped ionization front (NGC 4214) or a leaky \\ion{H}{2} region (NGC 2363).  The dwarfs show N$_{NaD}$ = 10$^{11.8-13.7}$ cm$^{-2}$, while the Na D columns in M82 and NGC 1614 are 10$^{13.7}$ cm$^{-2}$ and 10$^{14.0}$ cm$^{-2}$, respectively.  The mass of expelled gas is highly sensitive to outflow geometry, dust depletion, and ionization fraction, but with a simple shell model we estimate neutral outflow gas masses from $\\sim$10$^6$ M$_\\odot$ to $\\sim$10$^{10}$ M$_\\odot$. ",
        "introduction": "Galactic winds are driven by supernovae and stellar winds in starburst galaxies.  They drive outflows of metals and dust from the interstellar medium (ISM) into the galactic halo and even into the intergalactic medium (IGM).  These outflows may also heat the intracluster medium and enrich it with metals, although the amount of matter permanently escaping the galaxy is still debated.  The amount of interstellar gas expelled into the IGM by a starburst in a ``blowout'' is thought to depend on the star formation rate within the galaxy, the escape velocity of the galaxy, and the presence of an extended, low-density gaseous halo (Silich \\& Tenorio-Tagle 1998; Legrand et al. 2001).  In dwarf starbursts, the mass-loss measurements have focused on the ionized components of this multi-phase flow.  In this paper, we explore the kinematics of the cold, neutral clouds in dwarf starburst outflows.  Due to their low gravitational potential, dwarf starburst galaxies have small binding energies and may be particularly vulnerable to large mass loss via superwinds.  Previous works have studied X-ray emission from hot gas within the starbursts (Dahlem, Weaver, \\& Heckman 1998; Martin, Kobulnicky, \\& Heckman 2002), and optical line emission (e.g. Lehnert \\& Heckman 1996) and absorption (Martin \\& Armus 2004; Rupke, Vielleux, \\& Sanders 2002; Heckman et al. 2000 - hereafter HLSA) from larger, more luminous galaxies. Other studies have observed dwarf galaxies in emission (Marlowe et al. 1995; Martin 1999), but ours is among the first to study dwarf starbursts in neutral absorption lines at optical wavelengths.  We present a small sample of dwarf starburst galaxies and attempt to quantify mass flux in the cold neutral medium (CNM).  The advantage of looking at absorption lines is that we can eliminate all ambiguity in the direction of the radial flow of gas.  Using the starburst itself as a background continuum source, it is known absolutely whether the gas we see in absorption is falling in (redshifted) or expanding outward (blueshifted). Moreover, unlike emission lines, whose strength scales as density squared, the strength of spectral lines seen in absorption are directly proportional to the column density of the neutral gas along the line of sight.   Several previous studies (Phillips 1993;  Lequeux et al. 1995; Gonzalez-Delgado et al. 1998; HLSA; Martin \\& Armus 2004) have detected interstellar absorption lines in starburst galaxies of various sizes and morphologies, finding blueshifted neutral gas and related large-scale outflows of metals.  HLSA find an average outflow speed of $\\sim$100 km s$^{-1}$ for luminous infrared galaxies (LIRGs) and more recently, Martin \\& Armus (2004) and  Rupke et al. (2002) find outflow speeds of up to 700 km s$^{-1}$ for two samples of ultraluminous infrared galaxies (ULIRGs).  We focus our observations on the \\ion{Na}{1} absorption doublet (5889.95~\\AA\\, and  5895.92~\\AA, also called the Na D doublet) because it is stronger than any other optical resonance line (e.g. \\ion{K}{1}, \\ion{Ca}{2}); most other cosmically abundant resonance lines are produced in the UV and therefore cannot be observed in nearby galaxies with ground-based telescopes.  Na D is a good tracer of cold neutral gas because its ionization potential (5.14 eV) is less than that of hydrogen, so we can be sure any \\ion{Na}{1} absorption is occuring in a cold region of the ISM.  It is also important to note that sodium is not as strongly depleted in diffuse, low-velocity clouds as \\ion{K}{1} or \\ion{Ca}{2} (Spitzer 1968; Savage \\& Sembach 1996).   By using high resolution spectroscopy we can resolve individual, kinematically distinct doublet components within the absorption profiles. At lower resolution, deconvolving the absorption components of disparate origins is a more difficult task (see, e.g., HLSA).  It is our goal to observationally determine the kinematics of the CNM in our sample of starburst galaxies.  In the second section of this paper, we discuss the observations, data reduction, and basic data analysis.  In \\S3 we detail our method of separating stellar and interstellar sodium, discuss the kinematics and widths of the absorption features, and then use the interstellar sodium to determine a column density.  In \\S4 we present our interpretation of the data, including simple models for outflow scenarios in the individual galaxies.  The final section summarizes our major results. ",
        "conclusions": "We have obtained high-resolution Echelle spectra of six dwarf starburst galaxies, NGC 1614, and M82 in order to study the cold interstellar gas in dwarf starbursts.  Interstellar neutral sodium column densities were obtained by measuring the equivalent width of the Na D absorption doublet.  We find that out of the eight galaxies, NGC 1569, NGC 1614, NGC 4214-2, NGC 4449, and M82 unambiguously show interstellar sodium absorption, while NGC 2363, NGC 4214-1, NGC 5253, and I Zw 18 do not. The dwarf galaxies NGC 1569, NGC 4214, and NGC 4449 exhibit single-component outflows of neutral gas.  NGC 1614 and M82 have multiple interstellar components, some outflowing and some infalling, and both galaxies show more absorption from outflowing gas than any dwarf galaxy.  The dwarf galaxies show trends similar to brighter, larger galaxies, but on a smaller scale.  While samples of LIRGs and ULIRGs are shown to have average outflow speeds of $\\sim$100 km s$^{-1}$ (HLSA) and $\\sim$700 km s$^{-1}$ (Rupke et al. 2002; Martin \\& Armus 2004), respectively, the three dwarfs show an average outflow speed of only $\\sim$30 km s$^{-1}$.   Additionally, most ULIRGs (e.g. 73\\% in Rupke et al. 2002) show an outflow, whereas only half of the galaxies in our sample have an outflow region.  Spectral lines in dwarf galaxies also have  smaller velocity spreads.  It is particularly interesting to see how complex the absorption line spectra become at high resolution.  M82 is an excellent example of an absorption line system previously believed to be a single pair of lines, but within each member of the doublet we find a wealth of structure. We resolve at least five line pairs, showing the incredibly complicated nature of the CNM in this galaxy.  This is a very intriguing result and prompts us to wonder how many of the galaxies in previous samples could be resolved into complex systems of multiple components.  NGC 1614 shows complex absorption as well, though it has a far smoother profile in Na D (and \\ion{K}{1}) than M82. The three dwarfs with outflows do not seem to exhibit this behavior, showing far simpler spectra.  For the first time we are able to combine measurements of the kinematics of the warm, ionized gas in emission with absorption spectra of the cold, neutral gas.  For sightlines intersecting a single shell, we can determine whether the ionization front is trapped inside the shock front.  This works extremely well for NGC 2363, which is a simple expanding bubble with an ionization front that has expanded beyond the shock front.  NGC 4214-2 also presents a relatively simple shell, and this galaxy also shows Na D absorption. From the kinematics we postulate that the ionization front is expanding faster than the shock front, though it has clearly not ``caught up'' yet, as there is still a concentric, outer expanding bubble of cold, neutral gas.  In other words, the ionization front is still trapped in the shell.  The kinematics of Na D and H$\\alpha$ are very different in other galaxies, and cannot be easily explained by this straightforward picture.  Combining our measurements of \\ion{Na}{1} column density with previous observations of superbubbles and supershells in the sample galaxies, we have parameterized the mass of neutral gas (or a limiting case thereof) flowing out of the starbursting region.  Using a simple spherical shell model, we find the total neutral gas outflow masses to be roughly 10$^6$ to 10$^{10}$ M$_\\odot$.  NGC 4449 shows far more outflowing cold gas than any other dwarf galaxy.  Compared to the measurements of outflowing warm and hot gas from these galaxies (Martin 1998), it is likely that the bulk of the energy in the outflow is carried by the warm and hot gas, rather than cold gas."
    },
    "0404/astro-ph0404418_arXiv.txt": {
        "abstract": "We studied the X-ray variability of sources detected in the {\\it Chandra} Deep Field South \\citep{G02}, nearly all of which are low to moderate $z$ AGN \\citep{Tozzi01}. We find that 45\\% of the sources with $>100$ counts exhibit significant variability on timescales ranging from a day up to a year. The fraction of sources found to be variable increases with observed flux, suggesting that $>90\\%$ of all AGNs possess intrinsic variability. We also find that the fraction of variable sources appears to decrease with increasing intrinsic absorption; a lack of variability in hard, absorbed AGNs could be due to an increased contribution of reflected X-rays to the total flux. We do not detect significant {\\it spectral} variability ($\\Delta \\Gamma>0.2$) in the majority ($\\sim 70\\%$) of our sources. In half of the remaining $30\\%$, the hardness ratio is anti-correlated with flux, mimicking the high/soft--low/hard states of galactic sources. The X-ray variability appears anti-correlated with the luminosity of the sources, in agreement with previous studies. High redshift sources, however, have larger variability amplitudes than expected from extrapolations of their low-z counterparts, suggesting a possible evolution in the accretion rate and/or size of the X-ray emitting region. Finally, we discuss some effects that may produce the observed decrease in the fraction of variable sources from $z=0.5$ out to $z=2$. ",
        "introduction": "Active Galactic Nuclei (AGN) are found to vary on timescales ranging from minutes to years. Variability studies have proven to be an important tool to investigate the innermost regions of AGNs, which cannot be resolved with the current generation of astronomical instruments. The timescale of the variability, its spectral dependence and the correlation between different line and continuum components provide fundamental clues on the nature of the physical processes which occur near to the central Black Hole (BH) and on the size and relation between the different regions producing the observed emission \\cite[see][for a comprehensive review]{Ulrich97}.  In the X-ray band, AGNs show faster variability than in any other band, consistent with the X-ray emission originating from a small region very close to the central BH \\citep[see][]{Mushotzky93}. Several promising models have been proposed to explain the fast intrinsic variability and its characteristics: a single coherent oscillator \\cite[e.g.][]{Almaini00}, variable decay shot-noise due to a superposition of individual flares \\cite[e.g.][]{Lehto89}, and bright rotating spots spiraling around the BH \\citep{Abramowicz91}. All these models predict a dependence on the BH mass, accretion rate and the size of the X-ray emitting region which may in turn explain the observed dependence of variability on the AGN luminosity \\citep[e.g.][]{Law93,GML93,Nandra97} and spectral properties \\citep{Nandra97,Turner99,Fio98}.  Most variability studies in the X-ray band are based on samples of bright and nearby AGNs for which good quality data are available from past X-ray missions. Moreover they usually sample short timescales from hours to days. Only for the nearest sources have long monitoring campaigns allowed a detailed study of long-term variability up to several years. The {\\it Chandra} Deep Field South \\cite[hereafter CDFS]{G01} represents one of the deepest observations of the Universe in the X-ray band (the only deeper sample is in the Northern {\\it Chandra} Deep Field). \\cite{Rosati02} detected 346 sources in the CDFS, resolving $>90\\%$ of the hard X-ray background. The X-ray luminosities and spectral properties of the sources revealed that $\\sim 80\\%$ of them are AGNs \\citep{G01,Tozzi01}. Optical follow-up studies showed that they are equally divided among moderate redshift ($z<1.6$) Type II AGNs with large intrinsic absorption ($\\log N_H>21.5$ cm$^{-2}$), and unabsorbed Type I AGNs covering a large range of redshifts up to $z\\sim 3.6$, both hosted by a broad range of galaxy types \\citep{Schreier01,Koekemoer02,Rosati02}. The remaining $20\\%$ of the sample is composed of a mixture of normal galaxies, starburst galaxies, galaxy clusters \\citep{G01,Koekemoer02} and individual off-nuclear X-ray sources \\citep{Horn04}.  Being composed of several concentric exposures obtained over a period of $\\sim 1$ year, the CDFS represent an ideal dataset to study the X-ray variability of faint and distant AGNs, on timescales ranging from a day up to several months. In the current study we measure the variability properties of the AGN sample and correlate the variability with other X-ray properties (flux, luminosity, hardness ratio, etc.), compare the properties of these moderate-$z$ AGN to those of the well studied nearby AGN population and study their evolution with look-back time. The paper is organized as follows: in $\\S$ 2 we describe the dataset, in $\\S$ 3 we illustrate the variability analysis, in $\\S$ 4 and 5 we present and discuss our results and in $\\S$ 6 we draw the conclusions.  Throughout the paper we assume a cosmological model with $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m=0.3$ and $\\Omega_\\Lambda=0.7$. ",
        "conclusions": "We exploited the deep (1 Ms) and high-resolution data of the {\\it Chandra} Deep Field South to study the temporal variability of AGNs on timescales ranging from a fraction of a day up to one year. While similar studies have been performed on samples of nearby AGNs, our sample primarily contains moderate luminosity AGNs out to $z\\simeq 3.5$.  We detect significant variability in $>50\\%$ of the AGN population in the 0.5-7 keV band and suggest that as many as $90\\%$ of all AGN in our sample may be variable. The variability exists on all timescales that we can observe, but our sources are twice as variable on timescales of $>100$ days than within a few ($<10$) days. All sources show some variability on the shortest timescales indicating that part of the flux variations are produced on spatial scales $<2\\times 10^{-3}$ pc, corresponding to the inner part of the accretion disk and of the Broad Emission Line region.  We also find that unabsorbed sources ($\\log N_H<22.5$ cm$^{-2}$) appear intrinsically more variable than absorbed ones at the $2\\sigma$ significance level, suggesting that the lack of variability in hard, absorbed sources is due to an increased contribution of reflected X-rays to the total flux. On the other hand we do not find strong ($\\Delta\\Gamma >0.2$) spectral changes associated with the flux variations in the majority of our sources.  Our data are consistent with an anti-correlation between X-ray luminosity and variability, similar to that found in local AGNs. However, this relation seems to evolve with redshift, with X-ray sources having been more variable in the past. This trend may reflect an evolution in the accretion rate and/or size of the X-ray emitting region with look-back time.  Finally we find a possible decrease in the fraction of variable sources with redshift for $0.5\\lesssim z<2$. This result could be explained by assuming that AGN variability has not evolved since $z\\sim 2$. Alternatively, environmental effects triggering X-ray variability or an evolution in the Type I/Type II AGN ratio could produce the observed trend. Future work using spectrosopic redshift and an enlarged sample may be able to address this issue.  \\clearpage \\appendix \\centerline{\\bf\\large Appendix}"
    },
    "0404/astro-ph0404132_arXiv.txt": {
        "abstract": "{Spatially-resolved gas pressure maps of the Coma galaxy cluster are obtained from a mosaic of XMM-Newton observations in the scale range between a resolution of 20\\,kpc and an extent of 2.8\\,Mpc. A Fourier analysis of the data reveals the presence of a scale-invariant pressure fluctuation spectrum in the range between 40 and 90\\,kpc and is found to be well described by a projected Kolmogorov/Oboukhov-type turbulence spectrum. Deprojection and integration of the spectrum yields the lower limit of $\\sim 10$ percent of the total intracluster medium pressure in turbulent form. The results also provide observational constraints on the viscosity of the gas. ",
        "introduction": "In hierarchical structure formation scenarios clusters grow via accretion and merging of smaller subclumps. Gas accreting onto clusters of galaxies has bulk velocities of about $v=1900\\,(T/6.7\\,{\\rm keV})^{0.52}\\,{\\rm km}\\,{\\rm s}^{-1}$ at 1\\,Mpc (e.g. Miniati et al. 2000), where $T$ is the mean X-ray temperature of the intracluster medium (ICM). This velocity is comparable to the expected sound speed of 1000-1500\\,km/s of the ICM. Accretion flows through filaments and sheets are highly asymmetric and produce complex patterns which can survive for long time-scales in the ICM (Miniati et al. 2000).  Simulations by Norman \\& Bryan (1999) predict that the turbulent pressure in the ICM can account for up to 20\\% of the thermal pressure.  We thus expect some measurable effects of turbulence in the ICM of clusters of galaxies.  Concerning X-ray data, Inogamov \\& Sunyaev (2003) propose a study of spectral line profiles as a useful diagnostic tool of turbulent flows in the ICM which could be measured with the future ASTRO-E2 satellite. Vogt \\& En{\\ss}lin (2003) propose the application of Faraday Rotation measures to test turbulence in the ICM, and claim that for a few clusters a Kolmogorov spectrum seems to be plausible.  In the present investigation we show that turbulence in the ICM can be probed directly with pressure maps provided by the XMM-Newton satellite as a result of its high sensitivity and excellent spectral capabilities.  Section\\,\\ref{PHENOMEN} summarizes the basic phenomena related to turbulent flows. In Sect.\\,\\ref{EFFECTS} we give a simple analytic treatment of projection effects introduced through observation. In Sect.\\,\\ref{MAPS} we present the observational data and describe how our X-ray temperature and pressure maps are constructed. Based on the direct comparsion of local temperature and density measurements, we give in Sect.\\,\\ref{GENERAL} some arguments that their observed fluctuations appear to be almost adiabatic. The same statistical analysis also suggests the absence of pronounced contact discontinuities and strong shocks. These observations provide a baseline consistent with the presence of a turbulent flow. Therefore, we study in detail in Sect.\\,\\ref{STRUCT} the measured pressure spectrum in Fourier-transformed $k$ space and discuss its interpretation in Sect.\\,\\ref{DISCUSS}. For all computations a flat geometry and a Hubble constant of $H_0=50\\,{\\rm km}\\,{\\rm s}^{-1}\\,{\\rm Mpc}^{-1}$ are used. We assume a distance of $139$\\,Mpc to the Coma cluster so that 1\\,arcmin corresponds to about 40\\,kpc.  \\begin{figure} \\vspace{-0.2cm} \\centerline{\\hspace{-0.5cm} \\psfig{figure=1039fg1a.ps,height=6.0cm,width=8.0cm}} \\vspace{0.0cm} \\centerline{\\hspace{-0.5cm} \\psfig{figure=1039fg1b.ps,height=6.0cm,width=8.0cm}} \\vspace{-0.0cm} \\caption{\\small Angular profile of the weight function ($\\beta$ profile squared) integrated along the $z$ direction for various integration lengths. The lowest/highest contours indicate 10/100\\% contributions, and the intermediate lines the contributions in steps of 10\\%. The upper panel shows the contributions on cluster scales and the lower panel the contributions on the maximum scales of the observed turbulence structures (see Sect.\\,\\ref{DISCUSS}). For the brightness distribution we assume a $\\beta$ profile with the parameter values $\\beta=0.75$ and $r_{\\rm c}=420$\\,kpc.} \\label{FIG_PROF} \\end{figure}  \\begin{figure} \\vspace{-0.5cm} \\centerline{\\hspace{0.0cm} \\psfig{figure=1039fig2.ps,height=8.5cm,width=8.5cm}} \\vspace{-0.45cm} \\caption{\\small Pressure power spectra with an intrinsic slope of $n=-7/3$ as expected for a Kolmogorov/Oboukhov turbulence ($r_{\\rm c}=0$), and its projection (420\\,kpc), as seen along the $z$ direction through the central plane of a cluster with a core radius ($r_{\\rm c}=420$\\,kpc) and slope parameter ($\\beta=0.75$) as measured for the Coma cluster. In order to illustrate the projection effects over a large scale range, we did not introduce any characteristic scale which limits the spectra at large and small scales.} \\label{FIG_MODEL} \\end{figure} ",
        "conclusions": "\\label{DISCUSS}  The present investigation aims to detect turbulence in the ICM of the Coma cluster using the pressure spectrum. Under certain approximations, one also expects a scale-invariant spectrum of temperature fluctuations to be a probe of ICM turbulence (e.g. Lesieur 1997). However, this relies on the assumption that temperature behaves as a passive scalar. Once this has been verified empirically, the almost uniform distribution of the temperature over scales $\\gg r_{\\rm c}$ allows a cleaner distinction between small-scale turbulent substructures and the large-scale cluster profile. However, in reality temperature maps are affected by cold fronts and other contact discontinuities which contaminate the diagnostic maps. On the other hand, pressure maps have a clear relation to velocity ($P=\\rho v^2$) and are not significantly contaminated by contact discontinuities. Therefore, we regard pressure as a more direct probe of ICM turbulence.  The mosaic of XMM-Newton observations is well-suited for the detection of turbulence in the Coma cluster because it allows a better geometric discrimination between pressure variations originating from the overall cluster profile, and substructure superimposed onto it. This transition occurs at about 150\\,kpc.  The measured temperature and density gradients (Fig.\\,\\ref{FIG_GAMMA}) suggest that the substructures have an adiabatic exponent of $\\gamma\\,\\approx\\,4/3$, which is close to the adiabatic case of an ideal monoatomic gas. On the other hand, contact discontinuities and strong shocks seem to be less likely in the core region, consistent with hydrodynamical simulations (Miniati et al. 2000, Miniati 2003). In addition, the statistics of the residual pressure fluctuations appear quite Gaussian (Fig.\\,\\ref{FIG_HISTO}) emphasising their random nature. Their Fourier power spectrum thus completely summarizes the fluctuating pressure field and can be used to obtain observational evidence for the presence of turbulent flows which are characterized by a Kolmogorov/Oboukhov-like spectrum.  Figure\\,\\ref{FIG_PKFALL} shows the combined power spectrum of the Coma cluster on scales between 40 and 90\\,kpc. A scale-invariant range of the spectrum is indicated and suggests the detection of an inertial range of a turbulent ICM. Theoretical three-dimensional power spectra, \\begin{eqnarray}\\label{FIT} {\\cal P}_{\\rm 3D}(k)\\,&=&\\,2\\pi^2\\,k^{-2}\\,E_{\\rm 3D}(k)\\,=\\, 2\\pi^2\\,C_{\\rm P}\\,\\epsilon^{4/3}\\,k^{n-2}\\,\\\\ &=&\\,2\\,\\pi^2\\,C\\,k^{n-2}\\,, \\end{eqnarray} are transformed with Eq.\\,(\\ref{E30}) into their two-dimensional counterparts (dashed lines in Fig.\\,\\ref{FIG_PKFALL}). From the comparison of observed and theoretical spectra we see that on scales between 40 and 60\\,kpc, the observed power spectrum has a slope between $n=-7/3$ and $-5/3$. This slope corresponds to the spectral slope of the Fourier-transformed Kolmogorov/Oboukhov structure function (Eq.\\,\\ref{OB1}). On scales between 60 and 90\\,kpc the spectrum bends towards smaller slopes between $n=-5/3$ and $-1/3$. The normalization constants range from $C=C_{\\rm P}\\epsilon^{4/3}=0.0063\\,{\\rm kpc}^{-4/3}$ for $n=-7/3$, to $C=0.470\\,{\\rm kpc}^{2/3}$ for $n=-1/3$. An absolute calibration of the pressure used in our analyses is in preparation and will give a direct estimate of $\\epsilon$. In this respect it would be interesting to compare this rate of kinetic energy transport with the observed X-ray luminosity.  \\begin{figure} \\vspace{-0.0cm} \\centerline{\\hspace{-0.3cm} \\psfig{figure=1039fg13.ps,width=9.0cm}} \\vspace{-0.00cm} \\caption{\\small Contribution of the turbulent pressure to the total thermal pressure (contour lines of equal percentage) for a Kolmogorov/Oboukhov spectrum with the slope $n=-7/3$ and the amplitude $C=C_{\\rm P}\\epsilon^{4/3}=0.0063\\,{\\rm kpc}^{-4/3}$. The power spectrum is integrated between the injection and the dissipation scale.} \\label{FIG_CONTRI} \\end{figure}  The power spectrum shown in Fig.\\,\\ref{FIG_PKF10} allows a first estimation of the location of the characteristic scale where the spectrum sharply drops towards larger scales. This scale is at approximately $\\lambda_{\\rm i}\\,\\approx\\,100\\,{\\rm kpc}$ and should be regarded as a lower limit because of possible contaminations by the global cluster profile. This scale should also roughly correspond to the injection scale (e.g., Lesieur 1997), and it is similar to estimates for the impact parameter of merging clusters based on kinematics and tidal torque-based arguments (e.g., Sarazin 2002).  The integral of the power spectrum (Eq.\\,\\ref{E50}) is expected to give important information about the energy deposited in turbulent motion. For the scale range between 40 and 90\\,kpc, the slope and amplitude parameters derived above yield relative contributions of the turbulent pressure to the thermal pressure between 7.4 percent for $n=-7/3$, and 6.6 percent for $n=-1/3$. The largest possible contributions are obtained with the $n=-7/3$ spectrum. Therefore, we computed for this spectrum the relative contribution for different minimum scales, i.e., lower integration limits of the inertial range (dissipation scale), and maximum scales, i.e., upper integration limits (injection scale).  Figure\\,\\ref{FIG_CONTRI} shows that for a fixed turbulence spectrum the relative contribution is mainly determined by the value of the injection scale $\\lambda_{\\rm i}$. We do not see any turbulent eddies of the size of the core radius of 420\\,pc in Fig.\\,\\ref{FIG_PWLZ} which could have erroneously been subtracted by the Fourier low-pass filter -- although they still might be present, but are difficult to discriminate from the global cluster profile. Therefore, the relative contribution of the turbulent pressure to the thermal pressure should be smaller than 25 percent. If we take the indication for a turnover in the power spectrum shown in Fig.\\,\\ref{FIG_PKF10} at $\\lambda_{\\rm i}=100\\,{\\rm kpc}$ as the injection scale, then we would get a lower limit of about 10 percent. The simulations of Norman \\& Bryan (1999) suggest additional support by turbulent pressure of about 20 percent, averaged over the cluster (5 to 35 percent between core and virial radius), which is apparently of the same order as the present observational limit. However, further study is definitely required in order to establish how the observational quantities relate to the simulation results.  For the observed turbulent ICM we can now estimate the kinematic viscosity by assuming that magnetic fields have a negligible effect (see below). For a turbulent flow the Reynolds number of the global fluid $\\Re$ measured at the injection scale $\\lambda_{\\rm i}$, and the Reynolds number $\\Re_{\\rm d}$ measured at the dissipation scale $\\lambda_{\\rm d}$ are related by $\\Re/\\Re_{\\rm d}=(\\lambda_{\\rm i}/\\lambda_{\\rm d})^{4/3}$. The power spectra do not show any tendency to decrease at $\\lambda=30$\\,kpc (Fig.\\,\\ref{FIG_PKF5}). Therefore, $\\lambda_{\\rm d}$ is smaller than 30\\,kpc. In the following we assume a fiducial value of $\\lambda_{\\rm d}=10$\\,kpc. The turbulent flow in the central region of Coma can thus be characterized by $\\Re/\\Re_{\\rm d}\\,>\\,20$. Reynolds numbers at dissipation scales are expected to be above unity so that $\\Re$ will have values in excess of 20. Although this estimate is rather conservative, this is the best that can be obtained by direct observations at the moment.  For the viscosity we further need the velocity at the injection scale. This number can be obtained from hydrodynamical simulations (Miniati et al., in preparation), which typically give for an 8\\,keV cluster a dispersion turbulent velocity of $v_{\\lambda_{\\rm i}}=250\\,{\\rm km}\\,{\\rm s}^{-1}$ on scales of $\\lambda_{\\rm i}\\approx 100$\\,kpc. This provides a quite reliable upper limit to the kinematic viscosity of \\begin{equation}\\label{NU} \\nu<3\\cdot 10^{29}\\left(\\frac{v_{\\lambda_{\\rm i}}}{250\\,{\\rm km}{\\rm s}^{-1}}\\right)\\left(\\frac{\\lambda_{\\rm i}}{100\\,{\\rm kpc}}\\right) \\left(\\frac{\\Re}{20}\\right)^{-1}\\, \\left[\\frac{{\\rm cm}^2}{s}\\right]\\,. \\end{equation} Note that the coherence lengths of magnetic fields in the cores of galaxy clusters as obtained from Faraday Rotation measurements are about 5-10\\,kpc (e.g., Taylor \\& Perley 1993) and thus below the scale range covered by the present data. Therefore, we regard the upper limit (Eq.\\,\\ref{NU}) as not significantly affected by magnetic fields.  Fabian et al. (2003, see also Reynolds et al. 2004) assume laminar flow of the ICM around the radio galaxy NGC\\,1275 with $v_\\lambda=700\\,{\\rm km}\\,{\\rm s}^{-1}$ in the centre of the Perseus cluster on $\\lambda=14\\,{\\rm kpc}$ scales. From the laminar appearance of the filaments they assume that the effective Reynolds number is less than 1000 so that they estimate the lower limit $\\nu>4\\cdot 10^{27}[\\frac{{\\rm cm}^2}{s}]$. The upper limit obtained from a turbulent regime and the lower limit obtained from a laminar regime can be used to estimate the range 10--30\\,kpc where the transition from a turbulent to a laminar flow could occur. This corresponds to a dissipation scale of the ICM in the same range. A remark of caution is, however, necessary here, because we compare two different situations (merger driven turbulence versus AGN driven turbulence, and a bulk ICM in Coma versus condensed warm HII-gas in the NGC\\,1275 halo), and it is not fully clear in how far they are comparable.  Shibata et al. (2001) determined the 2-point angular correlation function of hardness ratios as a measure of the temperature fluctuations detected with ASCA over an area of 19 square degrees in the Virgo cluster. A significant excess of the correlation amplitude is found at 300\\,kpc. They interpreted the random temperature fluctuations in Virgo-North as local heating of infalling galaxy groups.  Future investigations should measure the pressure spectrum of the Coma cluster more accurately down to 5\\,kpc so that the combination with the present measurements would give information about the ICM in the Coma cluster from 5--2800\\,kpc. This could give us tight constraints on the type of gas turbulence, its energy content, the importance of magnetic fields, and on the viscosity of the ICM."
    },
    "0404/astro-ph0404304_arXiv.txt": {
        "abstract": "We study in detail the properties of clusters of ultra high energy cosmic ray events, looking in particular to their angular correlation function, to the relative frequency of clusters with different multiplicities (e.g. doublets vs. triplets or quadruplets), as well as the way in which these quantities should evolve for different ultra high energy cosmic rays source scenarios as a function of the experimental exposure achieved. We identify some useful tools which can help to characterise the nature of the cosmic ray sources in a more precise way after a modest increase in statistics will be achieved in the very near future, even before strong signals from individual sources become eventually observable. ",
        "introduction": "In spite of the big effort done over the last decades to unravel the origin and nature of the cosmic rays (CRs) at ultra high energies (UHE), most of the relevant questions related to them remain still unanswered \\cite{cr04}. Certainly the main obstacle has been the scarcity of data, which is due to the smallness of the fluxes at the highest energies (e.g. $\\sim 1/{\\rm km}^2$yr for $E>10^{19}$~eV). The observed isotropy of CRs above the ankle (i.e. for $E>5\\times 10^{18}$~eV), suggests that at these energies CRs are of extragalactic origin, but the nature of these particles (whether they are protons, heavier nuclei, photons or even neutrinos with stronger interactions), the nature of their sources (e.g. steady ones like active galactic nuclei jets, transient ones like gamma ray burst (GRBs) or diffuse ones like decaying topological defects) and the mechanism responsible for their acceleration in bottom-up scenarios or production in top-down ones (and hence the original source spectra) are all under strong debate. All these questions are further challenged by the (lack of) observation of the theoretically expected GZK suppression above $\\sim 6\\times 10^{19}$~eV, due to the attenuation that CR protons coming from far away sources should suffer due to their interactions with CMB photons (or similarly due to the photodisintegration processes that would affect the propagation of heavier nuclei). Clearly a big boost in our understanding of these issues could be obtained if individual CR sources were identified, and this search is indeed the main purpose of the next generation of observatories, like the Pierre Auger one now under completion.  There has been a lot of discussion during the last years about possible hints in the existing data of an excess of clustering in the highest energy range ($E>4\\times 10^{19}$~eV), that could be already indicating the possible emergence of the first strong CR sources out of a more diffuse background of unresolved sources. This excess clustering (with respect to the expectations from chance coincidences of an isotropic distribution) has been observed by the AGASA group \\cite{agasa}(which detected 4 doublets and one triplet within 2.5$^\\circ$, out of a total of 57 events\\footnote{We are not including in our discussion one event with energy slightly below $4\\times 10^{19}$~eV, included by AGASA because it lead to an extra doublet.}) and also in a combined analysis of data from several experiments \\cite{uc00} (Yakutsk, Haverah Park, AGASA and Volcano Ranch, which obtained 8 doublets and 2 triplets within $4^\\circ$ out of a total of 92 events), although it does not seem to be present in Fly's Eye or Hi Res data \\cite{hires}.  In this work we will study in detail the properties of the event clusters expected in different scenarios of UHECR sources, with the aim of devising the most appropriate tools to discriminate among them. The results we obtain suggest that a modest increase in the present statistics (as could be achieved after one year of AUGER observations) will allow to clarify significantly the existing debate, even before any strong (and hence unambiguous) signal from an individual CR source is observed. ",
        "conclusions": "We have seen that the statistics of clusters of events with different multiplicities is a very useful tool that may help characterise the nature of UHECR sources. We have analysed the way in which the relative number of clusters with different multiplicities, their change with increasing exposure, and their distribution with declination can help discriminate among alternative source scenarios. In particular, one of the most striking differences between the individual sources and chance coincidences scenarios for clustering is the expected hierarchy of the number of clusters with different multiplicities. For example, the ratio between the number of triplets or higher multiplets and doublets is $\\sim 1/3$ for the steady source model (or 0.6 for the bursting one), while it is only $\\sim n/3N$ for the chance coincidences, where $N\\sim 10^3(2.5^\\circ/\\theta)^2$. In case the clusters are dominated by individual sources, we see that with significant statistics this may even give a handle to discriminate between bursting and steady source scenarios. Moreover, the evolution of this ratio (which is proportional to the exposure for chance coincidences, while independent of it for individual sources) will give a further tool to discriminate among scenarios. On the other hand, the relative fraction of events in clusters (with respect to the unclustered ones) is only $\\sim n/2N$ for the chance coincidences of an isotropic distribution, but could be significantly enhanced if the distribution of sources is anisotropic. This should be however testable by studying the overall event distribution once larger statistics are achieved. In the case of clusters produced by individual sources, this fraction depends instead sensitively on the amount of very faint sources present.  We have applied our analysis to the published set of 57 AGASA events above $4\\times 10^{19}$ eV just for illustrative purposes, since it is clearly insufficient to draw definitive conclusions. Nevertheless, our results indicate that a modest increase in the number of events may already give much stronger hints about the source properties. In this respect, even the release of the latest AGASA data, which represents a 50\\% increase with respect to the published data, could have some impact. For instance, it would be natural to expect that the triplet becomes a quadruplet, and some doublets become triplets, if they are due to individual sources, while this would be very unlikely if they arose from chance coincidences.  We have also applied an alternative statistical tool to analyse UHECR clustering: the nearest neighbour distribution. It has the advantage over the pair autocorrelation method of being less sensitive to binning, and that it depends not just on the pair correlations but on correlations to all orders \\cite{wh79}. Applied to AGASA data, it gives an alternative measure of the small angular scale clustering, and also points out to an excess of clustering with respect to random expectations which persists up to $15^\\circ$, that manifests in the second neighbours cumulative distribution (Figure~\\ref{nn}), and which is also hinted by the defect of doublets and excess of triplets at scales up to $15^\\circ$ in Figure~\\ref{pdt}. Angular spread of events in a cluster due to deflections by magnetic fields, and structure in the spatial distribution of sources, are among conceivable causes of such signal, which should be looked upon with better statistics."
    },
    "0404/astro-ph0404074_arXiv.txt": {
        "abstract": "We address some current theoretical issues around ultra-high energy cosmic rays. We recall that scenarios producing more $\\gamma-$rays than cosmic rays up to high redshift can in general only provide a sub-dominant contribution to the ultra-high energy cosmic ray flux. This includes extra-galactic top-down and the Z-burst scenarios. Finally we discuss the influence of large scale cosmic magnetic fields on ultra-high energy cosmic ray propagation which is currently hard to quantify. The views presented here represent the authors perspective. ",
        "introduction": "High energy cosmic ray (CR) particles are shielded by Earth's atmosphere and reveal their existence on the ground only by indirect effects such as ionization and showers of secondary charged particles covering areas up to many km$^2$ for the highest energy particles. In fact, in 1912 Victor Hess discovered CRs by measuring ionization from a balloon~\\cite{hess}, and in 1938 Pierre Auger proved the existence of extensive air showers (EAS) caused by primary particles with energies above $10^{15}\\,$eV by simultaneously observing the arrival of secondary particles in Geiger counters many meters apart~\\cite{auger_disc}.  After almost 90 years of research, the origin of cosmic rays is still an open question, with a degree of uncertainty increasing with energy~\\cite{crbook}: Only below 100 MeV kinetic energy, where the solar wind shields protons coming from outside the solar system, the sun must give rise to the observed proton flux. Above that energy the CR spectrum exhibits little structure and is approximated by broken power laws $\\propto E^{-\\gamma}$: At the energy $E\\simeq4\\times 10^{15}\\,$eV called the ``knee'', the flux of particles per area, time, solid angle, and energy steepens from a power law index $\\gamma\\simeq2.7$ to one of index $\\simeq3.0$. The bulk of the CRs up to at least that energy is believed to originate within the Milky Way Galaxy, typically by shock acceleration in supernova remnants. The spectrum continues with a further steepening to $\\gamma\\simeq3.3$ at $E\\simeq4\\times 10^{17}\\,$eV, sometimes called the ``second knee''. There are experimental indications that the chemical composition changes from light, mostly protons, at the knee to domination by iron and even heavier nuclei at the second knee~\\cite{kascade}. This is in fact expected in any scenario where acceleration and propagation is due to magnetic fields whose effects only depend on rigidity, the ratio of charge to rest mass, $Z/A$. This is true as long as energy losses and interaction effects, which in general depend on $Z$ and $A$ separately, are small, as is the case in the Galaxy, in contrast to extra-galactic cosmic ray propagation at ultra-high energy. Above the so called ``ankle'' or ``dip'' at $E\\simeq5\\times10^{18}\\,$eV, the spectrum flattens again to a power law of index $\\gamma\\simeq2.8$. This latter feature is often interpreted as a cross over from a steeper Galactic component, which above the ankle cannot be confined by the Galactic magnetic field, to a harder component of extragalactic origin. The dip at $E\\simeq5\\times10^{18}\\,$eV could also be partially due to pair production by extra-galactic protons, especially if the extra-galactic component already starts to dominate below the ankle, for example, around the second-knee~\\cite{bgh}. This latter possibility appears, however, less likely in light of a rather heavy composition up to the ankle suggested by several experiments~\\cite{kascade}. In any case, an eventual cross over to an extra-galactic component is also in line with experimental indications for a chemical composition becoming again lighter above the ankle, although a significant heavy component is not excluded and the inferred chemical composition above $\\sim10^{18}\\,$eV is sensitive to the model of air shower interactions and consequently uncertain presently~\\cite{watson}. In the following we will restrict our discussion on ultra-high energy cosmic rays (UHECRs) above the ankle.  Although statistically meaningful information about the UHECR energy spectrum and arrival direction distribution has been accumulated, no conclusive picture for the nature and distribution of the sources emerges naturally from the data. There is on the one hand the approximate isotropic arrival direction distribution~\\cite{bm} which indicates that we are observing a large number of weak or distant sources. On the other hand, there are also indications which point more towards a small number of local and therefore bright sources, especially at the highest energies: First, the AGASA ground array claims statistically significant multi-plets of events from the same directions within a few degrees~\\cite{teshima1,bm}, although this is controversial~\\cite{fw} and has not been seen so far by the fluorescence experiment HiRes~\\cite{finley}. The spectrum of this clustered component is $\\propto E^{-1.8}$ and thus much harder than the total spectrum~\\cite{teshima1}. Second, nucleons above $\\simeq70\\,$EeV suffer heavy energy losses due to photo-pion production on the cosmic microwave background --- the Greisen-Zatsepin-Kuzmin (GZK) effect~\\cite{gzk} --- which limits the distance to possible sources to less than $\\simeq100\\,$Mpc~\\cite{stecker}. For a uniform source distribution this would predict a ``GZK cutoff'', a drop in the spectrum. However, the existence of this ``cutoff'' is not established yet from the observations~\\cite{bergman} and may even depend on the part of the sky one is looking at: The ``cutoff' could be mitigated in the northern hemisphere where more nearby accelerators related to the local supercluster can be expected. Apart from the SUGAR array which was active from 1968 until 1979 in Australia, all UHECR detectors completed up to the present were situated in the northern hemisphere. Nevertheless the situation is unclear even there: Whereas a cut-off seems consistent with the few events above $10^{20}\\,$eV recorded by the fluorescence detector HiRes~\\cite{hires}, it is not compatible with the 8 events above $10^{20}\\,$eV measured by the AGASA ground array~\\cite{agasa}. It can be remarked, however, that analysis of data based on a single fluorescence telescope, the so-called monocular mode in which most of the HiRes data were obtained, is complicated due to atmospheric conditions varying from event to event~\\cite{cronin}. The solution of this problem may have to await more analysis and, in particular, the completion of the Pierre Auger project~\\cite{auger} which will combine the two complementary detection techniques adopted by the aforementioned experiments and whose southern site is currently in construction in Argentina.  This currently unclear experimental situation could easily be solved if it would be possible to follow the UHECR trajectories backwards to their sources. However, this may be complicated by the possible presence of extragalactic magnetic fields, which would deflect the particles during their travel. Furthermore, since the GZK-energy losses are of stochastic nature, even a detailed knowledge of the extragalactic magnetic fields would not necessarily allow to follow a UHECR trajectory backwards to its source since the energy and therefor the Larmor radius of the particles have changed in an unknown way. Therefore it is not clear if charged particle astronomy with UHECRs is possible in principle or not. And even if possible, it remains unclear to which degree the angular resolution would be limited by magnetic deflection. This topic will be discussed in Sect.~3.  The physics and astrophysics of UHECRs are also intimately linked with the emerging field of neutrino astronomy (for reviews see Refs.~\\cite{nu_review}) as well as with the already established field of $\\gamma-$ray astronomy (for reviews see, e.g., Ref.~\\cite{gammarev}). Indeed, all scenarios of UHECR origin, including the top-down models, are severely constrained by neutrino and $\\gamma-$ray observations and limits. In turn, this linkage has important consequences for theoretical predictions of fluxes of extragalactic neutrinos above about a TeV whose detection is a major goal of next-generation neutrino telescopes: If these neutrinos are produced as secondaries of protons accelerated in astrophysical sources and if these protons are not absorbed in the sources, but rather contribute to the UHECR flux observed, then the energy content in the neutrino flux can not be higher than the one in UHECRs, leading to the so called Waxman-Bahcall bound for transparent sources with soft acceleration spectra~\\cite{wb-bound,mpr}. If one of these assumptions does not apply, such as for acceleration sources with injection spectra harder than $E^{-2}$ and/or opaque to nucleons, or in the top-down scenarios where X particle decays produce much fewer nucleons than $\\gamma-$rays and neutrinos, the Waxman-Bahcall bound does not apply, but the neutrino flux is still constrained by the observed diffuse $\\gamma-$ray flux in the GeV range. This will be discussed in the following section. ",
        "conclusions": "We have reviewed two current issues in theoretical ultra-high energy cosmic ray research.  The first one concerns constraints on scenarios attempting to explain highest energy cosmic rays by extra-galactic sources producing not only cosmic rays but also photons: Improved data analysis and, in the future, improved data for example from GLAST can considerably reduce estimates of the true extra-galactic GeV $\\gamma-$ray background which acts as a calorimeter of electromagnetic energy injected above $\\sim10^{15}/(1+z)\\,$eV. Already current estimates imply that scenarios producing considerably more photons than hadrons, such as extra-galactic top-down scenarios and the Z-burst mechanism, can not explain all of the highest energy cosmic ray flux. A further reduced diffuse GeV $\\gamma-$ray background will start to constrain even normal acceleration scenarios.  As for the second issue we pointed out that the influence of large scale cosmic magnetic fields on ultra-high energy cosmic ray propagation is currently hard to quantify and may not allow to do ``particle astronomy'' along most lines of sight, especially if a significant heavy nucleus component is present above $10^{19}\\,$eV. In this case extensive Monte Carlo simulations including nuclei and based on constrained large scale structure simulations will be necessary to fully exploit data from future instruments such as the Pierre Auger~\\cite{auger} and EUSO projects~\\cite{euso}."
    },
    "0404/astro-ph0404597_arXiv.txt": {
        "abstract": "Central regions of superclusters are the ideal places where to study cluster merging phenomena: in fact the accretion activity is enhanced, as predicted by the cosmological simulations. In this paper I review the case-study of the Shapley Concentration, aimed to understand the effect of major mergings on the intracluster medium and the galaxy population of the involved clusters. ",
        "introduction": "Cluster mergings are known to be among the most energetic phenomena in the Universe, but until now studies at all wavelengths have not been extensively carried on: therefore it is still unclear how the collision energy is dissipated and which is the effect of merging on the emission properties of the galaxies and on the physics of the intracluster medium.  In cosmological N-body simulations the cluster accretion happens along specific directions defined by the density caustics and richer clusters form preferentially where the environment density is higher. Superclusters can be considered as the observational counterparts of these caustics and it is expected that at their centers the cluster accretion is still strongly active. Therefore, superclusters are the ideal places where to study merging phenomena, because the cross-section for cluster collisions is enhanced.  The best place for these studies is the central region of the Shapley Concentration supercluster, a huge concentration of clusters at $z\\sim 0.05$ (\\cite{ray89}, \\cite{plionis91}, \\cite[Raychaudhury et al. 1991]{raychaudhury91}, \\cite[Zucca et al. 1993]{zucca93}). This region is anomalously rich in clusters, considering that it has 25 members while the Great Attractor (with a similar mass overdensity) contains only 6 clusters (see Table 3 of \\cite[Zucca et al. 1993]{zucca93}): for some reasons, in this region the cluster formation efficiency has been enhanced also with respect to similar regions. Moreover, as noted by \\cite{raychaudhury91}, the fraction of clusters with substructures in this supercluster is higher than elsewhere, meaning that the process of cluster formation is still strongly active, suggesting that this region could be considered a ``nursery\" of rich clusters.  \\begin{figure} \\centering \\caption{Central region of the Shapley Concentration supercluster (figure from \\cite[Drinkwater et al. 2004]{drinkwater04}; reproduced by permission of CSIRO Publishing, Melbourne, Australia; copyright Astronomical Society of Australia). Note the two structures (cluster complexes) at $\\alpha=13^h 30^m$ and $12^h 55^m$, connected by a bridge of galaxies resembling the Great Wall. Circles indicate the position of clusters. } \\end{figure} ",
        "conclusions": "Central regions of superclusters give the possibility to find cluster mergings at different phases and strengths and to study the consequences of cluster collisions on the intracluster medium and the galaxy population. The case study of the Shapley Concentration which I presented here shows that the multiwavelength approach is the best way to analyse merging clusters."
    },
    "0404/astro-ph0404242_arXiv.txt": {
        "abstract": "We present observational results from studying the quasi-periodicities in global solar radio flux during periods of enhanced noise storm activity, over durations of $\\sim4$~hrs a day ( \"intra-day\" variations ), observed at 77.5 MHz with the newly commissioned log-periodic array tracking system at the Gauribidanur radio observatory. Positional information on the storm centers were obtained with the radio imaging data from the Nan\\c cay RadioHeliograph ( NRH ), while their active region ( AR ) counterparts on the photosphere ( and the overlying chromosphere ) were located from the $H{\\alpha}$ images of the Big Bear Solar Observatory ( BBSO ). The quasi-periodicity in flux was found to be 110 minutes, with the fluctuation in flux being $3 (\\pm 1.5)$ solar flux units ( sfu ). The results of such pulsations are interpreted qualitatively as evidence for coronal seismology. ",
        "introduction": "Type I solar noise storms, ever since their discovery in the year 1946 \\cite{hey46}, have proved to be amongst the most profilic of events to occur at metric wavelengths. They comprise of the bursts that are radio flux enhanced, narrowband ( df/f $\\approxeq$ 3 \\%  ), spiked ( 0.1 - 1 s ) events, with the broadband ( df/f $\\approxeq$~100\\% ) continuum, lasting from several hours to a few days, serving as their diffuse background radio emission, on the dynamic spectral records. According to \\cite{hey46},~\\cite{mcd47}, noise storms have their origins in the outer corona, and appear proximal to the sites of active regions in the photosphere and the chromosphere.\\\\ \\\\ In this paper, we present results from investigations on the variations of global radio flux, during periods of prolonged noise storm activity, with attributes to quasi-periodicities on short time scales, from observing the Sun continuously for about 6 hrs each day. Details on the time-delay control enabled, broadband operable antenna array system used for the study are provided, followed by a note on the observation schedule adopted. Interpretation of the imaging data from complementary observations on the noise storms and the associated underlying ARs constitute part of the latter section. The scheme deployed to analyse and determine the extent of global radio flux variations is described in \\S 4. The results of this study are presented in the Discussion \\& Conclusion section, along with qualitative remarks on the plausible origin for the observed flux oscillations at metric wavelengths, and their likely coronal implications. ",
        "conclusions": "The general consensus on the radiation mechanism of noise storms is one of plasma emission. New emergence of magnetic flux, and the changes in the topography of the magnetic field arising out of the reconnection~\\cite{bewen81} of ruptured, preexisting flux lines with the newly sprouting flux, about the sites of bipolar sunspot activity, leads to the formation of recurrent shock waves~\\cite{spbeh81},~\\cite{wen81}. The electrons of the ambient plasma within the reconnection loop, at the outer coronal layers, become unstable as a result of this realignment of magnetic lines of force, and generate the Langmuir (L) waves or the Upper hybrid (UH) waves. Such L-waves and UH-waves coalesce in turn with the ion acoustic waves, at the sites of density inhomogeneities, to emit the type I continuum noise storm radiation. On the other hand, bursts of type I are produced when the UH-waves, produced by the electrons trapped in magnetic flux loops, are scattered on the Lower hybrid (LH) waves generated at their shock wavefront.\\\\ \\\\ The determination of quasi-periodicity is subject to a lower cut-off, defined by the beam dwell-time. The adjacent beam positions are $9^{o}$ apart, and the dwell-time varies as a function of the apparent declination of the Sun, in a manner defined by Equation~(\\ref{spe1}), where $t_{e}$ is 36 minutes, and ($\\delta$) varied from $13.^{o}2$~N to $23^{o}$~N. Hence, $t_{d}$ is the minimum value for quasi-periodicity that can be realised by the tracking system in its current scheme of observation. The quasi-periodic distribution, for the 32 cases of absolute deviation from mean, is shown in the histogram of Figure~\\ref{spf6}. The peak in the distribution occurs at 110 minutes, which also is tantamount to the periodicity in global solar radio flux variation at 77.5 MHz. In addition, more than half ($\\approx 60 \\% $)~of the distribution occurs in the range of periodicities from 100-150 minutes.\\\\ \\\\ Regarding the spline interpolation method employed for determination of the absolute deviation from mean ( shown in the four sub-plots of Figure~\\ref{spf5} ), it needs to be duly emphasised that vast deviations, occuring on either side of the mean and at consecutive beam position observations, are bound to yield quasi half-periodicity values lesser than the minimum achieveable \"cadence\" between the two adjacent observations on the same day - a case of so-called \"super-resolution\", that needs to be cautiously approached, especially when the distribution were to peak at those values. In Figure~\\ref{spf6}, the first couple of values in the vicinity of 50 min are the case in point, corresponding to the first negative going half-waveforms on the upper-right and the lower-left sub-plots of Figure~\\ref{spf5}; they have values for quasi-periodicity less than the $t_{d}$ defined by Equation~(\\ref{spe1}) for the particular days.\\\\ \\\\ Quasi-periodic fluctuations in solar coronal emission, observed at the operating frequencies of the antennas on-board the Prognoz 1 high-apogee satellites~\\cite{gps77} equipped with kilometric radiation detectors, have estimated periods ranging from 6 sec to 2 hours. A specific quasi-periodicity value of 118 ($ \\pm 20$)~minutes has been quoted among many other values detected, for fluctuations in radio emission by 10-15 dB from the mean solar flux. In the present study done at 77.5 MHz, the quasi-periodic fluctuations have an absolute deviation from the intra-day mean ranging from 1.5 to 4 sfu.\\\\ \\\\ The absolute deviation from mean flux, for each of the 17 days of enhanced activity, is taken as a measure of the intra-day quasi-periodicity in solar radio flux, and found to be 110 minutes, with the fluctuations in flux being $3 (\\pm 1.5)$ sfu. Positional information from the Nan\\c cay RadioHeliograph data, and features of the causative ARs of the underlying photospheric disk from the full disk $H{\\alpha}$ images of the BBSO, along with the radio spectral data published in the SGD reports, lead to the conclusion that heightened flux emission, with global ramifications, are a result of type I noise storms.\\\\ \\\\ The occurrence of the noise storm source regions, and their contribution to enhanced radio emission on a global scale, has been corroborated from complementary evidence based on active region data obtained from the BBSO full disk $H{\\alpha}$ images, the NRH's 164 MHz imaging data on noise storms and the metric noise storm spectral observations as reported in the SGD reports. Quasi-periodic pulsations in global solar radio flux were observed, and their origin has been attributed to modulation of the plasma radiation by magnetohydrodynamic ( MHD ) disturbances in the corona, at sites above large ARs threaded by complex magnetic flux tubes. MHD disturbances are generated by the weak shocks associated with magnetic reconnection events, as MHD oscillations or resonance of MHD waves, and fluctuations ( pulsations ) in the plasma radio emission ensue at the source region for noise storms. MHD waves on all scales, ranging in wavelength from the coronal loop-size (fraction of $R_{\\odot}$) down to the gyroradii (a few meters) of coronal ions, are believed to play a key role in the transport of mechanical energy, from the denser regions of the chromosphere to the Sun's corona and further as a steady stream called the solar wind; by means of dissipation of the wave energy, the corona is heated and sustained at elevated temperatures. The emissivity of trapped epithermal particles in coronal magnetic arches is modulated by a propagating MHD wave~\\cite{asc87}. In the case of noise storms, the plasma waves get converted to transverse electromagnetic ( TEM ) waves at sites of trapped density inhomogeneities and magnetic reconnections. Theoretical interpretation of the oscillatory phenomena, observed in the outer solar corona, based on the MHD wave theory, along with supportive information, regarding the plasma parameters, from investigation of high spatial and temporal imaging data on the coronal magnetic structures above the active region complexes, would significantly demystify the underlying mechanisms involved in the plasma dynamics of the outer corona, coronal-seismology and coronal-heating."
    },
    "0404/astro-ph0404205_arXiv.txt": {
        "abstract": "\\singleandabitspaced  Two ground-based experiments have recently independently detected TeV $\\gamma$-rays from the direction of the Galactic center. The observations made by the VERITAS and CANGAROO collaborations are unexpected, although not impossible to interpret in terms of astrophysical sources. Here we examine in detail whether the observed $\\gamma$-rays may arise from the more exotic alternative of annihilations of dark matter particles clustered in the center of the Galaxy. ",
        "introduction": "Recently, the VERITAS \\cite{Kosack:2004ri} and CANGAROO \\cite{Tsuchiya:2004wv} collaborations, using the Whipple 10 meter and CANGAROO-II Atmospheric \\u{C}erenkov Telescopes (ACTs), respectively, have made significant detections of TeV $\\gamma$-rays from the Galactic center region. Although the origin of this emission is not yet known, there are several possible, although unlikely, astrophysical sources in the field of view. Alternatively, this may be a signature of annihilating dark matter particles.  The Galactic center is a complex and rich region. Its most notable inhabitant is a $2.6 \\times 10^{6} M_{\\odot}$ black hole, coincident with the radio source Sgr A$^*$, which also demonstrates variable emission at infra-red \\cite{Ghez:2003hb}, soft X-ray \\cite{xray} and hard X-ray \\cite{Belanger:2003se} wavelengths. Additionally, the region may contain massive X-ray binaries emitting relativistic plasma jets (microquasars) capable of producing high-energy $\\gamma$-rays \\cite{microquasar_gamma} by either hadronic ($\\pi^0$ production) \\cite{microquasar_proton} or leptonic (inverse Compton) \\cite{microquasar_lepton} processes. The region could also contain Supernova Remnants (SNRs) which are widely believed to be the source of Galactic cosmic rays. TeV $\\gamma$-rays have indeed been observed from several nearby supernova remnants such as the Crab \\cite{snr_crab} and Cas A \\cite{snr_casa}. The responsible mechanism remains unclear, but again could be either leptonic or hadronic (see Ref.~\\cite{Volk_snr} for a discussion of TeV $\\gamma$-ray production in SNRs). The SNR Sgr A East lies only a few parsecs away from Sgr A$^*$, but is not a likely TeV source in itself \\cite{SgrA_East_SNR}. Interactions of its expanding shell with molecular clouds in its environment could, in principle, produce high-energy $\\gamma$-rays \\cite{Aharonian94_SNRMOLECULAR}.  Alternatively, strong winds from massive O and B type stars could lead to hadronic interactions that may result in the emission of high-energy $\\gamma$-rays \\cite{TeVOBObservations,TeVOBTheory}. Two massive, compact, young star clusters which  contain such stars (Arches and Quintuplet) are located roughly 10 arcminutes away from the Galactic center. Chandra observations of the Arches cluster have revealed non-thermal emission attributed to relativistic electrons accelerated in colliding wind shocks from binary systems within the cluster or in the winds from single stars with the collective winds from the other stars in the cluster \\cite{YusefArches1,YusefArches2}. It is argued that the existence of non-thermal particles could result in X-ray/$\\gamma$-ray emission by inverse Compton scattering of ambient photons.  Observations by INTEGRAL \\cite{INTEGRAL_SOURCES} and EGRET \\cite{dingus} have revealed $\\gamma$-ray emission from the Galactic center region, although thus far no corresponding sources have been identified. X-ray surveys of the region have revealed a new population of discrete sources, many of which resemble X-ray binaries (see Ref.~\\cite{Galactic_XRAY} and references therein). For more information on the Galactic center region, see Refs.~\\cite{Galactic_XRAY,YusefZadeh00,LaRosa00}.    Despite the extensive body of evidence in favor of cold, non-baryonic, dark matter \\cite{cdm}, its identity remains elusive. Searches for particle dark matter have been carried out using a variety of methods. Direct searches attempt to observe the recoil energy as Galactic dark matter particles orbiting through the Solar system scatter elastically off nuclear targets \\cite{direct}. Indirect searches attempt to observe the products of dark matter annihilations such as neutrinos \\cite{indirectneutrino}, positrons \\cite{positrons}, anti-protons \\cite{antiprotons} and, in particular, high-energy $\\gamma$-rays \\cite{indirectgamma,Bergstrom:1997fj,Bergstrom:2001jj}.   The remainder of this article is organized as follows. In section 2, we summarize the observations of high-energy $\\gamma$-ray emission from the Galactic Center region. In sections 3, 4 and 5, we discuss the spectral and spatial features of these observations and assess whether annihilating dark matter could be responsible for these observations. In section 6, we consider particle dark matter candidates suggested by new physics beyond the Standard Model, in particular supersymmetry, in the light of these recent observations. We present our conclusions in section 7. In Refs.~\\cite{Kosack:2004ri} and \\cite{Tsuchiya:2004wv}, this possibility was briefly discussed. Our intention here is to explore this scenario in considerably more detail. ",
        "conclusions": "In this article, we have discussed the possibility that annihilating dark matter in the Galactic center has produced the flux of $\\gamma$-rays observed by the CANGAROO and VERITAS collaborations. Although it is possible that these $\\gamma$-rays are the result of astrophysical processes, we summarize here the characteristics required of a dark matter particle, if its annihilations are responsible for the observed $\\gamma$-ray emission.  \\subsection{The Spectrum}  It is difficult to reconcile the spectra observed by the CANGAROO-II and Whipple experiments. The spectrum measured by CANGAROO-II is consistent with an annihilating particle of mass in the range of 1--3 TeV. On the other hand, Whipple has observed a substantial flux above its rather high threshold of 2.8 TeV, requiring a much heavier dark matter particle. Future observations will be needed to conclusively determine the spectrum of $\\gamma$-rays from the Galactic center in the GeV-TeV range.   \\subsection{The Halo Profile and Annihilation Cross Section}  For annihilating dark matter to produce the flux measured by either the CANGAROO or VERITAS collaborations, very high annihilation rates are required. This, in turn, requires a very large annihilation cross-section {\\it and} a very concentrated dark matter distribution in the innermost region of our Galaxy. Even if we consider a particle with a rather large annihilation cross-section, say $\\sim 10^{-26} \\, \\rm{cm}^3/\\rm{s}$, extremely cusped halo models, such as a Moore {\\it et al.} with adiabatic compression would be required. Alternatively, halo profiles with a density spike most plausibly associated with the central SMBH  or a nearby VMBH could provide the observed flux.   \\subsection{Future Prospects}  With the current data, it is very difficult to determine whether the $\\gamma$-rays observed from the Galactic center region by CANGAROO-II and Whipple are the product of dark matter annihilations rather than other, less exotic, astrophysics. This state of affairs may change with improved data in the future. As the angular resolution of ACTs (as well as space-based $\\gamma$-ray experiments, such as GLAST) improves, it will become clear whether the observed TeV emission comes from our Galaxy's dynamical center rather than from nearby star clusters, X-ray binaries or other objects. This information will be crucial for confidently identifying TeV emission as the product of dark matter annihilations.  Moreover, as the $\\gamma$-ray spectrum in the GeV to multi-TeV range becomes more refined, it may become possible to ascertain whether the observed emission is the result of annihilating dark matter. In particular, evidence of line emission would provide a ``smoking gun'' signal for anihilations. Presently, we are eagerly awaiting results from the HESS collaboration, which has also been observing the Galactic center. HESS should be more sensitive in the direction of the Galactic center than either CANGAROO-II or Whipple. Also, with four telescopes, HESS's angular resolution should be superior to single-telescope ACTs."
    },
    "0404/astro-ph0404343_arXiv.txt": {
        "abstract": "We present $R$-band photometry of the X-ray transient and candidate black hole binary XTE J1650-500 obtained between 2003 May  and August with the 6.5m Clay Telescope.  A timing analysis  of these data reveals a photometric period of $0.3205\\pm 0.0007$ days (i.e.\\ 7.63 hr) with a possible alias at 0.3785 days (9.12 hr).  Our photometry completely rules out the previously published spectroscopic period of 0.212 days (5.09 hr).  Consequently, we reanalyzed the 15 archival ESO/VLT spectra (obtained 2002 June by Sanchez-Fernandez et al.) that were the basis of the previously published spectroscopic period.  We used a ``restframe search'' technique that is well suited for cases when the signal-to-noise ratio of individual spectra is low. For each of roughly 1.1 million binary ephemerides, we summed all of the spectra in a trial restframe of the secondary star, and each restframe spectrum was cross-correlated against a template spectrum. We then searched for the set of orbital parameters that produced the strongest cross-correlation value. The results confirmed the photometric period of 0.3205 days, and rule out the alias period near 0.38 days. The best value for the velocity semiamplitude of the companion star is $K_2 = 435 \\pm 30$ km s$^{-1}$, and the corresponding optical mass function is $f(M)=2.73 \\pm 0.56\\,M_{\\odot}$. The spectral type of the companion star is not well constrained because we only have six template spectra available to us.  The K4V template provides the best match; next best matches are provided by the G5V and K2III templates. We also find that the accretion disk dominates the light in the $R$-band where the disk fraction is 80\\% or higher, although this value should be treated with caution owing to the poor signal-noise-ratio and the limited number of templates. The amplitude of the phased $R$-band light curve is 0.2 magnitudes, which gives a lower limit to the inclination of $50\\pm 3^{\\circ}$ in the limiting case of no contribution to the $R$-band light curve from the accretion disk. If the mass ratio of XTE J1650-500 is similar to the mass ratios of other black hole binaries like A0620-00 or GRS 1124-683 (e.g.\\ $Q\\gtrsim 10$), then  our lower limit to the inclination  gives an upper limit to the mass of the black hole in XTE J1650-500  of $M_1 \\lesssim 7.3\\,M_{\\odot}$. However, the mass can be considerably lower if the $R$-band flux is dominated by the accretion disk.  For example, if the accretion disk does contribute 80\\% of the flux, as our preliminary results suggest, then the black hole mass would be only about $4\\,M_{\\odot}$. ",
        "introduction": "X-ray novae (XN) provide the strongest evidence for the existence of stellar mass black holes.  XN are interacting binaries where a neutron star or a black hole accretes matter from a companion (usually a dwarf-like K or M star). The accretion rate on to the black hole increases substantially during the outburst phase, and hence the X-ray luminosity can vary by large amounts  (e.g.\\ factors of $10^6$ or more).  The majority of XN spend most of their time in a ``quiescent'' state where the X-ray luminosity is on the order of the optical luminosity of the companion star.  It is in quiescence  where the companion star can be best studied in the optical, where the observed radial velocity and light curves of the companion lead to dynamical mass measurements for the compact primary. There are 18 X-ray binaries (15 of them XN) where the mass of the primary has been shown to exceed the maximum mass of a stable neutron star ($\\approx 3\\,M_{\\odot}$, Kalogera \\& Baym 1996), confirming the presence of black holes in these systems: GRO J0422+32 (Orosz \\& Bailyn 1995, Filippenko, Matheson, \\& Ho 1995); A0620-00 (McClintock \\& Remillard 1986); GRS 1009-45 (Filippenko et al.\\ 1999); XTE J1118+480 (McClintock et al.\\ 2001, Wagner et al.\\ 2001); GS 1124-683 (Remillard, McClintock, \\& Bailyn 1992); 4U 1543-47 (Orosz et al.\\ 1998); XTE J1550-564 (Orosz et al.\\ 2002); GRO J1655-40 (Bailyn et al.\\ 1995); GX 339-4 (Hynes et al.\\ 2003); H1705-250  (Remillard et al.\\ 1996); SAX J1819.3-2525 (Orosz et al.\\ 2001); XTE J1859+226 (Filippenko \\& Chornock 2001); GRS 1915+105 (Greiner, Cuby, \\& McCaughrean 2001); GS 2000+25 (Casares, Charles, \\& Marsh 1995); GS 2023+338 (Casares, Charles, \\& Naylor 1992); Cyg X-1 (Gies \\& Bolton 1986); LMC X-3 (Cowley et al.\\ 1983); LMC X-1 (Hutchings et al.\\ 1987). These sources open up the possibility of studying general relativity in the strong field regime.  For example, the study of high frequency quasiperiodic oscillations in the X-ray light curves of certain XN may lead to a measurement of black hole spin (e.g.\\ McClintock \\& Remillard 2004 and cited references).  We must press hard to obtain further observations of black hole masses in order to fully pursue these opportunities.   XTE J1650-500 (hereafter J1650) was discovered by RXTE on 2001 September 5 (Remillard 2001) and subsequently reached a peak X-ray intensity of 0.5 Crab. Based on subsequent observations, J1650 was established as a strong black hole candidate based on its X-ray spectrum and variability in the X-ray light curve (Markwardt, Swank, \\& Smith 2001; Revnivtsev \\& Sunyaev 2001; Wijnands, Miller, \\& Lewin 2001). The radio counterpart was discovered with the Australia Telescope Compact Array (ATCA) by Groot et al.\\ (2001). Further radio observations sampled the behavior of XTE J1650-500 along all its X-ray states (Corbel et al.\\ in preparation). We highlight two key results obtained during its outburst phase. First, RXTE observations during the third and fourth weeks of the outburst yielded a strong X-ray QPO with an rms amplitude of $5.0 \\pm 0.4\\%$ at a frequency of $\\nu=250\\pm 5$ Hz (Homan et al.\\ 2003b). Second, {\\em XMM-Newton} observed a broad, skewed emission line due to Fe K$\\alpha$ (Miller et al.\\ 2002). Those authors argue that their results imply the primary is a nearly maximal Kerr black hole that is delivering its spin-down energy to the accretion flow.   The first significant observational program in the optical was that of Sanchez-Fernandez et al.\\ (2002, hereafter SF2002). They observed J1650 on the night of 2002 June 10 with the fourth 8.2m telescope at the European Southern Observatory, Paranal, and reported the following orbital elements: an orbital period of $P=0.212\\pm 0.001$ days and a velocity semiamplitude of $K_2=309\\pm 4$ km s$^{-1}$, resulting in an optical mass function of $f(M)=0.64\\pm 0.03\\,M_{\\odot}$. The results in this paper contradict these findings.  In this paper we report the results of our photometric study of J1650.  A time series analysis of our photometry rules out the orbital period  reported by SF2002. Consequently we also report herein our reanalysis of the SF2002 data obtained from the ESO archives. We show that the spectroscopic period we derive from these data is consistent with our photometric period  $P=0.3205$ days, and we go on to determine the orbital elements of the system.  We outline below our observations and reductions, and our analysis techniques.  We end with a brief discussion of the implications of our results regarding the fast QPO  observed for J1650. ",
        "conclusions": "Is general relativity (GR) the correct theory of gravity in the strong fields found near a black hole?  One promising approach to answering this fundamental question is offered by the key discovery of the {\\em Rossi X-ray Timing Explorer} (RXTE) that seven Galactic black holes (including J1650) display quasiperiodic (QPO) X-ray oscillations in the range of $100-450$ Hz (McClintock \\& Remillard 2004 and cited references). These fast QPOs must be produced near the event horizon since the X-rays originate there and since such high frequencies are comparable to the Kepler frequency of the innermost stable circular orbit around a black hole ($\\nu_K=2199-16150\\,(M/M_{\\odot})^{-1}\\,{\\rm Hz}$ depending on the dimensionless spin parameter $a_*$, where the full range of $a_*$ is 0 to 1 e.g.\\ Kato, Fukue, \\& Mineshige 1998).  RXTE has made a further important discovery: four of these seven black holes produce pairs of stationary, fast QPOs that have frequencies in a 3:2 ratio (McClintock \\& Remillard 2004; Homan et al.\\ 2003a). Such commensurate frequencies are a hallmark of non-linear resonance phenomena.  Thus, this discovery has promoted a ``resonance model'' that invokes enhanced emissivity at the radius in the accretion disk where two of the three spatial coordinate frequencies (e.g.\\ Keplerian and radial) are commensurate (Abramowicz \\& Kluzniak 2001; Abramowicz et al.\\ 2003).  For the three black holes with the QPO pairs {\\em and} measured masses (e.g.\\ GRO J1655-40, XTE J1550-564, and GRS 1915+105), the QPO frequency is correlated with the black hole mass, where $\\nu\\propto M^{-1}$ (e.g.\\ McClintock \\& Remillard 2004). This scaling is expected for GR oscillations, but different sources can lie on the same curve only if they have similar values of the dimensionless spin parameter $a_*$.  Schnittman \\& Bertschinger (2004) performed ray tracing calculations in the Kerr metric for emitting blobs orbiting a black hole at a radius where there is a 3:1 resonance between the azimuthal ($\\Omega$) and radial coordinate frequencies. It was shown that general relativistic effects may impart a QPO at $\\Omega$ and a beat oscillation at $\\slantfrac{2}{3}\\, \\Omega$ with a relative strength that depends on the angular width of the emitting blob of material.  This interpretation, when combined with accurate mass measurements, yields values for the dimensionless spin parameter: $0.40 \\le a_*\\le 0.55$ for GRO J1655-40 and $0.3 \\le a_* \\le 0.6$ for XTE J1550-564 (Remillard et al.\\ 2002).  This illustrates the potential diagnostic power of fast QPOs, if we can specify the correct oscillation mechanism.  Furthermore, measurements of black hole mass and QPO frequency allow {\\em model-independent} comparisons of spin differences between black hole subclasses as distinguished, for example, by relativistic jets or binary period.  With our measurement of the orbital period and optical mass function of J1650, we have taken the first steps need to fully exploit the potential of this system.  If the QPO at $\\nu=250\\pm 5$ Hz observed by Homan et al.\\ (2003b) represents the $2\\nu_0$ oscillation, then the predicted mass of the black hole would be about $7.5\\,M_{\\odot}$, if J1650 lies on the same $\\nu:M$ curve on which GRO J1655-40, XTE J1550-564, and GRS 1915+105 lie (i.e.\\ if the black holes all have similar values of the dimensionless spin parameter $a_*$). Although there are substantial uncertainties in both our measured value for the optical mass function and in our inclination estimate, it appears that the black hole mass might be less than $7.5\\,M_{\\odot}$, which in the context of the orbital resonance model would imply that the value $a_*$ in J1650 is much smaller than it is for the other three sources (e.g.\\ Abramowicz et al.\\ 2004). On the other hand, on the basis of X-ray spectroscopy, Miller et al.\\ (2002) have argued that J1650 contains a maximal Kerr black hole.  Thus J1650 may be a challenging case for the orbital resonance model of the high frequency QPOs, although we again point out that our mass estimate has large uncertainties.  Recent models for black hole formation (e.g.\\ Fryer \\& Kalogera 2001) predict a continuous and roughly exponential distribution of masses for black holes in binary systems, where the number of black holes formed falls off as the mass increases.  Depending on their model assumptions, these authors predict the number of black holes with masses in the range  $3-5\\,M_{\\odot}$ should be roughly two to three times the number of black holes with masses in the range $8-10\\,M_{\\odot}$.  Bailyn et al.\\ (1998) performed a statistical analysis of seven black holes with dynamical mass estimates and found strong evidence for a peak in the mass distribution centered near $7\\,M_{\\odot}$ and a ``gap'' in the distribution between the neutron star masses and the peak near $7\\,M_{\\odot}$. Since that time, the sample of objects has more than doubled, and it now appears that the evidence for a peak in the distribution near $7 \\, M_{\\odot}$ is weak (see the references cited in \\S1 and Orosz 2003). However, black holes with masses in the $3-5\\,M_{\\odot}$ range still seem to be rare, with GRO J0422+32 being the only  good candidate (Gelino \\& Harrison 2003). In this regard, J1650 is an interesting system because of its probable low mass.  Clearly additional data will be required to make more definitive statements about the mass of the compact object in J1650. It should not be unduly difficult to obtain higher quality spectra. In hindsight, we now know that the companion star is a  K-star, and as such it should have relatively strong absorption features near 5170\\AA.  This region was not included in the VLT spectra owing to the relatively high resolving power of the 1200R grism. It should be easier to measure radial velocities in spectra with more wavelength coverage (at the expense of resolving power).  Also, spectral observations obtained in good to excellent seeing conditions ($\\lesssim 0.6$ arcseconds) obviously will have better signal-to-noise than the current spectra do (the seeing varied between 1 and about 1.4 arcseconds).  Additional photometry should be obtained, and it would be helpful if two or more bandpasses could be used, since the amount of contaminating light from the disk should vary with color."
    },
    "0404/astro-ph0404496_arXiv.txt": {
        "abstract": "The physics of the ``hot spots\" on stellar surfaces and the associated variability of accreting magnetized rotating stars is investigated for the first time using fully three-dimensional magnetohydrodynamic simulations. The magnetic moment of the star $\\rvecmu$ is inclined relative to its rotation axis $\\bf{\\Omega}$ by an angle $\\Theta$ (we will call this angle the ``misalignment angle\") while the disk's rotation axis is parallel to $\\bf{\\Omega}$. A sequence of misalignment angles was investigated, between $\\Theta=0^\\circ$ and $90^\\circ$. Typically at small $\\Theta$ the spots are observed to have the shape of a bow which is curved around the magnetic axis, while at largest $\\Theta$ the spots have a shape of a bar, crossing the magnetic pole. The physical parameters (density, velocity, temperature, matter and energy fluxes, etc.) increase toward the central regions of the spots, thus the size of the spots is different at different values of these parameters. At relatively low density and temperature, the spots occupy approximately $10-20\\%$ of the stellar surface, while at the highest values of these parameters this area may be less than $1\\%$  of the area of the star. The size of the spots increases with the accretion rate. The  light curves were calculated for different $\\Theta$ and inclination angles of the disk $i$. They show a range of variability patterns, including one maximum-per-period curves (at most of angles $\\Theta$ and $i$), and two maximum-per-period curves (at large $\\Theta$ and $i$). At small $\\Theta$, the funnel streams may rotate faster/slower than the star, and this may lead to quasi-periodic variability of the star. The results are of interest for understanding the variability and quasi-variability of Classical T Tauri Stars, millisecond pulsars and cataclysmic variables. ",
        "introduction": "In accreting magnetized stars the inflowing matter is channelled to regions near the magnetic poles of the star, forming ``hot spots\" on the star's surface (e.g., Ghosh \\& Lamb 1979; Uchida \\& Shibata 1985; Camenzind 1990; K\\\"onigl 1991; Shu et al. 1994; Hartmann et al. 1994). The hot  spots form under conditions in which the star's magnetic field is strong enough to form a magnetosphere. Examples include classical T Tauri stars (e.g., Herbst et al. 1986; Bouvier \\& Bertout 1989; Johns \\& Basri 1995;  1998;  Petrov, et al. 2001a,b; Alencar \\& Batalha 2002), cataclysmic variables (e.g., Wickramasinghe, Wu, \\& Ferrario 1991; Livio \\& Pringle 1992; Warner 1995; Warner 2000), X-ray pulsars (e.g., Ghosh \\& Lamb 1979; Tr\\\"umper et al. 1985; Bildsten et al. 1997) and millisecond pulsars (e.g., Chakrabarty, et al. 2003). These objects have different dimensions and magnetic field strengths, but the underlying physics of the magnetospheric accretion and hot spots' properties are expected to be similar.  \\begin{figure*}[t] \\epsscale{2.} \\plotone{f1.eps} \\caption{Matter flow near the star and hot spots for the case of a relatively warm disk ($T_d=0.03$). The square panels show the surfaces of the density in the magnetospheric flows (green color corresponds to $\\rho=0.3-0.4$, while yellow-green color to $\\rho=0.6$). Red lines are sample magnetic field lines; the dark-green lines show the sample streamlines of matter flow. The circular panels show density distribution in the hot spots for different misalignment angles $\\Theta$. For $\\Theta=15^\\circ, 30^\\circ$ and $45^\\circ$, the magnetic axis is directed towards the observer, while for $\\Theta=60^\\circ, 75^\\circ$ and $90^\\circ$, it is directed at an angle $45^\\circ$ relative to the observer. It is shown by the crosses in the circular panels. The direction of the rotation axis is shown with the open circles.} \\label{Figure 1} \\end{figure*}  The observed variability/quasi-variability of these objects can be due to a number of possible phenomena in the magnetosphere (see, e.g., Bouvier 2003; Petrov 2003). One of the  time-scales may be associated with the hot spots. In spite of the large volume of the observational data on variability of these stars, relatively little is known about the properties of the hot spots.   For example, in classical T Tauri stars  (CTTS), hot spots radiation is associated with veiling of the continuum radiation and was analyzed for a number of CTTS. Herbst \\& Koret (1988) and Bouvier et al. (1986, 1993) estimated parameters of the hot spots assuming a model with black body radiation, while Lamzin (1995), Calvet \\& Gullbring (1998), Gullbring et al. (2000), and Ardila \\& Basri (2000) used models of radiation from a shock wave. They concluded that the area of the star covered with the hot spots in different CTTS varies from $0.3\\%$ to $20\\%$ depending on the accretion rate and other factors. In all models, however, the hot spots were considered to be homogeneous, that is, of the same density and temperature. Further, some of the models did not take into account the effect of limb-darkening. Additionally, the shape and location of the hot spots were not known because they can be derived only from a three-dimensional analysis. Similar problems and questions appear during analysis of X-ray pulsars and cataclysmic variables. Thus, it is important to study the properties of the hot spots in greater detail. For analysis of the hot spots we use results of our previous three-dimensional simulations (Romanova et al. 2003, hereafter - R03), and also perform new simulations at a lower temperature in the disk and at a variety of parameters of the star and the accretion flow.  The main questions which can be answered are: (1) What are the shapes of the hot spots? (2) What is the distribution of physical parameters (density, temperature) in the spots? (3) What is the area covered by the spots at different physical parameters? (4) What are the observed light curves at different misalignment angles $\\Theta$ and different inclination angles of the disk $i$?   In \\S 2 of the paper we  describe the underlying model and dimensional examples for CTTS and millisecond pulsars. In \\S 3 we discuss the expected  physical properties of the hot spots. In \\S 4 we calculate the intensity of radiation from the hot spots and the associated light curves. In \\S 5 we discuss dependence of results on different parameters, limitations of the model and future work. \\S 6 gives the summary of our work.     \\begin{figure*}[t] \\epsscale{2.} \\plotone{f2.eps} \\caption{Same as Figure 1 but for cooler disk ($T_d=0.01$).} \\label{Figure 2} \\end{figure*} ",
        "conclusions": "In this paper we fixed parameters of the star and the flow. Below we discuss dependence of the results on $\\Omega_*$ and on the accretion rate $\\dot M$. In \\S 5.2 we discuss the  limitations of the model and future work.   \\subsection{Dependence of Results on Parameters of the Star and the Disk}  \\noindent{\\it $\\Omega$--Dependence:} For the results given here we have taken a relatively low angular velocity, $\\Omega_*=0.19$. This corresponds to CTTS with a period $P=9.4~{\\rm days}$. We did additional simulations for an intermediate angular velocity, $\\Omega_*=0.35$ ($P=3.3~{\\rm days}$), and a rapidly rotating star, $\\Omega_*=1.0$ ($P=1.8~{\\rm days}$). We observed that the hot spot shapes are similar for slowly and for rapidly rotating stars. For the intermediate angular velocity, the bow shape is less prominent. The bow shape of the hot spots reflects the typical shape of the funnel streams which have a thickness that is small compared with their width. This is typical for relatively small inclination angles, $\\Theta\\lesssim 30^\\circ$, where the stream should ``climb\" an appreciable distance above the equatorial plane before accreting to the vicinity of the magnetic pole. At small inclination angles, the bow shape is natural, because it may be considered as a part of the narrow cylindrical ring which occurs for $\\Theta=0$.   \\begin{figure*}[t] \\epsscale{1.7} \\plotone{f13.eps} \\caption {The fraction of the star covered by the hot spots at the density levels: $\\rho=0.4$, $\\rho=1$, and $\\rho=2$ as a function of the accretion rate $\\dot M$. Squares show separate runs at different $\\alpha$ parameters of the disk viscosity.} \\label{Figure 13} \\end{figure*}   \\smallskip \\noindent{\\it {$\\dot M$--Dependence}}: We studied the dependence of the size of the spots on the accretion rate $\\dot M$. We introduced an alpha-type viscosity analogous to one used in R02 and performed a set of runs at different values of accretion rate $\\dot M$ (different $\\alpha$). Figure 13 shows that the fraction ${\\it f}$ of the star covered by the hot spots increases with the accretion rate for a range of density levels, $\\rho=0.4, 1, 2$. Thus, the size of the spots increases as the accretion rate increases. This is in accord with recent results of Ardila and Basri (2000), who analyzed the UV variability of the CTTS BP Tau and have shown significant correlation between the accretion rate and the filling factor of the shocks. In R03 we noticed that at a larger accretion rate, the streams become wider and cover a larger area of the magnetosphere. This may possibly explain the observed variation of the shapes of magnetospheric lines with the accretion rate (Muzerolle et al. 2001).    \\subsection{Limitations of the Model and Future Work}  Variability in different spectral lines in the accreting magnetized stars may be associated with different regions: inner regions of the disk, magnetospheric streams,  hot spots, or outflows. In this paper we analyze only the variability due to the rotation of the hot spots,  while the variability associated with  {\\it other regions} will be investigated  in the future work. The light from the hot spots may be also {\\it obscured} by magnetospheric streams  or by the warped inner regions of the disk if the star is approximately edge-on (e.g., Bouvier, et al. 1999, 2003). This paper does not include the effect of obscuration. The warping of the inner regions of the disk (Aly 1980, Lipunov \\& Shakura 1980, Lai 1999) was not observed in the simulations, but a special set of simulations will be done for investigation of this possible phenomenon.   The magnetic field of the CTTS or millisecond pulsars may not be {\\it a pure dipole} field (Safier 1998; Smirnov et al. 2003). This can lead to more complicated geometry of the hot spots and more complicated variability patterns. However, the detailed analysis of the photometric and spectral variability of a number of CCTS  have shown that most of the observed features can be explained by models with a dipole magnetic field (e.g., Muzerolle, Hartmann, \\& Calvet 1998; Petrov, et al. 2001a,b; Alencar, Johns-Krull, \\& Basri 2001; Bouvier et al. 2003). This paper considers only the case where the star's field is a dipole field. Non-dipolar field geometries will be investigated separately.  We observed that the hot spots may rotate more rapidly or more slowly than the star, which will lead to quasi-periodic oscillations in the light curves and may explain QPOs observed in CTTS (see Smith, Bonnell, \\& Lewis 1995) and millisecond pulsars (see, e.g., Chakrabarty et al. 2003). QPO variability in the disk-magnetized star systems may also be associated with oscillations of the inner radius of the disk (e.g., Goodson, Winglee \\& B\\\"ohm 1997). We plan to further investigate the quasi-variability associated with such phenomena."
    },
    "0404/astro-ph0404339_arXiv.txt": {
        "abstract": "The CAIRNS (Cluster And Infall Region Nearby Survey) project is a large spectroscopic survey of the infall regions surrounding nine nearby rich clusters of galaxies.  I describe the survey and use the kinematics of galaxies in the infall regions to estimate the cluster mass profiles.  At small radii, these mass profiles are consistent with independent mass estimates from X-ray observations and Jeans analysis.  I demonstrate the dependence of mass-to-light ratios on environment by combining these mass profiles with Two-Micron All-Sky Survey (2MASS) photometry.  Near-infrared light is more extended than mass in these clusters, suggesting that dense cluster cores are less efficient at forming galaxies and/or more efficient at disrupting them.  At large radii, galaxy populations in cluster infall regions closely resemble those in the field.  The mass-to-light ratio at these radii should therefore be a good probe of the global mass-to-light ratio.  The mass-to-light ratio in the infall region yields a surprisingly low estimate of $\\Omega_m \\sim 0.1$. ",
        "introduction": "The relative distribution of matter and light in the universe is one of the outstanding problems in astrophysics.  Clusters of galaxies, the largest gravitationally relaxed objects in the universe, are important probes of the distribution of mass and light. \\cite{zwicky1933} first computed the mass-to-light ratio of the Coma cluster using the virial theorem and found that dark matter dominates the cluster mass.  Recent determinations using the virial theorem yield mass-to-light ratios of $M/L_{B_j}\\sim 250 \\mlsun$ (Girardi et al.~2000 and references therein).  Equating the mass-to-light ratio in clusters to the global value provides an estimate of the mass density of the universe; this estimate is subject to significant systematic error introduced by differences in galaxy populations between cluster cores and lower density regions (\\cite{cye97,g2000}).  Indeed, some numerical simulations suggest that cluster mass-to-light ratios exceed the universal value (\\cite{diaferio1999,kk1999,bahcall2000}).  Determining the global matter density from cluster mass-to-light ratios therefore requires knowledge of the dependence of mass-to-light ratios on environment.  Bahcall et al.~(1995) show that mass-to-light ratios increase with scale from galaxies to groups to clusters.  Ellipticals have larger overall values of $M/L_B$ than spirals, presumably a result of younger, bluer stellar populations in spirals. At the scale of cluster virial radii, mass-to-light ratios appear to reach a maximum value.  Some estimates of the mass-to-light ratio on very large scales ($>$10$\\Mpc$) are available (\\cite{bld95}), but the systematic uncertainties are large.  There are few estimates of mass-to-light ratios on scales between cluster virial radii and scales of 10$\\Mpc$ (Rines et al.~2000, Rines et al.~2001a, Biviano \\& Girardi 2003, Katgert et al.~2003, Kneib et al.~2003, Rines et al.~2004, Tully 2004 and references therein).  On these scales, many galaxies near clusters are bound to the cluster but not yet in equilibrium (\\cite{gunngott}). These cluster infall regions have received relatively little scrutiny because they are mildly nonlinear, making their properties very difficult to predict analytically.  However, these scales are exactly the ones in which galaxy properties change dramatically (e.g., Ellingson et al.~2001, Lewis et al.~2002, Gomez et al.~2003, Treu et al.~2003, Balogh et al.~2004, Gray et al.~2004). Variations in the mass-to-light ratio with environment could have important physical implications; they could be produced either by a varying dark matter fraction or by variations in the efficiency of star formation with environment.  In blue light, however, higher star formation rates in field galaxies could produce lower mass-to-light ratios outside cluster cores resulting only from the different contributions of young and old stars to the total luminosity (\\cite{bahcall2000,tully}).   Galaxies in cluster infall regions produce sharp features in redshift surveys.  Early investigations of this infall pattern focused on its use as a direct indicator of the global matter density $\\Omega_m$.  Unfortunately, random motions caused by galaxy-galaxy interactions and substructure within the infall region smear out this cosmological signal (Diaferio \\& Geller 1997, Vedel \\& Hartwick 1998). Instead of sharp peaks in redshift space, infall regions around real clusters typically display a well-defined envelope in redshift space which is significantly denser than the surrounding environment (\\cite{cairnsi}, hereafter Paper I, and references therein).  Diaferio \\& Geller (1997) and Diaferio (1999) analyzed the dynamics of infall regions with numerical simulations and found that in the outskirts of clusters, random motions due to substructure and non-radial motions make a substantial contribution to the amplitude of the caustics which delineate the infall regions. Diaferio \\& Geller (1997) showed that the amplitude of the caustics is a measure of the escape velocity from the cluster; identification of the caustics therefore allows a determination of the mass profile of the cluster on scales $\\lesssim 10\\Mpc$.  Diaferio \\& Geller (1997) and Diaferio (1999) show that nonparametric measurements of caustics yield cluster mass profiles accurate to $\\sim$50\\% on scales of up to 10 $h^{-1}$ Mpc.  This method assumes only that galaxies trace the velocity field. Indeed, simulations suggest that little or no velocity bias exists on linear and mildly non-linear scales (\\cite{kauffmann1999a,kauffmann1999b}).  The caustic method has been applied to systems as large as the Shapley Supercluster (\\cite{rqcm}) and as small as the Fornax cluster (\\cite{drink}) as well as to many nearby clusters (Paper I).  Biviano \\& Girardi (2003) applied the caustic technique to an ensemble cluster created by stacking redshifts around 43 clusters from the 2dF Galaxy Redshift Survey.  Rines et al.~(2000) found an enclosed mass-to-light ratio of $M/L_R \\sim 300 h$ within 4$~\\Mpc$ of A576.  Rines et al.~(2001) used 2MASS photometry and the mass profile from Geller et al.~(1999) to compute the mass-to-light profile of Coma in the K-band. They found a roughly flat profile with a possible decrease in $M/L_K$ with radius by no more than a factor of 3.  Biviano \\& Girardi (2003) find a decreasing ratio of mass density to total galaxy number density.  For early-type galaxies only, the number density profile is consistent with a constant mass-to-light (actually mass-to-number) ratio.  Here, we calculate the infrared mass-to-light profile within the turnaround radius for the CAIRNS clusters (Paper I), a sample of nine nearby rich, X-ray luminous clusters.  We use photometry from 2MASS, the Two Micron All Sky Survey (\\cite{twomass}) and add several new redshifts to obtain complete or nearly complete surveys of galaxies up to 1-2 magnitudes fainter than $M^*_{K_s}$ (as determined by Cole et al.~2001 and Kochanek et al.~2001).  Infrared light is a better tracer of stellar mass than optical light; it is relatively insensitive to dust extinction and recent star formation.  Despite these advantages, there are very few measurements of infrared mass-to-light ratios in clusters (\\cite{tustin,rines01a,lin03,cairnsii}).   \\begin{figure} \\centerline{\\psfig{file=combocaustics3.ps,angle=0,width=12.5cm}} \\caption{Redshift versus projected radius for the combined CAIRNS cluster.  The solid lines indicate the caustics and the errorbars show 1-$\\sigma$ uncertainties.} \\label{fig:caustics} \\end{figure} ",
        "conclusions": "\\label{sec:conclusions}  Cluster infall regions contain more galaxies than their virial regions.  If currently popular cosmological models are correct, the mass in infall regions will eventually accrete onto the parent clusters, and their final masses will increase by a factor of about 2.  Near-infrared luminosity functions depend only weakly on environment, at least at the bright end.  We show that near-infrared light is more extended than mass in cluster infall regions, suggesting environmental dependence of the efficiency of galaxy formation and/or disruption. If more efficient galaxy disruption is responsible, intracluster stars might be a significant component of stellar mass in clusters.  The mass-to-light ratios in infall regions suggest a low $\\Omega_m\\sim 0.1$.  Future work is needed to determine the significance of the conflict of this result with the currently favored $\\Omega_m\\sim 0.3$."
    },
    "0404/astro-ph0404080_arXiv.txt": {
        "abstract": "Reference systems and frames are crucial for high precision {\\it absolute} astrometric work, and their foundations must be well-defined.  The current frame, the International Celestial Reference Frame, will be discussed: its history, the use of the group delay as the measured quantity, the positional accuracy of 0.3 mas, and possible future improvements.  On the other hand, for the determination of the motion of celestial objects, accuracies approaching 0.01 mas can be obtained by measuring the differential position between the target object and nearby stationary sources. This {\\it relative} astrometric technique uses phase referencing, and the current techniques and limitations are discussed, using the results from four experiments.  Brief comments are included on the interpretation of the Jupiter gravity deflection experiment of September 2002. ",
        "introduction": "This paper is based on a talk at the JENAM2003 meeting in Budapest, Hungary in August 2003, and will cover several topics.  In \\S 2, reference systems are defined and a brief historical sketch is given. In \\S 3, the fundamental synthesis formula used for the determination of radio sources positions, and two basic observed quantities, the phase and the group delay, are discussed.  The International Celestial Reference Frame (ICRF) is described in some detail in \\S 4.  The measurement of relative positions to obtain the motion of radio sources, and calibrator choice concerns are then given in \\S 5, and in \\S 6, the description of four experiments highlight current techniques.  A brief conclusion is given in \\S 7, and an appendix concerning the controversy of the speed of gravity ends the paper. ",
        "conclusions": ""
    },
    "0404/astro-ph0404049_arXiv.txt": {
        "abstract": "{We present a systematic search for new neighbourhood stars among the Liverpool-Edinburgh high proper motion survey which we cross-identified with the DENIS survey. Their high proper motions ensure that they are not giant stars. The distances are estimated using DENIS photometry and we found that 100 stars probably lie within 25 parsecs from the Sun. They are mostly M-dwarfs, and 10 are probably white dwarfs. This is the first distance estimate for 84 stars among them. 10 stars, L 170-14A, DENIS J0146291-533931, DENIS J0235219-240038, LP 942-107, DENIS J0320518-635148, LTT 1732, DENIS J0428054-620929,  DENIS J2210200-701005, DENIS J2230096-534445, and NLTT 166-3 are estimated to be closer than 15 parsecs. In addition, one star, DENIS J2343155-241047, could also lie within 15~pc if it belongs to the halo. ",
        "introduction": "Whereas most stars more luminous than $M_V = 7$ have been detected in the solar neighbourhood \\citep{jahreiss1994}, intrinsically low luminosity objects such as white dwarfs, M dwarfs and brown dwarfs may have escaped our telescopes. Yet M stars are the dominant stellar constituent of the Milky Way, and the number density of brown dwarfs may be comparable to that of stars \\citep{reid1999}.  From a nearby-star inventory, \\cite{henry1997} estimated that so far 130 systems are missing from the solar neighbourhood sample within 10~pc. Within 25~pc, the distance limit of the Catalogue of Nearby Stars \\citep[CNS3,][]{gliese1991}, the missing fraction could be twice, that is $\\sim$70\\% \\citep{henry2002}. These numbers were obtained on the assumption that the census is complete within 5~pc. However, recent discoveries of a M5.5 star at trigonometric distance $d = 3.7$ pc \\citep{henry1997}, a M9 dwarf at spectroscopic distance $d \\sim 4$ pc \\citep{delfosse2001} and a M6.5 star at spectroscopic distance $d \\sim 3$ pc \\citep{teegarden2003} suggest that even members of the immediate solar neighbourhood remain undetected.  Projected on the sky, the apparent velocity of closer stars is greater that of farther stars. Thus, most of our known neighbours have high proper motions. From the NASA NStars database, we see that almost all of the known stars within 10~pc have a proper motion larger than 0.2$''$yr$^{-1}$ (see http://nstars.arc.nasa.gov). For many stars in high proper motion catalogues no distance has been determined yet. As part of the Research Consortium on Nearby Stars (RECONS) effort to discover new nearby stars, \\cite{henry2002} obtained spectral types for high proper motion stars and candidate nearby stars and found that six of them were closer than 25~pc.  Furthermore, the new surveys in the near-infrared DENIS and 2MASS provide unprecedented data for a systematic search for low luminosity cool dwarfs and brown dwarfs. The use of these data together with existing or new high proper motion catalogues is a powerful tool for discovering our neighbours. Hundreds of stars closer than 25~pc have been discovered this way \\citep{phanbao2001,phanbao2003,reid2002a,reid2002b,reyle2002,scholz2002,cruz2002,cruz2003}. Nearby stars have also been found in high proper motion catalogues from spectroscopic follow-up observations \\citep{lepine2003a,rojo2003}.   In a previous paper \\citep{reyle2002}, we presented the determination of the photometric distances of stars taken from the \\cite{scholz2000} high proper motion catalogue and cross-identified with the DENIS survey. Near-infrared photometry was used to determine the distances. 15 new stars were found to fall within the 25~pc CNS3 limit. Spectroscopic distances should soon be obtained for selected nearby candidates from low-resolution spectroscopic observations on the New Technologies Telescope at La Silla observatory in Chile.  Recently, \\cite{pokorny2003} published a large catalogue of high proper motion stars, the Liverpool-Edinburgh high proper motion survey. One of the purposes of this catalogue is to identify nearby objects. We have determined the photometric distance of the stars cross-identified with the DENIS survey.  This high proper motion catalogue is briefly described in Sect.~\\ref{lehpms}. Sect.~\\ref{denis} gives the result of the cross-identification with DENIS, presents the determination of photometric distances, and gives a list of the stars closer than 25~pc. The accuracy of our distance estimates is discussed in Sect.~\\ref{discuss}. ",
        "conclusions": "We made a systematic search for nearby stars by cross-identifying the Liverpool-Edinburgh high proper motion survey with the DEep Near-Infrared Survey DENIS. The photometric distance is determined using DENIS photometry. 100 stars are found to be closer than the 25~pc limit of the CNS3. 10 are white dwarfs, 3 may be halo subdwarfs, the remainder are disc M dwarfs, mainly from spectral type M3 to M6. It is the first distance determination for 83 among them. 20 stars are found to be closer than 15~pc, of which for 10 stars there was no previously known distance: L 170-14A, DENIS J0146291-533931, DENIS J0235219-240038, LP 942-107, DENIS J0320518-635148, LTT 1732, DENIS J0428054-620929,  DENIS J2210200-701005, DENIS J2230096-534445, and NLTT 166-3. If it is a halo subdwarf, DENIS J2343155-241047 may also be closer than 15~pc.  Given the large uncertainties in the photometric distances (up to 45\\% in some cases), follow-up observations of these candidates are needed. Low-resolution spectroscopy would assess the spectral type and provide more accurate determinations. This work should help in increasing the solar neighbourhood census. The nearest stars provide accurate data on fundamental parameters of stellar physics, such as luminosity, temperature, and mass. In particular, our understanding of low-mass stars relies upon the nearby stars, as they provide the only sample of intrinsically faint stars. They also are interesting targets for the future missions Terrestrial Planet Finder and DARWIN that will concentrate on very nearby stars to search for Earth-like exoplanets."
    },
    "0404/astro-ph0404094_arXiv.txt": {
        "abstract": "{% I investigate the self-irradiation of intensively accreting circumstellar discs (backwarmed discs).  It is modelled using the two-layer disc approach by \\citet{Lachaume03b} that includes heating by viscous dissipation and by an external source of radiation.  The disc is made of a surface layer directly heated by the viscous luminosity of the central parts of the disc, and of an interior heated by viscosity as well as by reprocessed radiation from the surface.  This model convincingly accounts for the infrared excess of some FU Orionis objects in the range 1--200\\,{\\micron} and supports the backwarmed disc hypothesis sometimes invoked to explain the mid- and far-infrared  excesses whose origins are still under debate. Detailed simulation of the vertical radiative transfert in the presence of backwarming is still needed to corroborate these results and spectroscopically constrain the properties of intensively accreting discs. } ",
        "introduction": "It is now widely accepted that circumstellar discs accompany the process of star formation, all the more as such discs have already been imaged in the millimetre \\citep[ex. DG~Tau,][]{Dutrey94} and in the infrared \\citep[ex. HK~Tau,][]{Stapelfeldt98}.  Matter accreting from the disc onto the star is supposed to build up solar-mass stars in the timescale of a \\Myr; yet recent studies of star forming regions indicate way too low accretion rates among T~Tauri stars (TTS) ---in the range $10^{-10}$--$10^{-7}\\,\\Msun/\\yr$, as shown by \\citet{Gullbring98} in the Taurus region and \\citet{Robberto03} in the Trapezium.  Part of the matter is supposed to be accreted in earlier phases of high accretion but young stellar objects (YSOs) are then so embedded that their disc cannot be optically observed \\citep[see][for the stages of evolution]{Andre94}.  The other part might also be accreted during brief, periodic phases of intense accretion of otherwise quiescent TTS \\citep{Hartmann96}:  FU~Orionis objects (FUors), which feature accretion rates up to $10^{-4}\\,\\Msun/\\yr$ and which are seen to undergo an increase in luminosity of more than 4 visual magnitudes in a few years \\citep{Herbig66,Herbig77}, may represent such a phase.  FUors are convincingly modelled by a self-heated viscous accretion disc (active disc) that overwhelms the stellar light and the properties of FUor outbursts are also well-studied \\citep{Hartmann85,Bell95}.  Yet active disc models fail to explain the mid- and far-infrared excess \\citep[eg.][]{Simon88} as well as weak silicate feature at $10\\,\\micron$ sometimes present in emission \\citep{Hanner98}; these properties seem more typical of an irradiated disc. Alternative models have tackled this problem:  \\citet{Lodato01} showed that the self-gravity of the disc can trigger an instability that produces additional warming and were able to reproduce the SED.  However, this model cannot account for the silicate feature, which requires a temperature inversion at the surface of the disc.  A more pleasing approach is the presence of a circumstellar envelope proposed by \\citet{Adams87}: on one (observational) hand, \\citet{Kenyon91} successfully fit the spectra of two FUors with such an envelope.  On the other (theoretical) hand, the envelope could serve as a reservoir of  infalling material, replenishing the disc between outbursts. On the third hand, the envelope can account for a part of the extinction observed in these objects.  \\citet{Kenyon91,Bell99} also proposed that the inner hot parts of the disc are bright enough to heat up the outer ones (backwarming) and produce an irradiated disc-like SED.  Their simulation uses a black-body disc model, ie.  without vertical temperature profile, and convincingly explains the order of the excess at 30--100\\,\\micron, yet it fails to reproduce the SED at 10--30\\,{\\micron} and does not predict a silicate feature in emission. \\citet{Malbet91,Calvet91,Chiang97} showed that irradiation can produce a hot disc surface radiating at shorter wavelengths and account for emission features (CO bands, silicates).  Using their results, \\citet{Lachaume03b} developed a two-layer disc model in which irradiation by a central star and viscosity are taken into account.  In this paper, I shall use this model with the disc itself as a source of radiation and investigate the backwarming in FUors. ",
        "conclusions": "Using the viscous and irradiated disc model by \\citet{Lachaume03b}, it is possible to account for the strength of the mid- and far-IR excess among a few FUors using the viscous luminosity of the inner parts of the disc as an irradiation source.  Yet the model makes several approximations in the determination of the structure:  the temperature profile is assimilated to two isothermal layers, a surface and an interior, which should not prove critical in the SED diagnostic; the flaring of the disc has been left as a free parameter, but it proves no less relevant than full numerical simulations, that assume the unknown amount of material in accretion discs at a scale $\\lesssim 100\\,AU$;  more critical is my leaving aside the self-gravity of the disc, which could ``unflare'' the disc and prevent irradiation, but the simulation of self-gravity in the presence of irradiation has not yet been performed, so its influence still remains speculation.  This work can be seen as a feasibility study for a future vertical structure simulation of irradiated discs that I am developping using the radiative transfer formalism presented in \\citet{Malbet01}.  Such a simulation is needed to issue a spectral diagnostic of irradiation (eg. silicate feature in emission), that forthcoming IR long-baseline insterferometers will ease with their AU-scale resolution for the closest FUors.  In particular, MIDI on the VLTI will be able to measure the silicate feature while disentangling the contribution of an envelope from that from an irradiated disc, which an SED diagnostic cannot do \\citep{Vinkovic03}.  One also expects to obtain constraints on the flaring with IR closure phases (AMBER on the VLTI) and on the disc mass with ALMA; their determination is not model-independent and also requires a reliable model.  \\begin{acknowledgement} This work has made use of NASA's Astrophysics Data System Bibliographic Services and of CDS's Vizier Catalogue Database.  Computations and graphics have been done with free software, in particular Yorick by D.~Munro. I also wish to thank C.~P.~Dullemond for helpful comments that improved the quality of the paper.  Language corrections have been suggested by K.~Smith. \\end{acknowledgement}"
    },
    "0404/astro-ph0404577_arXiv.txt": {
        "abstract": "The pulsating white dwarf star \\gone\\ (G~185-32) exhibits pulsation modes with peculiar properties that set it apart from other variable stars in the ZZ Ceti (DAV) class. These peculiarities include a low total pulsation amplitude, a mode with bizarre amplitudes in the ultraviolet, and a mode harmonic that exceeds the amplitude of its fundamental. Here, we present optical, time series spectroscopy of \\gone\\ acquired with the Keck II LRIS spectrograph. Our analysis has revealed that the mode with unusual UV amplitudes also has distinguishing characteristics in the optical. Comparison of its line profile variations to models suggests that this mode has a spherical degree of four.  We show that all the other peculiarities in this star are accounted for by a dominant pulsation mode of \\el=4, and propose this hypothesis as a solution to the mysteries of \\gone. ",
        "introduction": "While each member of the variable DA white dwarfs (DAVs) shows its own unique set of pulsations, this group of pulsators share many similar characteristics. They all have a pure hydrogen atmosphere and reside in a narrow temperature strip near 11,500~K, making them the coolest known class of white dwarf pulsators. The multi-periodic brightness variations of these stars are due to non-radial, g-mode pulsations \\citep{RKN} with periods between 100~s and 1000~s and amplitudes less than a few percent. Those pulsators near the blue end of the instability strip tend to show fewer modes with shorter periods and smaller amplitudes than the pulsators near the red edge \\citep{WF82}.  Understanding the characteristics of the pulsations of individual DAVs provides the opportunity to model their interiors. The ability to model the pulsations is limited by our ability to identify the modes. Pulsations are described by the spherical degree (\\el), azimuthal order ($m$), and radial order ($n$). Unfortunately, the paucity of observed modes in DAVs creates an obstacle to using period spacings for mode identification.  Yet, some successful mode identification has been performed on DAVs. Long data sets from a single site, or using the Whole Earth Telescope (WET) has resolved azimuthal splittings, yielding the identification of \\el\\ and $m$ for a few DAVs.  Time-resolved spectroscopy in both the UV and optical wavelengths, has been applied to the brighter DAVs \\citep[see][]{R95,VK00,Ke00,K02b,K02,T03,K03}. This method identifies the spherical degree by measuring how the amplitude of the mode changes with wavelength. Here, we apply the technique of optical, time-resolved spectroscopy to the bright DAV, \\gone\\ and attempt to decipher its prominent pulsation modes.  \\paragraph{About \\gone}  One of the brightest known DAVs (V$=12.97$ mag) with some of the smallest amplitude ($<0.3\\%$) pulsations is \\gone\\ (G~185-32, WD~1935+276). Since its discovery as a pulsator \\citep{M81}, it was noted as having a curious pulsation spectrum. Atypical of small amplitude pulsators, \\gone\\ displays a wide range of periods, including prominent modes near 370~s, 300~s, 215~s, 142~s, 72~s and 71~s \\citep{M81, Ke00, WET}. Given the star's shorter periods, low amplitudes, relatively stable modes, and temperature, this star is grouped with the pulsators near the blue edge of the instability strip.  The pulsation amplitudes on \\gone\\ are much lower than expected for DAVs of similar periods. DAVs follow a distinct trend: those with larger amplitude modes have larger mean periods \\citep{C94}. The trend created by these stars (see Figure~\\ref{davs}) reflects their similarities and follows the expectations of mode driving mechanisms \\citep{wg99,B92,Wi82}. \\gone\\ remains the exception to this trend. Both its average mode amplitude and the amplitude of its largest mode are approximately a factor of ten smaller than the other DAVs. Suggested explanations for the low amplitudes have included nonlinear pulsation modes with a relatively large number of surface nodes \\citep{M81}, a large magnetic field limiting the growth of its modes \\citep{C94}, and a large inclination causing cancellation of the modes \\citep{TC03}.  The mode at 71~s, the harmonic of the 142~s mode, is a curious feature of \\gone. First, the presence of large harmonic modes is normally reserved for the large amplitude pulsators. Second, the amplitude of this harmonic has an amplitude similar to the 142~s mode and has been observed to occasionally exceed the parent mode \\citep[see][]{M81}; all other DAVs show harmonics consistently smaller than the parent mode. The non-linear effects in the outer layers of DAVs, believed to be responsible for the presence of harmonics, should be small for low amplitude pulsations \\citep{B92, Wu01, IK}. Some possible explanations for why a low amplitude pulsator might display harmonics were discussed in general by \\citet{IK} in their study of nonlinear effects on pulsations. Nonlinear effects could appear with small amplitude pulsations if the star has an unusual surface convection zone, the star has a large inclination, or the mode has a large spherical degree (\\el).  The most recent addition to the mysteries of the pulsations of \\gone\\ came from time-resolved UV spectroscopy from the Hubble Space Telescope (HST) \\citep{Ke00}. They observed that each of the modes' relative amplitudes increased in the UV as an \\el\\ $= 1$ or $2$ mode except for the 142~s mode. It shows no appreciable increase in amplitude while its harmonic at 71~s still resembles a mode of \\el $\\le 2$. \\citet{Ke00} suggests that the 142~s mode is a result of nonlinear mixing while the other modes, including the 71~s mode, are real pulsations.  \\paragraph{The scope of this paper} In this paper, we add our analysis of optical time-resolved spectroscopy which offers an appealing explanation for all the mysterious observations of \\gone's modes. Our analysis of variations in the \\Hb\\ and \\Hg\\ lines show that the 142~s mode behaves like \\el=4. We therefore propose that this mode is the dominant mode in the star and that it has a spherical degree of four. This hypothesis can also account for the low amplitudes, the UV characteristics of the 142~s mode, and the presence of its large harmonic mode. We begin in \\S\\ref{data} by presenting time series spectra and discussing the variations of both the flux and velocity variations measured from the spectra. We then measure the pulsation amplitudes at each wavelength across the spectra, present a new method to perform this analysis, and compare the results with the models. In \\S3 we propose that this star is dominated by a mode with \\el=4 and show how this hypothesis is consistent with the observations. We discuss other explanations for the modes of \\gone\\ in \\S4 and finally, we summarize our conclusions in \\S5. ",
        "conclusions": "The DAV star \\gone\\ has a perplexing pulsation spectrum that has challenged attempts to understand the star. We have analyzed line profile variations that suggest the 142~s pulsation mode in \\gone\\ has a spherical degree of four.  If this mode is the star's dominant pulsation mode, then we can explain the star's overall low pulsation amplitude by geometric cancellation without the need to invoke improbable inclinations.  Furthermore, because the pulse shape harmonic of an \\el=4 mode has lower \\el\\ characteristics, it cancels less effectively, explaining the unusually large amplitude of the mode's harmonic.  Finally, an \\el=4 character of the 142~s mode and the consequent \\el=0,2 character of its harmonic are the very values most consistent with the UV amplitudes measured by the Hubble Space Telescope, although problems with pulsation phases remain.  Together, these results yield a satisfying and unified picture of the star's pulsations, and the only one consistent with all the data.  If our assignment of \\el=4 to the 142~s mode is correct, then it is a surprise since we have never before seen even \\el=3 and only seldom \\el=2.  Interestingly, the assignment of \\el=1 to most of the modes in hot DAVs by \\citet{C94} does not apply to the 142~s region of \\gone, which was one of two stars thought to have \\el$>$1. Based on period spacings alone, \\citet{C94} recognized that the 142~s or its nearby 148~s mode must be higher \\el\\ but had suspected \\el=2, not 4. If we are forced to consider higher \\el\\ during mode identification, then it will make that process difficult because the high \\el\\ modes are closely space in period.  Fortunately, it appears that high \\el\\ modes advertise that fact by having large harmonics which may assist in their identification.  Finally, we have developed a new analysis technique for time series spectroscopy that appears to work for relatively poor signal-to-noise.  The other DAVs studied by time series spectroscopy \\citep{VK00,K02,K02b,K03} might benefit from re-analysis by this technqiue, which we intend to do.  Also, we expect that by using this technique we will be able to extend mode identification with time-series spectroscopy to fainter stars, increasing the number of stars for which secure mode identification is possible."
    },
    "0404/astro-ph0404607_arXiv.txt": {
        "abstract": "Parker's interface dynamo is generalized to the case when a homogeneous flow is present in the high-diffusivity (upper) layer in the lateral direction (i.e.\\ perpendicular to the shear flow in the lower layer). This is probably a realistic first representation of the situation near the bottom of the solar convective zone, as the strongly subadiabatic stratification of the tachocline (lower layer in the interface dynamo) imposes a strong upper limit on the speed of any meridional flow there.  Analytic solutions to the eigenvalue problem are presented for the cases of vanishing diffusivity contrast and infinite diffusivity contrast, respectively. Unlike the trivial case of a homogeneous system, the ability of the meridional flow to reverse the propagation of the dynamo wave is strongly reduced in the interface dynamo. In particular, in the limit of high diffusivity contrast relevant to the solar case it is found that a meridional flow of realistic amplitude cannot reverse the direction of propagation of the dynamo wave. The implications of this result for the solar dynamo problem are discussed. ",
        "introduction": "While the operation of the solar dynamo is still far from understood (\\citenp{Weiss:CUPrev}, \\citenp{Petrovay:SOLSPA}), it is now generally believed that the strong toroidal magnetic field responsible for solar activity is generated and stored in the tachocline layer. This transitional layer between the differentially rotating convective zone (CZ) and the rigidly rotating solar interior is characterized by a strong radial shear, and is thus an ideal candidate for the production of the toroidal field.  The site and physical nature of the toroidal$\\rightarrow$poloidal flux conversion ($\\alpha$-effect), needed to close the cycle, is much less clear. One popular possibility is the interface dynamo (\\citenp{Parker:interface}---hereafter P93,\\citenp{Tobias:nonlin.interface}, \\citenp{Charbonneau+McGregor:IFdynamo}, \\citenp{Markiel+Thomas}, \\citenp{Mason+:competition}) where  $\\alpha$ is concentrated in the deepest part of the CZ, immediately above the layer of strong shear.  Helioseismic inversions (\\citenp{Kosovichev:tachocline}; \\citenp{Basu+Antia:tachovar}) indicate that at low latitudes the solar tachocline lies immediately below the adiabatically stratified CZ. This implies that the turbulent diffusivity in the tachocline is significantly suppressed in comparison to the layers immediately above it, resembling the situation envisaged in interface dynamo models.  An attractive feature of these models is that the sharp diffusivity contrast between the two layers gives rise to quite strong toroidal magnetic fields, in agreement with the requirements of flux emergence calculations. If the sign of $\\alpha$ is negative on the northern hemisphere, as expected near the bottom of the CZ, then the Parker--Yoshimura sign rule predicts an equatorward propagating dynamo wave at low latitudes and a poleward propagating wave at high latitudes, with the two belts departing at the corotation latitude of $\\sim35^{\\circ}$. It is tempting to identify these waves with the two well known branches of the extended butterfly diagram (e.g.\\ \\citenp{Makarov+Sivaraman}; cf. also \\citenp{Petrovay+Szakaly:2d.pol}).  The presence of a meridional circulation, however, significantly complicates the picture. As a poleward meridional flow of amplitude $\\sim 10$--$20\\,$m/s is clearly detected in the shallower layers of the CZ (\\citenp{Hathaway:merid.flow}), continuity seems to require a counterflow in the deep CZ (though a two-cell pattern is at present also not excluded). As this velocity amplitude is comparable to the speed of migration of the activity belts, the question arises whether the meridional flow can invalidate the Parker--Yoshimura sign rule, reversing the propagation of a dynamo wave. Indeed, in flux transport models of the solar dynamo (\\citenp{Wang+:1.5D}; \\citenp{Choudhuri+:mixed.transp}), a deep equatorward meridional flow is responsible for the migration of sunspot-forming latitudes during the solar cycle.  As the effect of a homogeneous flow can obviously be described by a Galilean transformation of the solution with no flow, it is indeed to be expected that an equatorward meridional flow pervading the whole dynamo region near the bottom of the CZ (as assumed in flux transport models) will reverse the poleward propagation of the dynamo wave if its speed is higher than the phase velocity. (In these models, $\\alpha$ is concentrated near the surface and is positive in the northern hemisphere, leading to poleward propagation at low latitudes ---hence the need for a reversal.) The actual situation is, however, more complex, as a meridional flow of significant amplitude cannot be expected to penetrate below the adiabatically stratified CZ. The strong subadiabatic stratification of the tachocline represents a serious obstacle in the way of meridional circulation. The timescale of any meridional flow here cannot be shorter than the (turbulent) heat diffusion  timescale, allowing downmoving fluid elements to get rid of their strong  buoyancy. Heat diffusivity in the subadiabatic layer is clearly much lower than above, so the same must be true for the amplitude of the poloidal flow. A more realistic representation of the meridional flow pattern in the dynamo layer, then, is an interface-type model with no meridional flow in the strongly sheared, low-diffusivity lower layer, and a homogeneous meridional flow in the highly diffusive upper layer.  This paper addresses the problem of how a meridional flow influences the properties of the dynamo wave by considering the arguably simplest nontrivial case: the effect of a homogeneous meridional flow in the top layer of Parker's interface dynamo. Section 2 describes the analytic model, Section 3 presents some solutions for important special cases. Finally, in Section 4 the implications of our findings are discussed for the problems outlined above. ",
        "conclusions": "We have studied by analytical methods how a meridional flow limited to the upper (i.e.~unsheared) layer affects kinematic interface dynamos. We found that the growth rate and period of the dynamo wave have a nontrivial dependence on the flow amplitude. In the case of strongly reduced diffusivity in the lower layer, relevant for the Sun, a flow parallel to the direction of propagation of the dynamo wave reduces the growth rate until, at a finite critical flow speed, it completely suppresses dynamo action. An antiparallel flow, in contrast, first increases the growth rate; then, after reaching a maximum, the growth rate tends to zero as $u_0\\rightarrow-\\infty$. Contrary to intuition, an antiparallel flow can neither suppress nor revert the direction of propagation of the dynamo wave.  This conclusion is apparently at odds with flux transport models (\\citenp{Choudhuri+:mixed.transp}, \\citenp{Dikpati+Charbonneau}) where an equatorward flow near the bottom of the CZ can turn the direction of propagation of the dynamo wave towards the equator at low latitudes. There is, however, no real contradiction here, as in those models the flow is assumed to penetrate quite deep into the layer of strong shear. In terms of the simple Cartesian model studied here, a closer counterpart of the flux transport models would be a case with $u_0=$const. throughout the region: then an antiparallel flow exceeding the phase speed can trivially turn around a dynamo wave. In Section~1 we argued that it is more realistic to assume that the meridional flow is limited to the adiabatic upper layer. In fact, the more recent flux transport model of \\citet{Dikpati+:polar.fields} admits that the circulation must be limited to the adiabatic layer. Therefore, in order to make the model work, they need to assume that a significant fraction of the tachocline overlaps the adiabatic SCZ. This may indeed hold at high latitudes, but at low latitudes, where a meridional transport of toroidal fields is most needed, it does not seem to be supported by helioseismic data (\\citenp{Basu+Antia:tachovar}).  These arguments indicate that our simplified models may have more relevance to the solar dynamo than many complex nonlinear spherical models where the assumed geometrical distribution of the flows and transport coefficient does not reflect the situation in the solar interior. Nevertheless, owing to the simplified  2D Cartesian geometry and the kinematical nature of our models our conclusions should be taken with proper reservation. In this context it may be mentioned that that details of the physical structure, such as the precise form of the rotation law, were shown to have a profound effect on the behaviour of dynamo models (\\citenp{Moss+:towards}, \\citenp{Phillips+:dyn.structure}). Extensions of this work to spherical geometry and to the 3D nonlinear domain are clearly needed to tell whether the present results remain valid in more realistic situations."
    },
    "0404/hep-ph0404042_arXiv.txt": {
        "abstract": "Solar neutrino physics enters a stage of precision measurements. In this connection we present a precise analytic description of the neutrino conversion in the context of LMA MSW solution of the solar neutrino problem. Using the adiabatic perturbation theory we derive an analytic formula for the $\\nu_e$ survival probability which takes into account the non-adiabatic corrections and the regeneration effect inside the Earth. The probability is averaged over the neutrino production region. We find that the non-adiabatic corrections are of the order $10^{-9}-10^{-7}$. Using the formula for the Earth regeneration effect we discuss features of the zenith angle dependence of the $\\nu_e$ flux. In particular, we show that effects of small structures at the surface of the Earth can be important. ",
        "introduction": "\\label{sec1}  The LMA MSW solution~\\cite{w,ms} has been identified ~\\cite{sno1}$-$\\cite{ped03} as the correct solution of the solar neutrino problem. The $2 \\nu $ conversion probability of this solution gives a very good description of all available data: no statistically significant deviation has been found so far. New physics effects beyond LMA, if exist, are below few per cent.  The program of future solar neutrino studies includes  1). further tests of the LMA solution, in particular, searches for signatures of this solution such as the Day-Night asymmetry and the distortion (``upturn\") of the boron $\\nu_e$ spectrum at low energies;   2). precise determination of the oscillation parameters, especially the 1-2 mixing angle;  3). searches for the sub-leading effects which originate from  - 1-3 mixing,  - sterile neutrino mixing,  - non-standard neutrino interactions,  - spin-flavor flip in the magnetic fields of the Sun,  - violation of the fundamental symmetries (CPT, equivalence principle, {\\it etc}.).\\\\   Already the present solar neutrino measurements have sensitivity at the level of few per cent. For instance, the predicted day-night asymmetry of the SuperKamiokande signal is about $2\\%$ which is comparable with the existing $1\\sigma$ experimental error~\\cite{sk-dn}. At SNO one expects the $2-4\\%$ asymmetry, consistent with the experimental result~\\cite{sno1} at the $1\\sigma$ level.  Future experiments will have substantially higher sensitivity~\\cite{BahPen,UNO,hyper-K,FREJUS}. The solar neutrino studies enter a phase of precision measurements. \\\\  In this connection it is important  \\begin{itemize}  \\item  to give precise description of the  LMA conversion, both in the Sun and in the Earth, taking into account various corrections;  \\item to estimate accuracy of the approximations made;  \\item  to find  the precise {\\it analytic} expressions for probabilities and observables as functions of the oscillation parameters ($\\Delta m^2$, $\\sin^2 \\theta_{12}$). This will help to test the LMA solution and to search for physics beyond LMA.   \\end{itemize}  We address these issues in the present paper. In section \\ref{sec2} we consider the non-adiabatic corrections to the LMA conversion probability. We calculate these corrections for propagation inside the Sun and the Earth. In section \\ref{sec3} we obtain the analytical formula for the probability averaged over the distribution of neutrino sources. In section \\ref{sec4} we derive the analytic formula for the $\\nu_e$ regeneration effect in the Earth. We present our conclusions in Section \\ref{sec5}. In the appendices A and B, alternative derivations of formulas for the regeneration factor are given. ",
        "conclusions": "\\label{sec5}  We have performed detailed analytic study of the LMA MSW conversion of the solar neutrinos. Our main result is the precise analytic formula for the survival probability which includes non-adiabatic corrections, averaging over the neutrino production region and the Earth regeneration effect. For the $K$ component of the solar neutrino spectrum ($K = pp, pep, Be, N, O, F,  B, hep$) it can be written as \\bea P_K = \\frac{1}{2}+\\frac{1}{2}(1-\\delta_K) \\cos 2\\theta_m({\\bar V}_K) \\cos 2\\theta -(1-\\delta_K) \\cos 2\\theta_m({\\bar V}_K) f_{reg}. \\eea Here the correction due to averaging effect, $\\delta_K$, is given in Eq. (\\ref{deltaN}); the average values of matter potential in the production regions of $K$ components, ${\\bar V}_K$, are defined in (\\ref{avepot}) and their numerical values are presented in the Table 1. The regeneration factor $f_{reg}$ is given in (\\ref{regnew}) for the symmetric density profile and in (\\ref{regasy}) for general asymmetric density profile.  Effect of averaging over the neutrino production region in the Sun is reduced to specific value of the initial mixing angle in matter which should be taken for the average value of the potential, $\\theta^0_m=\\theta_m({\\bar V}_K)$, and to the appearance of the correction $\\delta_K$. We have compared the analytic results with the results of numerical computation and found that maximal deviation $\\sim 1.8 \\%$ happens for the $hep$ neutrinos. For the boron neutrinos the precision is better than 0.2\\%.  We have obtained precise analytic formula for the regeneration effect in the Earth using the realistic  density profile. We present simple derivation of this formula which uses the adiabatic perturbation theory. Performing also explicit calculations of the evolution in sequent layers we show that this derivation is correct. The analytic formula reproduces results of numerical computations with accuracy determined by $\\eta \\sim 1-2 \\%$.  Essentially the regeneration effect is the sum of contributions from different shells which are determined by jumps of the potential at the borders  and by the adiabatic phase acquired inside the outer  borders of the corresponding shells. The dependence of regeneration factor on the zenith angle can be understood in terms of interference of contributions from different borders.  The derived analytical formula allows us to understand the effect of averaging over the neutrino energy. Using the analytical formula we have considered effects of small scale structures ($\\sim 10$ km) of the Earth profile. These effects can be important for small values of $\\cos \\theta_Z$. We stress that local ``perturbations'' of the density profile can produce sizable uncertainties in $f_{reg}$.\\\\  \\noindent {\\Large \\bf Acknowledgment}  One of the authors (P.C.H.) would like to thank FAPESP for financial support. A.Y.S. thanks Tokyo Metropolitan University where this work has been accomplished for hospitality."
    },
    "0404/astro-ph0404188.txt": {
        "abstract": "During long-term observations of the Galactic Centre region in hard X-rays (10-300 keV) in space experiments made on board Prognoz-9 satellite and \"Mir\" orbital station (GRIF experiment) some periodic sources were revealed. They include periodicities of hour and day range of period: 152 h, 98 h, 82 h (4U1700-37), 69 h, 62 h, 13.3 h, 9.36 h, 8.03h, 8.15 h (Cen X-4), 4.38 h, 4.35 h (4U1755-33) and 3.45 h. For all the observed periodic processes the mean phase profiles (light curves) in the different energy ranges were obtained. The mean phase profiles of most of the observed periodic sources differ from both sine and purely eclipse form. Among six sources with day ranges of period at least three were identified with X-ray Novas - black hole candidates: 152 h (H1705-25, Nov.Oph., 1977), 62 h (GRO J 1655-40, Nov. Sco., 1994), 13.3 h (4U1543-47). The results of the GRO J 1655-40 and 4U1543-47 observations in the Prognoz-9 and GRIF \"Mir\" experiments were compared with the observational data obtained in space observatories CGRO, RXTE, BeppoSax. It should be noted that the orbital periodicity is revealed in the hard emission of the X-ray transient - black hole candidates GRO J 1655-40 and 4U1543-47 even in the epochs between outbursts. Periodic processes with day ranges of periods can be typical for the transient source like X-ray Nova - black- hole candidate. Such periodicity is orbital but its origin is not connected with eclipses of the compact companion of a binary system. Periodic processes in soft and hard X-rays can be associated with orbital motion of binary system companions, but the physical mechanisms and the regions of generation of soft and hard X-rays can be quite different. ",
        "introduction": "The enigma of black holes is one of the most intriguing problems of modern astrophysics. It is well known that black holes (BH) as well as neutron stars (NS) are associated with hard X-ray sources, which are characterized by great energy release as the result of accretion. The temporal behavior of such sources was always a problem of great interest for high-energy astrophysicists. Modern observations are usually made in a wide range of wavelengths from radio waves to gamma rays. While optical observations are in some aspects more accurate, the observation of X-ray emitting objects in hard radiation allows us to study the processes in the region of great energy release. Moreover, the information on temporal features is necessary for the exact identification of a hard source.  Various instabilities in the accretion process lead to the flaring activity of many hard X- ray sources. The properties of matter and fields surrounding a compact object determine quasi- periodic variations, while the rotation of binary system components produce orbital and super- orbital periodic changes in X-ray source luminosity. These factors are significantly different in dependence on characteristic time. The short-term variability is caused, particularly by quasi- periodic oscillations (QPO), the properties of which in some cases are different for two classes of objects - X-ray binaries with neutron star (NSXB) and with Black Hole (BHXB). Only NS can produce QPO with frequency up to 1 kHz \\citep{sun00}. Long-term variability of most of BHXB-s is caused by their flaring activity. It is known that most of soft X-ray transients (X-ray nova) manifest recurrent flares. The typical duration of an X-ray flare is about 50-100 days and the time interval between flares varies around 1.5 years. However, a number of X-ray nova flared only once during all the period of observations. It is not yet clear, whether there are any soft X-ray transient objects with no recurrent flares. The shape of X-ray flares strongly varies from source to source and even from one flare to another in the same source. Most of them have standard shape with a \"fast rise and an exponential decay\" but other forms are more complicated with several maximums. An outstanding example is the 1994-1997 activity of BHXB GRO J1655-40, which demonstrated a large series of quasi-periodic flares, separated from one another by several months \\citep{zha94,tav96,zha97,hyn98}.  Variations with minute's, hour's and day's characteristic time scales can occur when X-rays are produced or absorbed in some large structures surrounding the compact object as well as in the compact object itself. Accretion instabilities or the motion of binary system components may cause them. Sometimes they can appear as regular variations of spectra or dips observed at the determined orbital phases. The study of such medium-scale temporal phenomena is of great importance. Monitoring observations are the traditional way of studying temporal phenomena in hard X-rays and soft gamma rays such as periodic processes, outbursts and transients. Periodic variations of the hard emission from some galactic sources were discovered long ago. The first observations were made on the UHURU space observatory \\citep{gia71,for78}, and then studied in details during several missions (see for example, \\citep{cor86,nag89,pre87} and others). Nevertheless, to the present time not so many objects manifesting as periodic sources specifically in the hard range of electromagnetic spectrum are known. This is caused by some methodical problems, in particular, the necessity of long-term continuous observations of each source, which is rather difficult even in the case of monitoring observations particularly for long- period processes of the orbital and accretion disk precession type. Extensive data on periodicity in hard emissions from galactic sources were obtained during the BATSE experiment onboard the Compton Orbital observatory (CGRO) \\citep{rob97}. A study of temporal phenomena in hard X-rays and soft gamma rays was also made in the OSSE CGRO experiment \\citep{joh93}. More or less regular observation of the Galactic Centre region (in the limits of $\\sim10^o$x$10^o$) were carried out for over 5 years by means of the \"SIGMA\" telescope on-board the \"GRANAT\" observatory. This experiment allowed to detect new sources at energies of about 50-100 keV and to give the notion about the temporal properties of these and previously known objects \\citep{gol94}. Currently monitoring observations of high-energy properties of X-ray sources are being conducted on the BeppoSAX mission \\citep{fro98}. The most complete catalogue of the orbital periodicities observed in the X-ray range as well as the transient sources was obtained as the result of observations during the Rossi XTE (RXTE) mission \\citep{lev99,bra00a,bra00b}.  Although a large amount of observational data have been obtained by the present time, it seems that the experimental studies of periodicity as well as the other temporal phenomena (primarily outbursts) in X-ray sources have not yet exhausted their potential. One of the main problems is that not all of the observed periodicities can be unambiguously interpreted. In particular, what is the nature of the so-called non-eclipsing periodicities in the orbital range of periods? Several sources are known, which manifested periodic variations in the X-ray flux during which the X-ray luminosity does not fall to zero. Such a periodicity (8.2 h) was discovered in the emission from Cen X-4 during the well-known outburst of 1978, May \\citep{kal80} observed in the Ariel-5 mission. The dip-type periodicity was detected during the EXOSAT observations. It was observed in the X-rays from several low-mass binaries including MXB1759-29 (7.1 h), 4U1822-37 (5.6 h), 4U1755-33 (4.4 h), X1323-62 (3.8 h), XB1254-69 (2.4 h), etc \\citep{com84,mas81,whi84,par89,cou86}. Although periodic dips are not full eclipsing the common view is that all these periodicities are orbital. The 62-h periodicity in hard emission of GRO J1655-40 source also is not full eclipsing. Moreover in soft X-rays 62-h periodicity reveals itself as a number of deep drops in the intensity occurring between orbital phases 0.7 and 0.9 \\citep{kuu98}. The light curve corresponding to this periodicity looks like a dip-type one with minima dislocated from the optical occultation. Thus, simple mechanisms such as a total or partial eclipse of the emitting hot region by the optical component do not give a complete explanation of the temporal behavior of the tight binary system's hard emission. Another problem, which still remain is the lack of data concerning the observation of periodic processes in the high-energy part of the electromagnetic spectrum. For most of the known tight binaries their orbital period values and corresponding light curves were obtained as the result of optical observations. On the contrary, only complete data from both optical and high-energy observations can give us the necessary possibility, which allows to determine the orbital parameters as well as the mass of binary system components. Thus, it seems important to find the manifestation of orbital periodicity in hard emission for the larger number of tight binaries.  Search for periodicity in the emission from objects with great irregular (noise-type) flux variations or sporadic flux rises (transients) can also bring us closer to the understanding of the nature of such sources. The interest to the latter has increased in connection with the observations on the CGRO of the so-called transient pulsars \\citep{fin96,zha96,wil97} and sources with periodically recurring transient rises in the X-ray flux \\citep{bay93,dek93}. Study of temporal and spectral properties of sources known as X- ray transients can clarify the nature of such objects and give us important information about high energy processes responsible for hard emission generation. From this point of view the search for periodic processes in hard emission during outbursts of sources known as X-ray novas is of particular interest. Most of X-ray binaries - black-hole candidates are just X-ray novas or transient sources \\citep{sun91,cher97,van98}. Thus the study of temporal behavior of such objects can give us independent criteria which permit to identify black holes. For example, identification or, on the contrary, non-identification of pulsation-type periodic process during the outburst can clarify the nature of a compact companion, because the presence of pulsation excludes the black-hole variant. On the other hand, the spectral information, such as the presence of absorption lines at cyclotron frequencies, can also indicate the presence of a magnetic neutron star. On the contrary, the presence of an annihilation line is a strong indication of a black hole as the compact companion \\citep{sun91}. To the present time the information on transient X-ray sources was obtained primarily during their outbursts \\citep{chen97}. However, the observations of such sources between outbursts can also be rather fruitful. There are some predictions that X-ray fluxes from transient sources even in quiescent state are non-zero. In particular, according to \\citep{las00} the X-ray luminosity of quiescent black hole low-mass X-ray binary transients should be correlated with their orbital period. It is very unlikely that X-rays from such systems are emitted by coronae of companion stars. More convincing models associate their X-ray luminosity with accreting mechanisms, providing the flux generation at the observation level of such missions as Chandra or XMM. The measurements by Chandra showed that typical quiescent intensity of BH X-ray nova is about 10-6 Crab in 0.5-7 keV energy range. There are some indications that quiescent X-ray flux from BH X-ray nova is significantly lower than from NS ones and this fact can be caused by the existence of event horizon \\citep{gar00}. Thus, the observations of X-ray transients in quiescent state can give unique information clarifying the nature of such objects. Below we will discuss the results of the search and study of periodic processes in hard X- rays from some Galactic sources including the black hole candidates during the observations which were carried out on the Prognoz-9 and \"Mir\" orbital station (OS) missions. ",
        "conclusions": "Most of the periodic sources of hard emission, which were observed in Prognoz 9 and GRIF \"Mir\" experiments, are characterized by non-eclipsing mean light curves. An exception is the well- known X-ray eclipsing binary 4U1700-37. Nevertheless there is no doubt that most of these periodicities have an orbital origin. From this point of view, it is of particular interest that except 4U1700-37 and 4U1755-33 all identified sources of hour and day periodicity (H1705-22, GRO J1655-40, 4U1543- 47, Cen X-4) appear to belong to transient objects. Because sources of 98-h and 67-h periodicities were not observed in the GRIF experiment, we can assume that they were not active in 1995-1997, which could be the consequence of strong variability of their hard emission and, hence, the transient character also. We can also see that among the periodic sources with periods in day range there are three X- ray novae (Ophiuchis 1977; Scorpius, 1994, 4U1543-47), which are known as black-hole candidates. However, there were no direct indications of overall activity of X-ray nova Ophiuchis 1977 and Scorpius, 1994 as well as 4U1543-47 in 1983-1984. Detection of the above mentioned periodicities was made when their activity was not very large. Thus, we can conclude that some X- ray transients - black hole candidates between outbursts can be active at the level $\\sim 100$ mCrab, sufficient to detect their hard emission by sensitive detectors. Moreover, some kind of orbital, but non-eclipsing periodicity could be observed in high-energy emission in the time of \"intermediate\" activity of these objects, i.e. such X-ray transients could be not quite quiescent between outbursts. It is very likely that this periodicity can be caused by features of the accretion disk surrounding the black hole and may be inherent to other sources of X-ray transient - black hole candidate. To the present time there were no clear observations of the development of an orbital periodicity in hard emission. Our observations of non-eclipsing orbital periodicity in hard X-rays from black-hole candidates just when they were not outburst-like active could be connected with some features of the zones of the X-ray formation in these objects. This can be discussed in some details using the example of GRO J 1655-40. The Prognoz 9 and GRIF \"Mir\" observations allow to conclude that there is a strict periodicity in the hard emission from GRO J 1655-40, which is traceable during difference cycles of the source's activity separated by a long time interval ($\\sim 12$ years). This periodicity is apparently attributable to the orbital motion of the binary's components. The energy spectrum of the 62-h periodicity in the emission from GRO J 1655-40 \\citep{kud01} agrees with various spectral measurements of the total hard flux from this source. Thus, it is consistent with the model of a two-component spectrum typical of compact binaries, black-hole candidates \\citep{sha73}. The observation of the 62-h periodicity in the hard emission from GRO J 1655-40 during different epochs suggests that even between its X-ray outbursts, the source was not completely quiescent in the hard energy range. In particular, it was active between the X-ray outbursts in August 1995 and February 1996. This gives some indications that it could be possible to reduce the typical recurrence time, which is important for estimating the mean rate mass transfer from the binary's optical companion tip to its compact object $\\dot M_T$ \\citep{esi00} by at least half. The GRIF data also suggest, that taking the source's distance of $\\sim 3.5$ kpc \\citep{hje95}, its luminosity between X-ray outbursts in the periodicity of hard emission at energies $E \\geq 25$ keV is $\\sim 4 \\cdot 10^{35}$ erg/s, which is more than three orders of magnitude higher than the luminosity that has been believed to be typical of the quiescent state of such objects \\citep{esi00}.  The intensity of the periodic component of GRO J1655-40 in soft X-rays obtained from ASM RXTE observations increases with the increase of the full intensity. It reaches a $3\\%$ level of the full flux in the whole range of its values. On the other hand, there are some indications on that during the time, when the source brightness decreasing after outbursts, in particular, between outbursts of 1996 and 1997 years the periodic component in hard X-rays was not low, but on the contrary was becoming more contrast. The results of GRIF \"Mir\" observations indicate that orbital periodicity revealed in the GRO J1655-40 hard emission even in the interval between time of its high X-ray activity. In this case the emission of GRO J1655-40 in the phase of minimum intensity has not fallen to zero. It should be noted that high flux variability during epochs of the source high activity makes it difficult to reveal the periodic component with high significance. The source's variability is associated not only with periodic variations because the rms deviations of the instrument's outputs in all the intervals of observations are several times higher than the amplitude of periodic component, which can be revealed only due to averaging over many periods. The BATSE CGRO data revealed no manifestations of the periodicity in the hard emission from GRO J1655-40 \\citep{zha96}. This could be because the amplitude of the 62-h periodicity in the hard energy range weakly depends on the total flux. In this case, the periodicity would show up most clearly when the source's activity is relatively low (to all appearances, the Prognoz 9 and GRIF \"Mir\" observations occurred precisely at such epochs). However, during outbursts, the total flux can increase by more than an order of magnitude. This severely hampers the observation of the periodicity, because it proves to be strongly noised by fluctuations in the total flux. Thus, we can conclude, that although the time profiles of mean light curves, which characterize the periodic component in GRO J1655-40 emission are similar in soft and hard X-rays, 62-h periodicity manifests quite different behaviour in these energy ranges. This could be caused not only by different physical mechanisms of generation of soft and hard X-rays in the source, but even by emission of hard and soft X-rays by different regions of the accretion disk. This assumption does not contradict the spectral features of GRO J1655-40 hard emission, which consists of, at least, two or more components. The soft thermal component, which is dominant at energies less than 10 keV and hard non-thermal \"comptonisation\" component, which becomes dominant at higher energies, were revealed clearly. It also should be noted that the minimum on the mean phase dependence as in soft as in hard X-rays is displaced relatively to the phase zero of the optical light curve. Such non- coincidence of optical and X-ray light curves for emission periodic components caused by the orbital motion could be typical not only for GRO J1655-40. For example, we saw before, the periodic component of 4U1543-47 emission also demonstrates different behaviour in optical and X- ray ranges. The optical light curve gives the orbital period, which is twice larger than the period observed in X-rays. The quasi-ellipsoidal form of the star-companion of a binary system can explain the non-coincidence of these periods. In any event, the features mentioned above as well as the complicated form of mean light curves of periodic component in hard X-rays give no opportunity to explain the observed periodicity in X-ray binaries - black-hole candidates using the simple eclipsing geometry in orbital motion of binary system companions. It indicates that processes, which provide such periodicity, depend also on the accretion disk dynamics, and they may reflect the accretion disk structure non-uniformity. Thus, the main conclusions drawn from observations of periodic processes in hard emission of X-ray binaries - black-hole candidates, which were made on Prognoz 9 and GRIF \"Mir\" experiments with the use of data of CGRO and RXTE observations are: - periodic processes of days range of periods are inherent to the transient source like a X- ray Nova - black-hole candidate; - such periodicity is orbital but its origin is not connected with eclipses of the compact companion of a binary system; - periodic processes in soft and hard X-rays are associated with orbital motion of binary system companions, but the physical mechanisms and the regions of generation of soft and hard X-rays are quite different."
    },
    "0404/astro-ph0404482_arXiv.txt": {
        "abstract": "For the first time in literature, we provide an exact analytical scheme for calculating the analogous Hawking temperature and surface gravity for general relativistic accretion onto astrophysical black holes. Our calculation may bridge the gap between the theory of transonic astrophysical accretion and the theory of analogous Hawking radiation, and may allow one to compare the properties of the two types of horizons, electromagnetic and acoustic, in connection to the black hole physics. We show that the domination of the analogous Hawking radiation over the actual Hawking radiation may be a real astrophysical phenomena. We also discuss the possibilities of the emergence of analogous white holes around astrophysical black holes. Our calculation is general enough to accommodate accreting black holes with any mass. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404161_arXiv.txt": {
        "abstract": "{A high-accuracy multisite campaign was carried out from 2002 January to May with a photometric coverage of 398 hours at five observatories. The concentration on a few selected sites gives better consistency and accuracy than collecting smaller amounts from a larger number of sites. 23 frequencies were detected with a high statistical significance. 6 of these are new. The 17 frequencies found in common with the 1992--1995 data are the modes with highest amplitudes. This indicates that the pulsation spectrum of FG~Vir is relatively stable over the ten-year period. Two frequencies have variable amplitudes and phases from year to year as well as during 2002. These were both found to be double modes with close frequencies. For the mode at 12.15 c/d this leads to an apparent modulation with a time scale of $\\sim$129d. The close frequencies at 12.15 c/d are composed of a radial and a nonradial mode, suggesting a similarity with the Blazhko Effect seen in RR Lyrae stars. ",
        "introduction": "Asteroseismology of $\\delta$~Scuti stars has reached a stage where the choice between different models of stellar structure and evolution requires a large number of known pulsation frequencies. While the discovery of many new frequencies cannot be accompanied by successful pulsation mode identifications, the latter is important to refine the stellar pulsation models. So far, the observations could not yet be matched with perfect models due, in part, to lack of enough observational constraints.  Consequently, the Delta Scuti Network (DSN) specializes in the intensive study of a few selected pulsating stars. One of these objects is the star FG~Vir, which was observed for several months during 2002. The motivation for additional photometry was:  (i) Successful mode identification relies on both photometric and spectroscopic techniques. While the first method is used to determine the pulsational $\\ell$ values through phase differences and to a lesser extent, amplitude ratios, the line-profile technique allows accurate determinations of the $m$ values and determines the rotational velocity as well as the aspect of the rotation axis, $i$. In 2002, for the first time the photometric campaign has been paired with an intensive multisite spectroscopic campaign.  (ii) A comparison of the photometric amplitudes with those obtained during previous campaigns from 1992--1995 would yield valuable information on the long-term stability of the pulsation spectrum.  (iii) Furthermore, a long campaign would lead to a high frequency resolution within the year of observation. This would permit the search for close frequency pairs, such as those that have been reported for several other $\\delta$~Scuti variables.  (iv) The mode selection mechanism operating in $\\delta$~Scuti stars is still unknown. Theoretical pulsation models predict considerably more modes than are observed. Consequently, it is important to study the stars in more detail to search for modes with small amplitudes and to increase the number of known frequencies.  The star FG~Vir (=~HD~106384) has been observed before: its variability was originally discovered by Eggen (1971) in one night of observation. During 1992, Mantegazza et al. (1994) measured FG~Vir photometrically for 8 nights. They were able to identify six frequencies of pulsation, while a seventh mode of pulsation was also suggested. Two multisite photometric campaigns were organized by the Delta Scuti Network during 1993 (170 hours, Breger et al. 1995) and 1995 (435 hours, Breger et al. 1998). The multisite campaigns led to the discovery of 24 frequencies of pulsation, of which 21 were independent pulsation modes. A reasonable agreement between different attempts towards mode identifications was found (Guzik et al. 1998, Viskum et al. 1998, Breger et al. 1999, Mantegazza \\& Poretti 2002).  The Delta Scuti Network is engaged in a long-term study of FG~Vir, using both photometric and spectroscopic techniques. Consequently, even larger and better analyses are expected in the future. These will include photometric frequency determinations, photometric and spectroscopic mode identifications as well as stellar modelling. The purpose of the present paper is to report the 2002 photometry and its implications for our understanding of FG~Vir. ",
        "conclusions": "16 out of the 17 frequencies found to be statistically significant in both the 1992--1995 and 2002 data are the modes with the largest photometric amplitudes. While this is hardly surprising, it also provides an argument in favor of the stability of the pulsation spectrum of FG~Vir. In this regard the star differs from some other $\\delta$~Scuti stars such as 4~CVn (see Breger 2000), in which measurements taken ten years apart may even provide a first impression of belonging to two different stars! Actually, in 4~CVn, the 'disappeared' modes have been shown to be still present, but with much smaller amplitudes.  The 'missing' modes in FG~Vir all had $V$ amplitudes between 0.4 and 0.8 mmag during 1995 with a statistical uncertainty of $\\pm$ of 0.11 mmag. The 2002 data show that most may still be present with the peaks at the correct frequencies, but with amplitudes less than 0.5 mmag. This is below the significance limit. We conclude that the differences can be explained by observational uncertainties in most cases."
    },
    "0404/astro-ph0404357_arXiv.txt": {
        "abstract": "The properties of stellar populations in the centers of nearby lenticular galaxies are investigated by means of 2D spectroscopy. All the galaxies are divided into 4 groups depending on the environment type; every subsample contains more than 10 galaxies. Clear distinctions between the mean stellar ages and abundance ratios both for the nuclei and for the bulges of the S0s in the different environments are found. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404211_arXiv.txt": {
        "abstract": "Recent results have shown that many of the known extrasolar planetary systems contain regions which are stable for massless test particles.  We examine the possibility that Saturn-mass planets exist in these systems, just below the detection threshold, and attempt to predict likely orbital parameters for such unseen planets.  To do this, we insert a Saturn-mass planet into the stable regions of these systems and integrate its orbit for 100 million years.  We conduct 200-600 of these experiments to test parameter space in HD37124, HD38529, 55Cnc, and HD74156. In HD37124 the global maximum of the survival rate of Saturns in parameter space is at semimajor axis $a$ = 1.03 AU, eccentricity $e \\sim$ 0.1.  In HD38529, only 5\\% of Saturns are unstable, and the region in which a Saturn could survive is very broad, centered on $0.5<a<0.6$, $e<0.2$. In 55Cnc we find three maxima at ($a$, $e$) = (1.0 AU, 0.02), (2.0 AU, 0.08), and (3.0 AU, 0.17). In HD74156 we find a broad maximum with $a$ = 0.9-1.2 AU, $e \\leq$ 0.15.  Several of these maxima are located in the habitable zones of their parent stars and are therefore of astrobiological interest.  We suggest the possibility that companions may lie in these locations of parameter space, and encourage further observational investigation of these systems. ",
        "introduction": "There are currently 110 known extrasolar planets, including ten systems containing two or more planets.  These planets are known to be Jovian both from their large masses, which range from 0.11 Jupiter masses (HD49674; Butler \\etal 2003) to 17.5 Jupiter masses (HD202206; Udry \\etal 2002), and from their sizes, measured in HD209458 to be 1.27 Jupiter radii (Charbonneau et al. 2000).  The vast majority of these planets were discovered by the radial velocity technique, which is sensitive to roughly 3-10 $m \\, s^{-1}$(Butler \\etal 1996; Baranne \\etal 1996).  All planetary systems must be dynamically stable for at least the age of their host star.  Recent work by Barnes \\& Quinn (2004) suggests that a large fraction of systems are on the edge of stability: a small change in semimajor axis $a$ or eccentricity $e$ causes the system to become unstable. The ``packed planetary systems'' (PPS) hypothesis presented in Barnes \\& Raymond (2004; hereafter Paper 1) predicts that all planetary systems are ``on the edge.''  This leads to speculation that those systems which appear stable may harbor unseen planets which push them to the edge of stability.  The PPS hypothesis suggests that if a region exists in a planetary system in which the orbit of a massive planet is stable, then its presence is likely.  The first paper of this series (Paper 1) used integrations of massless test particles to map the stability of regions in certain extrasolar planetary systems in ($a$, $e$) space.  Of the five systems examined, three (HD37124, HD38529, and 55Cnc) were found to contain zones between the giant planets in which test particles were dynamically stable for 5-10 Myr.  Stable regions have been found in $a$ space (assuming circular orbits) for $\\upsilon$ And (Rivera \\& Lissauer 2000), GJ876 (Rivera \\& Lissauer 2001) and 55Cnc (Rivera \\& Haghighipour 2003).  In this work we test for the presence of unseen Saturn-mass planets in four known extrasolar planetary systems: HD37124 (Butler \\etal 2003), HD38539 (Fischer \\etal 2003), 55Cnc (Marcy \\etal 2002), and HD74156 (Naef et al 2004).  We choose Saturn-mass planets because they lie roughly at the detection threshold for the current radial velocity surveys (Butler \\etal 1996).  The reflex velocity caused by a Saturn-mass planet at 1 AU on a solar-mass star is 8.5 $m \\, s^{-1}$, and scales with the planet's semimajor axis as $a^{-1/2}$.  For comparison, the smallest amplitude reflex velocity of any detected planet is 11 $m \\, s^{-1}$ (HD1641; Marcy \\etal 2000). Although seven sub-Saturn mass planets have been discovered as of November 2003 (e.g. Fischer \\etal 2003), none has $a >$ 0.35 AU.\\footnote{Data from http://www.exoplanets.org}  Paper 1 found that no test particles survived in HD74156 for longer than a few Myr.  However, Dvorak \\etal (2003) found orbits stable for test particles between 0.9 and 1.4 AU. We therefore include HD74156 in our sample.  Table 1 shows the orbital parameters for the four extrasolar planetary systems we investigate.  Note that the best fit orbital elements for some systems, especially HD74156c, have changed many times.  We therefore adopt elements as of a given date, with the knowledge that they may fluctuate. In $\\S$2 we describe our initial conditions and numerical method.  We present the results for each planetary system in $\\S$3, and compare these with other work in $\\S$4.  We present our conclusions in $\\S$5. ",
        "conclusions": "We have found specific locations in four known extrasolar planetary systems in which Saturn-mass planets could exist on stable orbits.  Such a planet would lie just below the detection threshold of current radial velocity surveys, and may be detected in the near future.  Table 3 summarizes our results, detailing the location in ($a,e$) space of each maximum in the survival rate for each of our four candidate systems.  If an additional planet is discovered in the stable region of one of these systems, it would mark the first successful prediction of a planet since John Couch Adams predicted the existence of Neptune in 1845 based on perturbations to Uranus' orbit.  Does the presence of a stable region imply the presence of a planet?  Must all systems contain as many planets as they can?  Laskar (1996) speculated that ``a planetary system will always be in this state of marginal stability, as a result of its gravitational interactions.''  The ``packed planetary systems'' (PPS) hypothesis, presented in Paper 1 (see also Barnes \\& Quinn, 2004), extends this idea by suggesting that all systems contain as many planets as they can dynamically support without self-disrupting. All systems may be on the edge of stability, but observational constraints prevent the detection of smaller or more distant bodies which push apparently stable systems to this edge.  The formation scenario of a planet of any size in between two gas giant planets is of great interest.  In the Solar System no stable regions exist between the orbits of the gas giants.  The detailed formation scenario of a smaller giant planet between two others is unclear, be it through gravitational instability (e.g. Mayer \\etal 2002) or core-accretion (Pollack \\etal 1996).  Gas giant planets at small orbital radii may have formed farther out in the protoplanetary disk and migrated inward, which further complicates this formation scenario.  Certain stable regions in HD37124, 55Cnc and HD74156 are located in the habitable zones of their parent stars (see Table 3).  Clearly, the discovery of a planet of any size in these regions is of great astrobiological importance, as any giant planet would likely have one or more large moons.  Understanding the formation of terrestrial planets in these systems is vital. In the upcoming third paper of the ``predicting planets'' series (Raymond \\& Barnes 2004) we present results of simulations of terrestrial planet formation in between the known giant planets in the same four systems examined here."
    },
    "0404/astro-ph0404283_arXiv.txt": {
        "abstract": "{ New radio data at 330 MHz are presented for the rich clusters Abell 665 and Abell 2163, whose radio emission is characterized by the presence of a radio halo.  These images allowed us to derive the spectral properties of the two clusters under study.  The integrated spectra of these halos between 0.3 GHz and 1.4 GHz are moderately steep: $\\alpha^{1.4}_{0.3}$ = 1.04 and $\\alpha^{1.4}_{0.3}$ = 1.18, for A665 and A2163, respectively.  The spectral index maps, produced with an angular resolution of the order of $\\sim$ 1\\arcmin, show features of the spectral index (flattening and patches), which are indication of a complex shape of the radiating electron spectrum, and are therefore in support of electron reacceleration models.  Regions of flatter spectrum are found to be related to the recent merger activity in these clusters.  This is the first strong confirmation that the cluster merger supplies energy to the radio halo. In the undisturbed cluster regions, the spectrum steepens with the distance from the cluster center.  This is interpreted as the result of the combination of the magnetic field profile with the spatial distribution of the reacceleration efficiency, thus allowing us to set constraints on the radial profile of the cluster magnetic field. ",
        "introduction": "Recent observations of clusters of galaxies have revealed a new and complex scenario in the structure of the intergalactic medium.  The clusters are not simple relaxed structures, but are still forming at the present epoch.  Substructures, commonly observed in the X-ray distribution of a high number of rich clusters (Henry \\& Briel 1993, Burns \\etal 1994), are evidence of the hierarchic growth of clusters from the merger of poorer subclusters.  \\begin{table*} \\caption{Observed clusters of galaxies} \\begin{flushleft} \\begin{tabular}{lllllllllll} \\hline \\noalign{\\smallskip} Name  & z  &  RA (J2000) & DEC  & T & L$_X$ (bol) & r$_c$ & r$_c$ & $\\beta$ & P$_{1.4}$ & LS \\\\ &    &  h~~ m~~ s  & \\degrees~~~~ \\arcmin~~~~ \\arcsec~~ & keV & erg s$^{-1}$ & \\arcsec  &  kpc & & W Hz$^{-1}$ & Mpc\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} A~665  & 0.1818 &  08 30 57.4 & +65 51 14.4  & 9.03  & 4.17 10$^{45}$ & 96 & 379 & 0.66  & 6.6 10$^{24}$  & 1.8  \\\\ A~2163 & 0.203  &  16 15 49.4 & --06 09 00.0  & 13.83 & 1.33 10$^{46}$ & 72 & 305 & 0.62  & 3.0 10$^{25}$  & 2.9  \\\\ \\noalign{\\smallskip} \\hline \\label{olog} \\end{tabular} \\end{flushleft} \\par\\noindent Caption. Col. 1: cluster name; Col. 2: redshift; Cols. 3 and 4: coordinates of the X-ray cluster center from Ebeling et al. (1996); Col. 5: temperature from Allen \\& Fabian (1995); Col. 6: bolometric X-ray luminosity from Wu et al. (1999); Cols. 7, 8 and 9: cluster core radius (angular and linear) and $\\beta$ from Birkinshaw et al. (1991)  for A665 and Elbaz et al. (1995) form A2163 ; Cols. 10 and 11: Total radio power and maximum linear size at 1.4 GHz from Giovannini \\& Feretti (2000) for A665 and Feretti et al. (2001) for A2163. \\end{table*}  An important problem in cluster phenomenology involves cluster-wide radio halos, whose prototype is Coma C (Giovannini \\etal 1993, and references therein). These are extended diffuse radio sources located at the cluster centers, with typical sizes of \\gtsim 1 $h_{50}^{-1}$ Mpc, regular shape, steep radio spectra and no significant polarization.  According to recent suggestions, the cluster merger process may play a crucial role in the formation and energetics of these sources (see Giovannini \\& Feretti 2002, and references therein).  The origin and evolution of halos is still a matter of debate. Several suggestions for the mechanism transferring energy into the relativistic electron population and for the origin of relativistic electrons themselves have been made: in-situ reacceleration of relativistic electrons by plasma and by shock waves, particle injection from radio galaxies, acceleration out of the thermal pool, secondary electrons resulting from hadronic collisions of relativistic protons with the ICM gas protons, and combinations of these processes (see e.g. Brunetti 2003, Blasi 2003, Petrosian 2003).  It is therefore important to carry out new observations aimed at discriminating between these theoretical models.  An important observable in a radio halo is the spectral index distribution.  It could allow us to test predictions of these different models, since it reflects the shape of the electron energy distribution. Spectral index maps represent a powerful tool to study the properties of the relativistic electrons and of the magnetic field in which they emit, and to investigate the connection between the electron energy distribution and the ICM.  By combining high resolution spectral information and X--ray images it is possible to study the thermal--relativistic plasma connection both on small scales (e.g. spectral index variations vs. clumps in the ICM distribution) and on the large scale (e.g. radial spectral index trends).  The prototypical example of cluster radio halos is the diffuse source Coma C in the Coma cluster (Willson 1970).  Nowadays Coma C is the only radio halo for which a high resolution spectral index image has been presented in the literature (Giovannini et al., 1993).  The spectral index trend in Coma C shows a flat spectrum\\footnote {$S_{\\nu}\\propto \\nu ^{-\\alpha}$ through this paper} in the center ($\\alpha\\simeq$ 0.8) and a progressive steepening with increasing distance from the center (up to $\\alpha\\simeq$ 1.8). Since the diffusion velocity of relativistic particles is low in relation to their radiative lifetime, the radial spectral steepening cannot be simply due to ageing of radioemitting electrons. Therefore the spectral steepening must be related to the intrinsic evolution of the local electron spectrum and to the radial profile of the cluster magnetic field.  It has been shown (Brunetti et al. 2001) that a relatively general expectation of models invoking reacceleration of relic particles is a radial spectral steepening in the synchrotron emission from radio halos.  The steepening, that is difficult to reproduce by other models such as those invoking secondary electron populations, is due to the combined effect of a radial decrease of the cluster magnetic field strength and of the presence of a high energy break in the energy distribution of the reaccelerated electron population.  In the framework of reacceleration models the radio spectral index map can be used to derive the physical conditions prevailing in the clusters, i.e. reacceleration efficiency and magnetic field strength.  This method has been successfully applied to the case of Coma C by Brunetti et al. (2001) who applied a {\\it two phase} reacceleration model and obtained large scale reacceleration efficiencies of the order of $\\sim 10^{8}$ yr$^{-1}$ and magnetic field strengths ranging from 1--3 $\\mu$G in the central regions down to 0.05--0.1 $\\mu$G in the cluster periphery.  In this paper we present the radio images at 90 cm of Abell 665 and Abell 2163. These clusters host powerful and giant radio halos studied at 20 cm with the VLA (Giovannini \\& Feretti 2000, Feretti et al. 2001).  From the comparison between the 20 and 90 cm data, we derive the spectral index maps of these radio halos and discuss their implications.  We adopt H$_0$=50 km s$^{-1}$ Mpc$^{-1}$ and q$_0$ = 0.5. With these values, 1 arcsec corresponds to 3.94 kpc at the distance of A665 and to 4.27 kpc at the distance of A2163.   \\begin{table} \\caption{Observing log} \\begin{flushleft} \\begin{tabular}{llllllll} \\hline \\noalign{\\smallskip} Name & $\\nu$ &  $\\Delta \\nu$  & Conf  &  Date  & Dur \\\\ & MHz & MHz  & & & h \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} A665 & 321.5/327.5 & 3.125 & B & 4/2001 & 3  \\\\ &  `` & `` & C & 9/2001 & 6 \\\\ & `` & `` & D & 12/2001 & 1.5 \\\\ A2163 &  321.5/327.5 & 3.125 & B & 4/2001 & 3.5   \\\\ &  `` & `` & C & 8/2001 & 6  \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{flushleft} Caption. Col. 1: cluster name; Col. 2: observing frequencies; Cols. 3: bandwidth; Col. 4: VLA configuration; Col. 5: month and year of observation; Col. 6: observing time duration. \\end{table}   \\begin{figure*} \\special{psfile=fig1.ps hoffset=-70 voffset=-600 hscale=100 vscale=100 } \\vspace{14 cm} \\caption{ Radio map at 90 cm of A665. The beam FWHM is 68\\arcsec $\\times$ 59\\arcsec~ in PA= 25\\degrees. Contour levels are -2, 2, 3, 4, 6, 8, 10, 15, 30, 50 mJy/beam.} \\end{figure*}  \\begin{figure*} \\special{psfile=fig2.ps hoffset=30 voffset=-500 hscale=80 vscale=80} \\vspace{14 cm} \\caption{ Radio map at 90 cm of A2163. The beam FWHM is 57\\arcsec $\\times$ 51\\arcsec~ in PA= 2\\degrees. Contour levels are -1, 1, 2, 3, 5, 7, 10, 25, 50 mJy/beam.} \\end{figure*} ",
        "conclusions": "In this paper we have presented new images at 0.3 GHz of the radio halos in the clusters A665 and A2163 and we have derived the spectral index maps between 0.3 and 1.4 GHz, using maps already published at 1.4 GHz.  Our results show that the integrated radio spectra of these halos are moderately steep in this frequency range: $\\alpha^{1.4}_{0.3}$ = 1.04 and $\\alpha^{1.4}_{0.3}$ = 1.18, for A665 and A2163, respectively.  The spectral index maps, produced with an angular resolution of the order of $\\sim$ 1\\arcmin, show a clumpy distribution with significant variations, which are indication of a complex shape of the radiating electron spectrum, and are therefore in support of halo models invoking the reacceleration of relativistic particles.  We find that the regions of flatter spectrum appear to trace the geometry of recent merger activity as suggested by X-ray maps.  These results prove that the radio spectral index can be a powerful tracer of both the current physical properties and past history of clusters.  We find no evidence of spectral flattening at the location of the hot shock detected in A665 (Markevitch \\& Vikhlinin 2001).  This favours the scenario that cluster turbulence might be the major responsible for the electron reacceleration.  In the undisturbed cluster regions, the spectrum steepens with the distance from the cluster center.  Brunetti et al. (2001) explain the spectral steepening in the framework of electron reacceleration models, as due to the combined effect of a radial decrease of the cluster magnetic field strength and of the presence of a high energy break in the energy distribution of the reaccelerated electron population. This is more difficult to reproduce by other models, such as those requiring secondary electron populations.  The spectral steepening detected in A665 and A2163 in the direction not presently affected by the merging processes allowed us to constrain the profile of the product between the cluster magnetic field and the reacceleration efficiency, $B \\chi^2$, under simple assumptions.  The magnetic field profile in both clusters is flatter than predicted in the case of a constant reacceleration and magnetic field frozen to the cluster thermal gas.  The ongoing violent mergers may play a crucial role in determining the conditions of the radiating particles and of the magnetic fields in clusters."
    },
    "0404/astro-ph0404556_arXiv.txt": {
        "abstract": "We present  results from a set of high ($512^3$ effective resolution), and ultra-high ($1024^3)$  SPH adiabatic cosmological simulations of cluster formation aimed at  studying  the internal  structure of the intracluster medium (ICM). We discuss the radial structure and scaling relations expected from purely gravitational collapse, and show that the choice of a particular halo model can have important consequences on the interpretation of observational data. The validity of the approximations of hydrostatic equilibrium and a polytropic equation of state are checked against results of our  simulations. We also show  the first  results from an unprecedented large-scale  simulation of 500 $h^{-1}$ Mpc and $2\\times 512^3$ gas and dark matter particles. This experiment will make possible a detailed  study of the large-scale distribution of clusters  as a function of their  X-ray properties. ",
        "introduction": "Galaxy clusters are  a unique laboratory to test the hierarchical paradigm of structure formation. They are the best probes  of the large scale structure of the Universe and have often been used as a diagnostic of the cosmological parameters. The intrinsic non-linear nature of gravitational collapse and gas dynamics makes numerical simulations the most useful tool to study in detail the process of cluster formation and evolution.   Simple analytical models for the structure of the ICM can be derived from  the  hypotheses of  hydrostatic equilibrium and polytropic equation of state, $P\\propto\\rhogas^\\gamma$. But real clusters might not be well described by these two hypothesis. For instance, kinetic energy makes a significant contribution to the energy budget of merging systems, and therefore thermally-supported hydrostatic equilibrium ceases to be a valid approximation. Even in relaxed systems, this assumption is not very accurate in the outermost parts, where gas motions become more important. Departures from spherical symmetry can also play a role in the final structure of the ICM \\citep{LeeSuto03} and, last but not least, there is no obvious physical reason for the gas to follow a polytropic relation. From our numerical experiments, we showed (see \\citet{Ascasibar03}) that hydrostatic equilibrium is fulfilled  within $\\sim20$ per cent accuracy by all simulated clusters, as long as they are not heavily disturbed. A polytropic equation of state seems to be a good approximation as well, although its reliability near the centre is still a matter of debate \\citep[see e.g.][]{Rasia_03}. From our data we derive a polytropic index of $\\gamma\\sim 1.18$.   We compared four different analytical  halo models with  our simulations. The first two models assume that haloes are well described by  \\citet{NFW97} and  \\citet{Moore99}  fitting formulae, while the other two assume that the gas follows a \\bm \\citep{CavaliereFusco76}. We consider a 'canonical' version of the \\BM, in which the gas is isothermal ($\\gamma=1$) and $\\beta=2/3$, and a polytropic version with $\\gamma=1.18$ and $\\beta=1$. The same value of the polytropic index has been used for the first two models as well. Hereafter we will use the abbreviations NFW, MQGSL, BM and PBM to refer  to these models. For a detailed description, the reader is referred to \\citet{Ascasibar03}. ",
        "conclusions": ""
    },
    "0404/hep-th0404161_arXiv.txt": {
        "abstract": "Braneworld inflation is a phenomenology related to string theory that describes high-energy modifications to general relativistic inflation. The observable universe is a braneworld embedded in 5-dimensional anti de Sitter spacetime. When the 5-dimensional action is Einstein-Hilbert, we have a Randall-Sundrum type braneworld. The amplitude of tensor and scalar perturbations from inflation is strongly increased relative to the standard results, although the ratio of tensor to scalar amplitudes still obeys the standard consistency relation. If a Gauss-Bonnet term is included in the action, as a high-energy correction motivated by string theory, we show that there are important changes to the Randall-Sundrum case. We give an exact analysis of the tensor perturbations. They satisfy the same wave equation and have the same spectrum as in the Randall-Sundrum case, but the Gauss-Bonnet change to the junction conditions leads to a modified amplitude of gravitational waves. The amplitude is no longer monotonically increasing with energy scale, but decreases asymptotically after an initial rise above the standard level. Using an approximation that neglects bulk effects, we show that the amplitude of scalar perturbations has a qualitatively similar behaviour to the tensor amplitude. In addition, the tensor to scalar ratio breaks the standard consistency relation. ",
        "introduction": "In recent years, there has been considerable interest in the possibility that our observable four-dimensional (4D) universe may be viewed as a brane hypersurface embedded in a higher-dimensional bulk space. Of particular importance is the Randall-Sundrum (RS) model, where a single, positive-tension brane is embedded in a five-dimensional (5D) anti de Sitter $({\\rm AdS}_5)$ spacetime~\\cite{RSII}. (For recent reviews, see Ref.~\\cite{royreview}.) Although the fifth dimension may be infinite in extent, the zero-mode of the 5D graviton, corresponding to 4D gravitational waves, is localized at low energies on the brane due to the warped geometry of the bulk. This property can also be understood within the context of the AdS/CFT correspondence~\\cite{adscft}, where the RS model is viewed as 4D gravity coupled to a conformal field theory (CFT)~\\cite{hrr}.  A natural extension of the RS model that is motivated by string/M~theory considerations, is to include higher-order curvature invariants in the bulk action. Such terms arise in the AdS/CFT correspondence as next-to-leading order corrections to the CFT~\\cite{largeN}. The Gauss-Bonnet (GB) combination of curvature invariants is of particular relevance in five dimensions, since it represents the unique combination that leads to second-order gravitational field equations in the bulk metric and since the field equations contain only linear second derivatives~\\cite{d86}. A GB term may also arise as the next-to-leading order correction in the heterotic string effective action, and it is ghost-free about flat spacetime~\\cite{stringGB}. Moreover, the graviton zero mode remains localized in the GB braneworld~\\cite{onlylocal} and deviations from Newton's law at low energies are less pronounced than in the RS model~\\cite{milder}.  From an observational perspective, there is now strong evidence that the very early universe underwent an epoch of accelerated (inflationary) expansion~\\cite{wmap}. During inflation, light fields such as the graviton are quantum-mechanically excited and acquire nearly scale-invariant fluctuations~\\cite{gw}. The resulting spectrum of primordial gravitational waves could be detectable from its imprint on the polarization of the cosmic microwave background (CMB)~\\cite{Bmode}. Such a detection would open a unique window into the physics of the very early universe. The evolution of gravitational waves during slow-roll inflation has been determined in the RS scenario~\\cite{lmw}. At high energies, the amplitude is enhanced relative to the standard result in 4D Einstein gravity. In view of the above developments, therefore, it is important to determine the properties of tensor perturbations generated during inflation in the GB braneworld. We show that significant changes to the RS case are introduced by the GB term, even when the GB corrections are very small relative to the Einstein-Hilbert terms. ",
        "conclusions": "Brane inflation offers a phenomenology that allows us to explore some of the cosmological implications of ideas arising from string and M~theory. The effects on inflationary perturbations from the extra dimensional nature of gravity introduce new features that need to be computed and then subjected to the constraints from high-precision cosmological data. Here we have concentrated on computing the corrections to the standard results for tensor and scalar perturbations that are generated during slow-roll inflation at energies where brane effects become dominant. This has previously been done for the Randall-Sundrum braneworld, based on 5-dimensional Einstein gravity. We have introduced a Gauss-Bonnet term, since string theory arguments indicate that this term is a high-energy perturbative correction to the gravitational action. This correction leads to significant qualitative changes, even in the perturbative regime $\\beta\\equiv 4\\alpha\\mu^2 \\ll 1$.  For the tensor perturbations, we have given an exact analysis, including the 5D modes. The wave equation and its fundamental solutions are not changed by the GB term. The spectrum contains a normalizable zero-mode and a continuous tower of massive modes after a mass gap, $m>{3\\over 2}H$, as in the RS case. The massive modes are not excited during inflation, as in the RS case. However, the GB term changes the boundary conditions at the brane, and therefore changes the normalization of the zero-mode, as shown by Eq.~(\\ref{F}): \\[ {A_T^2 \\over [A_T^2]_{4\\rm D}}= \\left[\\sqrt{1+x^2} -\\left(\\!\\frac{1-\\beta} {1+\\beta}\\!\\right) x^2\\sinh^{-1}{1\\over x}\\right]^{-2}\\!. \\] This leads in the GB regime to a suppression of tensor perturbations relative to the standard result, unlike the enhancement that arises in the RS case $\\beta=0$.  For the scalar perturbations, we used an approximation where bulk perturbations decouple from the density perturbations. We showed that the GB modifications to the Friedman equation lead to a significant change from the RS case, as given by Eq.~(\\ref{gg}): \\begin{eqnarray*} {A_S^2 \\over [A_S^2]_{4\\rm D}}= \\left[{3(1+\\beta) x^2 \\over 2\\sqrt{1+x^2}(3-\\beta +2\\beta x^2)+2(\\beta-3)}\\right]^{3}\\!. \\end{eqnarray*} These perturbations are again suppressed in the GB regime relative to the standard result, unlike the RS enhancement.  Because the scalar suppression is stronger than the tensor suppression in the GB regime, the relative tensor contribution, as a fraction of the scalar amplitude, $R=A_T^2/A_S^2$, is enhanced in the GB regime, in comparison with the standard result. This is shown by Eq.~(\\ref{ratio}): \\begin{eqnarray*} &&{R \\over R_{4\\rm D}}= \\\\&&{}{[2\\sqrt{1+x^2}(3-\\beta +2\\beta x^2)+2(\\beta-3)]^3\\over 27(1+\\beta)^2x^6[(1+\\beta)\\! \\sqrt{1+x^2}-(1-\\beta) x^2\\sinh^{-\\!1}x^{-\\!1}]}\\!. \\end{eqnarray*} By contrast, in the RS case the relative tensor contribution is suppressed.  Furthermore, the consistency relation between the tensor/ scalar ratio and the tensor spectral tilt is different in the GB case, i.e., \\[ R=-\\left({ 1+\\beta + 2\\beta x^2 \\over 1+\\beta + \\beta x^2} \\right){n_T \\over 2}\\,, \\] by Eq.~(\\ref{Qx}). The RS model by contrast has the same consistency relation as the standard case, $R=-{1\\over2}n_T$.  Our results provide a basis on which to confront the GB braneworld with observational constraints, and this is under investigation.    \\[ \\] {\\bf Acknowledgements:}  We thank Pierre Binetruy, Christos Charmousis, Naresh Dadhich, Stephen Davis, Lev Kofman, David Langlois, Jihad Mourad, Sergei Odintsov, Renaud Parentani, Danielle Steer and David Wands for helpful discussions. The work of RM is supported by PPARC."
    },
    "0404/astro-ph0404360.txt": {
        "abstract": "We report 3 - 13 $\\mu$m spectroscopy of 4 comets observed between August 2002 and February 2003: C/2002 O4 (H\\\"{o}nig) on August 1, 2002, C/2002 V1 (NEAT) on Jan. 9 and 10, 2003, C/2002 X5 (Kudo-Fujikawa) on Jan. 9 and 10, 2003, and C/2002 Y1 (Juels-Holvorcem) on Feb. 20, 2003. In addition, we include data obtained much earlier on 69P/Taylor (February 9, 1998) but not previously published. For Comets Taylor, H\\\"{o}nig, NEAT, and Kudo-Fujikawa, the silicate emission band was detected, being approximately 23\\%, 12\\%, 15\\%, and 10\\%, respectively, above the continuum. The data for Comet Juels-Holvorcem were of insufficient quality to detect the presence of a silicate band of comparable strength to the other three objects, and we place an upper limit of 24\\% on this feature. The silicate features in both NEAT and Kudo-Fujikawa contained structure indicating the presence of crystalline material. The shape of the silicate feature at a projected distance of 1900 km from the nucleus of Kudo-Fujukawa was nearly identical to that centered on the nucleus, indicating that the grain size population had not been measurable modified by the time it had reached that distance. Combining these data with those of other comets, we confirm the correlation between silicate band strength and grain temperature of Gehrz \\& Ney (1992) and Williams et al. (1997) for dynamically new and long period comets, but the majority of Jupiter family objects may deviate from this relation. Despite the weakness of the silicate band in Kudo-Fujikawa, its structure resembles the bands seen in dynamically new and long period objects with substantially stronger features. The limited data available on Jupiter family objects suggest that they may have silicate bands that are slightly different from the former objects. Finally, when compared to the silicate emission bands observed in pre-main sequence stars, the dynamically new and long period comets most closely resemble the more evolved stellar systems, while the limited data (in quantity and quality) on Jupiter family objects seem to suggest that these have spectra more like the less-evolved stars. Higher quality data on a larger number of Jupiter family objects are needed to confirm (or reject) this trend. ",
        "introduction": "One of the major goals of astronomy today is to understand the origin and evolution of solar systems, including our own. In the same way that stellar astrophysics is tied to our detailed understanding of the Sun for guidance, our ability to understand the evolution of extrasolar planetary systems in anything more than a general and tentative way will require detailed knowledge of how the most easily accessed system has developed. One of the key pieces of information that is required to model the formation and evolution of the solar system is the \u00d2boundary condition\u00d3 of its initial state. And while understanding the specific details of the initial collapse of the protosolar cloud may remain an unreachable goal at the present time, much of the physical and chemical nature of our solar system during its first $10^{6}$ years, when it was only 0.02\\% of its present age, remains in the form of solids stored beyond the orbit of Neptune where the equilibrium temperature drops below 50 K. These materials, true interstellar material that survived intact, material that condensed out of the early nebular gas, and material processed during the earliest stages of nebular development, are accessible in the form of comets.  The presence of silicate material in comets was first confirmed by \\citet{maas70}, who unambiguously detected the emission band of silicates in the spectrum of Comet C/1969 Y1 Bennett using filter photometry with a spectral resolution R = $\\lambda$/$\\Delta\\lambda\\sim 10$. Since then, numerous bright comets have been observed using broad-band and intermediate-band photometry \\citep{ney74a,ney74b,gn92,att86a,att86b,msh87}, and this technique continues to be used to this day. While it lacks the mineralogical diagnostic power obtainable with higher spectral resolution, its broad wavelength coverage (2-23 $\\mu$m and often extending to 0.7 $\\mu$m) has made it a powerful tool for understanding the nature of cometary dust grains. Two important discoveries made with this technique are that the strength of the silicate band varies from object to object, and that most bright comets have grains that radiate at temperatures in excess of that of a blackbody at the expected equilibrium temperature for their heliocentric distance. Both effects are expected for a grain population that includes a significant fraction of the grains having sizes smaller that the wavelength of light being radiated, in this case from 3-20 $\\mu$m, and the two have been found to be correlated in the emission from bright comets \\citep{gn92,wil97}.  The data relevant to the population of grain sizes may be dependent on the position within the coma where the observations are made, and therefore also dependent on the beam size used in the measurement. The grain size distribution measured in the coma of 1P/Halley by the Giotto and Vega spacecraft varied noticeably as a function of distance from the nucleus and between pre- and post-encounter tracks \\citep{mcd87,mcd91}. This is not surprising, as the entrainment of dust by subliming volatiles will tend to sort grains by mass (or the ratio of mass to surface area for highly porous grains), and fragmentation of grains during outflow will tend to produce small grains from large ones. Ground-based photometry of C/1973 E1 Kohoutek indicated that the grains in the anti-tail were considerably larger than those in the central coma and tail \\citep{ney74a}. Such behavior is expected in comets due to the sorting by radiation pressure (see, for example, \\citet{gj90}).  To properly characterize the grain sizes, thermal properties, and mineralogical composition of relic solar system grains, spectroscopic data of a large number of comets are required. Objects ranging from short-period comets to dynamically new ones need to be included. Comets that have entered the inner solar system for the first time will have surfaces that have been exposed for billions of years to processing by cosmic rays, and may not necessarily be shedding the more pristine materials to be found deeper in their interiors. Short-period objects that have endured thousands of years (or more) of surface activity will have had this outer layer removed. However, because smaller grains are more easily entrained by sublimating volatiles from the nucleus than larger ones, their population of observable grain sizes will tend to be skewed toward larger values than what they were born with. Hence, data on a large number of objects at various stages of their evolution will be required to disentangle the various factors affecting their grain properties. In this paper we report our observations of 5 comets, and analyze the data for these and 14 others for which data at comparable spectral resolution and wavelength coverage exist. ",
        "conclusions": "We have reported previously unpublished mid-IR spectra of 5 comets: C/2002 O4 (H\\\"{o}nig), C/2002 X1 (NEAT), C/2002 X5 (Kudo-Fujikawa), C/2002 Y1 (Juels-Holvorcem), and 69P/Taylor. In 4 of these object (H\\\"{o}nig, NEAT, Kudo-Fujikawa, and Taylor) we detected the presence of the emission band due to silicates. In the other (Juels-Holvorcem) the band was weak or absent, but the quality of the data was low. Four objects (H\\\"{o}nig, NEAT, Taylor, and Juels-Holvorcem) contained grains radiating at temperatures generally above those expected of a blackbody at the heliocentric distance of the comet, although for any individual object the results are within 2 standard deviations of the blackbody temperature.  In the one object where we obtained data both on and offset from the nucleus of Comet Kudo-Fujikawa, the dust emission characteristics seemed unchanged. The timing of the observations were such that they were sampling the same grains if all were traveling at a constant velocity equal to that observed in Comet Hale-Bopp. At least in this one object, any possible evidence for significant modification of the grain size by grain size sorting or disintegration was not evident. The near lack of excess temperature and weakness of the silicate band indicated that the grains ejected in Kudo-Fujikawa were not dominated by small silicate grains seen in most  ``young'' comets. However, the shape of the silicate band itself requires it be produced by whatever population of small grains that were present.  Merging this sample with objects observed earlier at similar spectral resolutions and wavelength coverage, a few notable characteristics emerge. The strong correlation between excess temperature and silicate band strength previously reported is confirmed here for the majority of the dynamically new objects in our sample. However, there is a tendency for Jupiter family objects to deviate from this correlation in the sense that they tend to have temperature characteristics close to the middle of the range exhibited by the sample as a whole, but silicate bands that are consistently weaker than the average. If this trend is confirmed with a larger sample, then it would suggest that our solar system possessed a radial gradient in the size distribution of its silicate grains, similar to that observed in $\\beta$ Pic today.  It is apparent that the spectra of solar system comets possess many features in common. Most well-observed objects exhibit temperatures in excess of that expected from a blackbody at the equilibrium temperature expected for their heliocentric distances. Most show a silicate emission feature, although the strength can be quite different from object to object, or even from epoch to epoch for a single object. When strong enough to measure reliably, the silicate band usually has a trapezoidal shape above the continuum, usually with a rather distinct peak at 11.2 $\\mu$m feature attributed to crystalline olivine. However, it would be a mistake to say that they are all alike (Figure 12).  We also find that in some cases, most notably that of Comet Kudo-Fujikawa, even though its silicate band is weak, the shape and structure after subtracting the continuum is similar to that of other dynamically new and long period comets. It also resembles the emission of more evolved pre-main sequence stars, such as HD 35187. For the same band strength at 11.2 $\\mu$m, the 9-10 $\\mu$m fluxes  of the JF comets tend to be slightly larger than that of Kudo-Fujikawa (and other dynamically new and long period comets), and are more similar to pre-main sequence stars that are less evolved than HD 35187. It would seem that the dynamically new and long period comets, which come from the Oort Cloud but whose origins were probably inside the orbit of Neptune, were endowed with grains of a somewhat different nature than the Jupiter family objects, which probably originated from the Kuiper-Edgeworth Belt beyond the orbit of Neptune. During the first tens of millions of years of the evolution of the solar system, these two regions of comet formation would have been processed over time to various degrees with solar wind exposure, silicate grain annealing, small grain expulsion, grain coagulation, and settling of grains in the protosolar disk. Alien observers (if they existed) might, at various times, have seen a system that progressively evolved from one that looked like UX Ori to HD 31648 to HD 35187."
    },
    "0404/astro-ph0404376_arXiv.txt": {
        "abstract": "{ We investigate the bright cut-off of the \\oiiiline\\ planetary nebula luminosity function (PNLF), that has been suggested as a powerful extragalactic distance indicator in virtue of its observed invariance against populations effects. Theoretical PNLFs are constructed via Monte-Carlo simulations of populations of PNe, whose individual properties are described with the aid of recent PN synthetic models (Marigo et al. 2001), coupled to a detailed photoionisation code (CLOUDY). The basic dependences of the cut-off magnitude $M^*$ are then discussed. We find that : (i) In galaxies with recent or ongoing star formation, the modelled PNLF present $M^*$ values between $-4$ and $-5$, depending on model details. These are very close to the observationally-calibrated value for the LMC. (ii) In these galaxies, the PNLF cut-off is produced by PNe with progenitor masses of about 2.5 \\Msun, while less massive stars give origin to fainter PNe. As a consequence $M^*$ is expected to depend strongly on the age of the last burst of star formation, dimming by as much as 5 magnitudes as we go from young to 10-Gyr old populations. (iii) Rather than on the initial metallicity of a stellar population, $M^*$ depends on the actual [O/H] of the observed PNe, a quantity that may differ significantly from the initial value (due to dredge-up episodes), especially in young and intermediate-age PN populations. (iv) Also the transition time from the end of AGB to the PN phase, and the nuclear-burning properties (i.e. H- or He-burning) of the central stars introduce non-negligible effects on $M^*$. The strongest indication derived from the present calculations is a serious difficulty to explain the age-invariance of the cut-off brightness over an extended interval, say from 1 to 13 Gyr, that observations of PNLFs in galaxies of late-to-early type seem to suggest. We discuss the implications of our findings, also in relation to other interpretative pictures proposed in the past literature. ",
        "introduction": "\\label{sec_intro}  The Planetary Nebula luminosity function (PNLF) constitutes one of the most attractive indicators to set out the extragalactic distance scale. {From} the observational side, PNe can be identified in galaxies from the LMC to as far as the Virgo cluster, and are numerous in some of the key galaxies that link the ``local'' distance scale given by Cepheids, with the more distant scale given by the surface brightness fluctuations (SBF) method, the Tully-Fisher relation, and SNe Ia. The PNLF is suitable for determining distances as far as 20--25~Mpc, beyond which exposure times for reasonable S/N-ratios become prohibitively long. The importance of the PNLF for the cosmological distance ladder is fully described by e.g. Ciardullo (2003), Ciardullo \\& Jacoby (1999) and Jacoby, Ciardullo \\& Feldmeier (1999), Jacoby (1997), to whom we refer for all details.  In practice, the PNLF as a candle works in the following way: Once the flux in the \\oiiiline\\ emission line, $F_{5007}$ (in cgs units), is measured for a sample of PNe in an external galaxy, it is converted to \\moiii\\ magnitude, \\beq \\moiii = -2.5\\,\\log F_{5007} - 13.74 \\label{eq_moiii} \\eeq (Jacoby 1989). The number of objects at any given \\moiii\\ bin constitutes the PNLF, $N(\\moiii)$, which can be fitted over the brightest magnitudes by the function \\beq N(\\moiii) = e^{0.307\\,m_{5007}} \\left[ 1-e^{3(m^*-m_{5007})} \\right] \\label{eq_cutoff} \\eeq (see Ciardullo et al.~1989). The key quantity is the cut-off magnitude $m^*$. According to several arguments (see below), its corresponding absolute value, $M^*$, appears to be constant in most galaxies, hence providing an excellent (secondary) standard candle. Therefore, a measure of $m^*$ gives a distance estimate in a single step, provided that interstellar reddening is suitably taken into account. The present-day calibration of $M^*$, based on a sample of galaxies with SBF and Cepheid distances, and correcting for a small metallicity dependence, is $M^* = -4.47$ (Ciardullo et al. 2002).  The possibility of systematic variations in the value of $M^*$ among galaxies has been investigated in several different ways. From the observational side, the evidences for a nearly constant maximum luminosity for PNe are numerous, as reviewed by e.g.\\ Ciardullo (2003) and Jacoby (1997). The main points are: (1) $m^*$ does not significantly change with galactocentric radius in the galaxies M31, M81, NGC\\,4494 and NGC\\,5128, where significant population gradients should be present. (2) There is no hint of large systematic variations in $m^*$ among galaxies of different types in the galaxy clusters/groups of M81, NGC\\,1023, NGC\\,5128, Fornax, Leo~I, and Virgo. Overall, this is indicating that age and metallicity variations have little effect on $M^*$. On the other hand, comparison with Cepheid distances reveals a small non-monotonic dependence on metallicity, expressed by the system's oxygen abundance, which can be compensated for by a relation quadratic in [O/H] (Ciardullo et al.\\ 2002).  Although these observational arguments seem compelling, they are not accompanied by a solid, unique theoretical framework. In other words, different authors do not agree on the reasons why $M^*$ should be constant, reaching seemingly contradictory conclusions over the years. First, early theoretical arguments by Jacoby (1989) and PN modelling by Dopita et al.\\ (1992) explain why the \\oiiiline\\ flux depends little on the PN metallicity. At the same time, the PNLF simulations by M\\'endez et al. (1993) indicate that a substantial dependence on population age might be present, as indicated  by the differences in the best fits of the LMC and Milky Way PNLFs. Dopita et al. (1992) also suggest that a dependence on age should be present. While M\\'endez et al.\\ (1993) claim that the bright end of the PNLF is populated by optically-thin objects, Dopita et al.\\ (1992) model optically-thick objects only and, after metallicity and population corrections, succeed in reproducing the PNLF and the constant cutoff brightness. Similar open questions concern the importance of hydrogen- (M\\'endez et al.\\ 1993) vs.\\ helium-burning (Ciardullo et al.\\ 1989; Jacoby 1989) central stars.  More recently, the apparent lack of predicted population age effects is explained by  Ciardullo \\& Jacoby (1999) by compensating circumstellar extinction. Young and massive central stars should be very luminous and result in over-luminous PNe (i.e. beyond the standard $M^*$), which is not observed. The proposed solution by Ciardullo \\& Jacoby (1999) is  that, since the most massive asymptotic giant branch (AGB) stars likely become dust-enshrouded because of  heavy mass loss, the related circumstellar extinction would reduce the  \\oiiiline\\ emission from these objects, always to at least  $M^*$ or below. Provided that PNe with lower-mass progenitors are present, then the observed cut-off magnitude would be $M^*$. If this is the case, one has yet to understand which is the typical age of the PNe that contribute to the cut-off, and explore the metallicity effects in the relevant age intervals.  The theoretical investigations of the PNLF, quoted above, have different levels of sophistication in the models. Because of the complexity of the problem, simplifications and severe assumptions about the properties of PNe and their progenitor stars have always been introduced, like for instance: an initial--final mass relation (hereinafter also IFMR) independent of metallicity (e.g. M\\'endez et al.\\ 1993), either optically-thin or thick PNe (Dopita et al.\\ 1992; M\\'endez et al.\\ 1993; Stanghellini 1995; M\\'endez \\& Soffner 1997), a narrow mass range being responsible for the bright end of the PNLF (Jacoby 1989; Dopita et al.\\ 1992), and the restriction to either H- or He-burner central stars only. All these assumptions can be improved in some way, by using the full results of stellar evolution and the population synthesis theory.  In the present paper, we elaborate on the theoretical framework for the PNLF predictions, based on our own simulations. The starting point for such a project has been the construction of an updated synthetic code, that follows the PN evolution as a function of basic stellar parameters, such as progenitor mass, metallicity, type of post-AGB evolution, etc. The complete code is described in Marigo et al. (2001; hereafter Paper~I), together with basic comparisons with the available data for Galactic PNe. Note that this code is much more complex than those previously used to investigate the PNLF, as it takes into account the chemical composition and wind velocity of the ejected material as predicted by AGB models, the interaction of the multiple winds as the PN expands, the effect of shell thickening due to ionisation, etc. Moreover, there is no artificial assumption about the conditions of optical thickness of the nebulae. The underlying model used to produce theoretical PNLFs is therefore more realistic and much less dependent on explicit or hidden assumptions than previous approaches. We will recapitulate the key features of our theoretical model in Section~\\ref{sec_overview}.  Despite all improvements, the Marigo et al.\\ (2001) models still treated some characteristics of the ionized nebulae in a simplified way. For instance, the electronic temperature \\Te\\ was introduced as a parameter and kept constant during the complete PN evolution. The emission line fluxes were then derived by simple formulas based on this temperature, on element abundances, and on the estimated sizes of the Str\\\"omgren spheres for different ions. Since the flux of the optical forbidden lines -- and in particular the \\oiiiline\\ one -- depends strongly  on \\Te\\ and the actual amount of ionized material, this simplification is clearly inadequate for an investigation of the PNLF. Hence, in this paper (Section~\\ref{ssec_cloudy}) we relax the above-mentioned approximations by incorporating the photoionisation code CLOUDY (Ferland 2001) into our PN models, thereby improving them further, and particularly towards the direction of accurate predictions of spectral features.  Then, in Sect.~\\ref{sec_samples} we describe the generation of simulated PN samples in galaxies, firstly performing some basic comparisons with the empirical data of Galactic PNe. We show how present models can reproduce, to a large extent, some well-known properties of observed PNe, such as the correlations between the ionized mass and nebular density with the radius, and those between particular emission line ratios.  In Sect.~\\ref{sec_pnlf} we move to analyse the predicted synthetic PNLF, with the aim to discuss the expected behaviour of the bright cut-off, $M^*$. We will in particular focus  on the nature of the objects responsible for the bright cut-off, analysing the dependence of $M^*$ on physical and technical factors.  Finally, Sects.~\\ref{sec_discussion} and \\ref{sec_conclusion} summarise the results of present models in terms of basic dependences of the cut-off, and draw a few conclusions about the PNLF method, at the same time highlighting the main points of discrepancy between our findings and the picture used so far to interpret the PNLF results. ",
        "conclusions": "\\label{sec_conclusion}  In Paper~I we had developed a theoretical model for synthetic PNe. In the present paper we use this model to predict properties of synthetic samples of PNe. In particular we investigated the PNLF of synthetic PN populations with regard to its shape, maximum \\oiiiline\\ brightness, and how it depends on age and composition of the population as well as on assumptions necessarily entering our model. With respect to Paper~I and in view of the importance of a reliable prediction of the emission properties of PNe, we extended our model by including the photoionisation program CLOUDY, thereby eliminating one of the free parameters, the electron temperature \\Te\\ in the nebulae. The detailed radiation transfer treatment confirms that the assumption of an approximately constant \\Te\\ of about $10000\\,\\rm{K}$ in the dynamical part of the model is justified, at least within the accuracy required at this stage. Apart from free assumptions about age and composition of the population causing PNe, and about the SFH and IMF, the transition time remains the sole physical parameter for which we have only rough indications, and for which we would need additional theoretical calculations of the superwind phase of AGB stars, and their transition to the post-AGB phase. We are also forced to assume spherical symmetry of the circumstellar matter.  With this improved model for PNe populations we have simulated galactic PNe and compared the results with observations. We concentrated on emission lines and ionisation states, and find excellent agreement without the necessity to vary any model parameters. Based on these successful comparisons, we then concentrated on the luminosity function. The main results concerning the PNLF are: \\begin{itemize}  \\item Our models predict a PNLF shape towards the cut-off very similar to those observed in spiral and irregular galaxies.  \\item In galaxies with recent star formation (either continued up to now, or with a burst occurred $\\sim 1$ Gyr ago), the cut-off is expected at $-5 \\le \\Moiii \\le -4$, depending on model parameters. This is consistent with the observed values at about $-4.5$ (e.g. for the LMC).  \\item In galaxies without recent star formation, the cut-off is predicted much fainter than observed. This is due to the lack of stars with $M>2$ \\Msun, responsible for the brightest PNe ($\\Moiii\\le-3.0$).  \\item  From the two previous points it follows that we can hardly explain the constancy of the cut-off observed in some ellipticals (dominated by old populations) in terms of the claimed age-invariance. Rather, a possible way to explain the bright PNLF cut-off observed in some late-type galaxies is just to invoke a recent burst of star formation.  \\item One  main parameter determining the PNLF cut-off is the transition time, to be constrained on the base of observed PN properties (e.g. $N_{\\rm e}$, $T_{\\rm e}$, ionized masses, radii, expansion velocities, etc). Our simulations indicate that, in order to recover the observed brightness of the cut-off,  the  CSPNe populating it (with masses around $0.70-0.75$ \\Msun) should have a \\ttr\\ of the order of few hundred years (typically $500-1000$ yr). On galaxies without recent SFR, the cut-off position becomes little sensitive to \\ttr.  \\item The predicted dependence of the PNLF on metallicity is the result of a number of factors, as we need to know (i) how the evolutionary properties of the stellar progenitors (i.e. the efficiency of mass loss on the AGB, the details of the initial--final mass relation) vary as a function of the initial chemical composition, (ii) the initial main-sequence abundance of oxygen, and iii) the actual oxygen abundance in PNe since it may differ significantly from the initial value due to dredge-up episodes during the TP-AGB phase. \\end{itemize}  Overall, {\\em our results explain in a very natural way the successful calibration of the PNLF using Cepheid distances} (Feldmeier et al.\\ 1997; Jacoby 1996). In the list given by Jacoby (1996), 7 out of 8 galaxies in the sample are classified as spiral or irregular galaxies, only NGC5253 is of E/S0 type. However, it is well documented that this galaxy has undergone a recent starbust (see, e.g.\\ Verma et al.\\ 2003) and thus, as for all the other objects star formation has taken place during the last billion year, and all PNLFs reach $M^*$. In fact, the argument can be made even stronger and more general: in case there is a Cepheid distance to a galaxy, there must be a young population and thus the PNLF will reach the canonical cutoff brightness.  Furthermore, {\\em our models do not support the idea of a population-invariant PNLF cut-off}. Indeed, the suggestions for an age-invariant cut-off poses a serious problem to the theory. In order to explain it, {we may need to explore the properties of very low-$Z$ PNe, or to invoke} a very different and difficult-to-imagine scenario for the evolution of PNe and their progenitor stars than here considered.  Considering all these aspects, we find it worth questioning if the indications for the cut-off invariance are really as solid as claimed. Until definitive evidence is not brought, and until there is not a solid theoretical framework supporting such evidence, we suggest that the PNLF should still be considered as a potentially useful probe of stellar populations in galaxies, rather than a standard candle."
    },
    "0404/astro-ph0404140_arXiv.txt": {
        "abstract": "Evidence suggests that gamma-ray burst (GRB) ejecta are likely magnetized, although the degree of magnetization is unknown. When such magnetized ejecta are decelerated by the ambient medium, the characteristics of the reverse shock emission are strongly influenced by the degree of magnetization. We derive a rigorous analytical solution for the relativistic 90$^{\\rm o}$ shocks under the ideal MHD condition. The solution is reduced to the Blandford-McKee hydrodynamical solution when the magnetization parameter $\\sigma$ approaches zero, and to the Kennel-Coroniti solution (which depends on $\\sigma$ only) when the shock upstream and downstream are ultra-relativistic with each other. Our generalized solution can be used to treat the more general cases, e.g. when the shock upstream and downstream are mildly relativistic with each other. We find that the suppression factor of the shock in the strong magnetic field regime is only mild as long as the shock upstream is relativistic with respect to the downstream, and it saturates in the high-$\\sigma$ regime. This indicates that generally strong relativistic shocks still exist in the high-$\\sigma$ limit. This can effectively convert kinetic energy into heat. The overall efficiency of converting ejecta energy into heat, however, decreases with increasing $\\sigma$, mainly because the fraction of the kinetic energy in the total energy decreases. We use the theory to study the reverse shock emission properties of arbitrarily magnetized ejecta in the GRB problem assuming a constant density of the circumburst medium. We study the shell-medium interaction in detail and categorize various critical radii for shell evolution. With typical GRB parameters, a reverse shock exists when $\\sigma$ is less than a few tens or a few hundreds. The shell evolution can be still categorized into the thick and thin shell regimes, but the separation between the two regimes now depends on the $\\sigma$ parameter and the thick shell regime greatly shrinks at high-$\\sigma$. The thin shell regime can be also categorized into two sub-regions depending on whether the shell starts to spread during the first shock crossing. The early optical afterglow lightcurves are calculated for GRBs with a wide range of $\\sigma$ value, with the main focus on the reverse shock component. We find that as $\\sigma$ increases from below the reverse shock emission level increases steadily until reaching a peak at $\\sigma \\siml 1$, then it decreases steadily when $\\sigma > 1$. At large $\\sigma$ values, the reverse shock peak is broadened in the thin shell regime because of the separation of the shock crossing radius and the deceleration radius. This novel feature can be regarded as a signature of high $\\sigma$.  The early afterglow data of GRB 990123 and GRB 021211 could be understood within the theoretical framework developed in this paper, with the inferred $\\sigma$ value $\\simg 0.1$. The case of GRB 021004 and GRB 030418 may be also interpreted with higher $\\sigma$ values, although more detailed modeling is needed. Early tight optical upper limits could be interpreted as very high $\\sigma$ cases, in which a reverse shock does not exist or very weak. Our model predictions could be further tested against future abundant early afterglow data collected by the Swift UV-optical telescope, so that the magnetic content of GRB fireballs can be diagnosed. ",
        "introduction": "Extensive broad-band observational campaigns and theoretical modeling of gamma-ray burst (GRB) afterglows have greatly advanced our understanding of these mysterious cosmic explosions. Yet, the origin of the GRB prompt emission itself and the nature of the relativistic flow (which are directly connected to the function of the central engine) is still unknown (e.g. \\Mesz~2002; Zhang \\& \\Mesz~2004). In particular, it is unclear how important the role of magnetic fields is in producing GRBs. Recently two independent pieces of evidence suggest that the GRB central engine is likely strongly magnetized. First, the claimed detection of the very high degree of linear polarization of gamma-ray emission in GRB 021206 (Coburn \\& Boggs 2003, see however Rutledge \\& Fox 2004), if true, could be readily interpreted by assuming that the magnetic field involved in the synchrotron radiation is globally ordered (e.g. Waxman 2003; Granot 2003), although some alternative explanations remain (e.g. Waxman 2003). Second, recently we (Zhang, Kobayashi \\& \\Mesz~2003, hereafter ZKM03) developed a method to perform combined reverse and forward shock emission study for GRB early optical afterglows, and revealed that a stronger magnetic field in the reverse shock region than in the forward shock region is needed to interpret the early afterglow data of GRB 990123 and GRB 021211. This claim was confirmed by independent detailed case studies for both bursts (Fan et al. 2002; Kumar \\& Panaitescu 2003). These findings suggest that magnetic fields may play a significant role in the GRB physics, as has been suggested by various authors previously (e.g. Usov 1994; Thompson 1994; \\Mesz~\\& Rees 1997b; Wheeler et al. 2000; Spruit, Daigne \\& Drenkhahn 2001; Blandford 2002). Within the framework of the currently favored collapsar progenitor model for GRBs (MacFadyen \\& Woosley 1999), the ejecta are found to be magnetized when MHD simulations are performed (Proga et al. 2003).  The degree of magnetization of the ejecta, however, is unknown. This is usually quantified by the parameter $\\sigma$ (see eq.[\\ref{sigma}] for a precise definition), the ratio of the electromagnetic energy flux to the kinetic energy flux. Current GRB models are focused on two extreme regimes. In the first regime, it is essentially assumed that the GRB fireball is purely hydrodynamical. Magnetic fields are introduced only through an equipartition parameter $\\epsilon_B$ (which is of the order of 0.001-0.1) for the purpose of calculating synchrotron radiation. This is the $\\sigma \\rightarrow 0$ regime. In this picture, the GRB prompt emission is produced from internal shocks (Rees \\& \\Mesz~1994) or sometimes from external shocks (\\Mesz~\\& Rees 1993; Dermer \\& Mitman 1999). This is currently the standard scenario of GRB emission. The second is the $\\sigma \\rightarrow \\infty$ regime. This is the regime where a Poynting-flux dominates the flow, and GRB prompt emission is envisaged to be due to some less familiar magnetic dissipation processes (e.g. Usov 1994; Spruit et al. 2001; Blandford 2002; Lyutikov \\& Blandford 2003). In principle, a GRB event could include both a ``hot component'' as invoked in the $\\sigma=0$ model (e.g. due to neutrino annihilation) and a ``cold component'' as invoked in the $\\sigma=\\infty$ model, the interplay between both components may result in a $\\sigma$ value varying in a wide range (Zhang \\& \\Mesz~2002). It is an important but difficult task to pin down the degree of magnetization of GRB ejecta.  GRB early afterglow data (especially in the optical band) potentially contain essential information to diagnose the magnetic content of the fireball. The reason is that an early optical afterglow lightcurve is believed to include contributions from  both the forward shock (which propagates into the ambient medium) and the reverse shock (which propagates into the ejecta itself). Since the magnetization degree of the ejecta influences the emission level of the reverse shock (or maybe even the level of the forward shock), by studying the interplay between the reverse shock and the forward shock emission components, one could potentially infer the degree of magnetization of the ejecta. In all the current analyses, the reverse shock emission is treated purely hydrodynamically (e.g. \\Mesz~\\& Rees 1997a; Sari \\& Piran 1999; Kobayashi 2000; Kobayashi \\& Zhang 2003a; ZKM03). When confronted with the available early afterglow data (four cases so far: GRB 990123, Akerlof et al. 1999; GRB 021004, Fox et al. 2003a; GRB 021211, Fox et al. 2003b, Li et al. 2003a; and GRB 030418, Rykoff et al. 2004), the model works reasonably well for two of them (GRB 990123 and GRB 021211), although a good fit requires that the magnetic field in the reverse shock region is much stronger than that in the forward shock region (ZKM03). For the other two, the lightcurves are not easy to explain with the simplest reverse shock model. On the other hand, GRB ejecta could in principle have an arbitrary $\\sigma$ value. When $\\sigma$ is large, the conventional hydrodynamical treatment is no longer a good approximation, and a full treatment involving MHD shock jump condition is desirable.  It is generally believed that a GRB involves a rapidly-rotating central engine. If the magnetic dissipation processes are not significant, field lines are essentially frozen in the expanding shells. The radial component of the magnetic field decays with radius as $\\propto R^{-2}$, while the toroidal magnetic field decays as $\\sim R^{-1}$. At the external shock radius, magnetic field lines are essentially frozen in the plane perpendicular to the moving direction. The MHD shock Rankine-Hugoniot relations are greatly simplified in such a 90$^{\\rm o}$ shock. Such relations have been studied extensively before both analytically and numerically. Kennel \\& Coroniti (1984) derived some simplified analytical expressions applicable for strong 90$^{\\rm o}$ degree shocks whose upstream and downstream are ultra-relativistic with each other. The model was used to treat the pulsar-wind nebula problem. In this regime, the strength of the shock is essentially characterized by only one parameter, i.e. the $\\sigma$ parameter. The conclusion was confirmed later by numerical simulations (e.g. Gallant et al. 1992)  Within the context of GRBs, since a GRB invokes a transient release of energy, the ejecta shell has a finite width (in contrast to the long-standing pulsar wind). Under some conditions, the reverse shock upstream and downstream could never become relativistic with each other when the reverse shock crosses the ejecta shell. In the $\\sigma=0$ limit, whether the reverse shock becomes relativistic depends on the comparison between the time scale ($T$) of the central engine activity (essentially the duration of the burst) and the time scale ($t_\\gamma$) when the mass of the ambient medium collected by the fireball reaches $1/\\gamma_0$ times the mass of the ejecta (e.g. Sari \\& Piran 1995; Kobayashi, Piran \\& Sari 1999). Both times are measured by the observer. The case of $T>t_\\gamma$ is called the thick shell regime, and the reverse shock is relativistic. In many cases, however, one has $T<t_\\gamma$. i.e. the thin shell regime\\footnote{If a GRB contains several well-separated emission episodes, the whole burst may be even separated into several discrete shells. In such cases, even a long-duration burst may be treated as the superposition of several thin shells rather than one single thick shell. See Zhang \\& \\Mesz~2004 for more discussions.}. The reverse shock is initially non-relativistic and only becomes mildly relativistic as the shock crosses the shell. For magnetized ejecta (e.g. a shell with a finite width but an arbitrary $\\sigma$ value), the separation between the thick and thin shell regimes becomes more complicated, but the non-relativistic reverse shock case (for the thin shell regime) is even more common (see \\S\\ref{sec:regimes}). The Kennel-Coroniti approximation can not be directly used. The theory developed in this paper becomes essential to discuss the reverse shock physics in this parameter regime.  In this paper, we present a detailed treatment of reverse shock emission for an arbitrarily magnetized ejecta under the ideal MHD condition. The reverse shock emission in the mildly magnetized regime is also discussed by Fan, Wei \\& Wang (2004a) recently. Here we will develop a theoretical framework to include discussions for the reverse shock emission in a wider $\\sigma$ range, as well as for various ejecta-medium interaction parameter regimes. We first (\\S2) present a rigorous analytical solution for the MHD 90$^{\\rm o}$ shock jump conditions, which is applicable for an arbitrary $\\sigma$ value and for an arbitrary Lorentz factor $\\gamma_{21}$ between the upstream and the downstream. The detailed derivation and the relevant equations are presented in the Appendix A. We then (\\S3) discuss the ejecta - medium interaction within the context of GRB fireball deceleration and re-investigate the critical fireball radii and re-categorize the thick vs. thin shell regimes. This leads to a more complicated picture than in the pure hydrodynamical case (Sari \\& Piran 1995). In \\S4, we calculate the synchrotron emission from the shocks under the conventional assumptions about the particle acceleration in collisionless shocks, and present the predicted GRB early optical lightcurves for a wide range of $\\sigma$ value. We discuss how the early afterglow data may be used to diagnose the magnetic content of GRB ejecta. Our results are summarized in \\S5 with some discussions. ",
        "conclusions": "We have derived a rigorous analytical solution for the relativistic 90$^{\\rm o}$ shocks under the ideal MHD condition (eq.[\\ref{eq}]). Generally, the solution depends both on the magnetization $\\sigma$ parameter and the Lorentz factor of the shock, $\\gamma_{12}$. The solution can be reduced to the Blandford-McKee hydrodynamical solution when $\\sigma=0$, and to the Kennel-Coroniti solution (which depends on $\\sigma$ only) when the $\\gamma_{21} \\gg 1$. Our generalized solution can be used to treat the more general cases, e.g. when the reverse shock upstream and downstream are mildly relativistic with each other. Since GRBs invoke a shell with finite width, this latter possibility is common (e.g. the parameter space for thin shell greatly increases in the high-$\\sigma$ regime), so that our generalized solution is essential to deal with the GRB reverse shock problem.  Several interesting conclusions emerge from our analysis. (1) Strong shocks still exist in the high-$\\sigma$ regime, as long as the shock is relativistic. Figs.\\ref{fig:sig}c,d indicate that as $\\sigma$ increases, both the downstream ``temperature'' $e_2/n_2$ and the ``shock compression factor'' $n_2/n_1$ decreases with respect to the $\\sigma=0$ values. However, the suppression factors in both cases are only mild (a factor of $\\sim$ 0.5), and they saturate when $\\sigma \\gg 1$. In the relativistic shock regime, the results are actually consistent with Kennel \\& Coroniti (1984). However, these authors did not calculate the suppression factor with respect to the $\\sigma=0$ case, and did not explore further into the high-$\\sigma$ regime, so that their results leave the impression to the readers that the shock is completely suppressed when $\\sigma$ reaches higher values. For typical GRB parameters, we found that a reverse shock still exists when $\\sigma$ is as high as several tens or even hundreds. When the reverse shock exists, its emission level decreases when $\\sigma$ gets higher. This is not only because the reverse shock becomes weaker since $\\gamma_{34}$ gets close to unity in the high-$\\sigma$ regime, but also because the total kinetic energy in the flow (which is the energy reservoir for shock dissipation) gets smaller given a same total energy. (2) When discussing ejecta-medium interaction, somewhat surprisingly, some important parameters, such as $F$ and $C_\\Delta$, are very insensitive to the reverse shock Lorentz factor, $\\gamma_{34}$, and can be regarded as the function of $\\sigma$ only (\\S\\ref{sec:general}, \\S\\ref{sec:regimes}). This greatly simplify the problem, and is essential to characterize the parameter regimes. (3) The triple coincidence of the first three critical lines in eq.(\\ref{lines}) is very crucial for a self-consistent description of the problem (\\S\\ref{sec:regimes}).  Comparing with the conventional hydrodynamical treatment, we reveal several novel features for the early lightcurves. First, as $\\sigma$ increases, the reverse shock peak flux increases rapidly initially, reaching a peak around $\\sigma \\sim (0.1-1)$, and starts to decrease when $\\sigma \\geq 1$. Second, due to the inability of tapping the Poynting flux energy during the shock crossing process, the fireball deceleration radius for the thin shell case decreases as $\\sigma$ increases $[\\propto (1+\\sigma)^{-1/3}]$. The forward shock emission level is also lower right after shock crossing. Third, in the high-$\\sigma$ thin shell regime, the reverse shock peak is broadened due to the separation of the shock crossing radius and the deceleration radius. This is a signature for a high-$\\sigma$ flow, which can be used to diagnose the magnetic content of the fireball. Fourth, as $\\sigma$ becomes large enough (larger than several tens or several hundreds), the condition for forming a reverse shock is no longer satisfied, and there should be no reverse shock component in the early afterglow lightcurves. This could be consistent with very early tight optical upper limits for some GRBs, such as the recent Swift dark burst GRB 050319a (Roming et al. 2005). In summary, the above new features allows the current theory to potentially interpret known GRB early afterglow cases collected so far as well as the case of the dark bursts, if one allows $\\sigma$ to vary in a wide enough range (say, from $0.01$ to $100$). The Swift GRB mission, launched on November 20, 2004, is expected to detect many early optical afterglow lightcurves with the UV-optical telescope on board. We expect to further test our theoretical predictions against the abundant Swift data, and to systematically diagnose the magnetic content of GRB fireballs.  If the GRB ejecta is indeed magnetized, as inferred from the early afterglow data, the internal shocks should also be corrected by the magnetic suppression factor. This aspect has been investigated by Fan, Wei \\& Zhang (2004b) recently.  Throughout the paper, we have treated the problem under the ideal MHD limit. In the high-$\\sigma$ case, strong magnetic dissipation may occur. The magnetic dissipation effect has been included in the internal shock study of Fan et al. (2004b). Our treatment in this paper presents a first order picture to the early afterglow problem (see also Fan et al. 2004a, whose treatment in the mildly magnetized regime is consistent with ours), and further considerations are needed to fully delineate the physics involved. Also our whole discussion is relevant when a reverse shock is present. It does not apply to the regime for an even higher $\\sigma$ value (e.g. Lyutikov \\& Blandford 2003). Finally, as discussed in \\S4.1, further studies on particle acceleration in MHD shocks are essential to give a more accurate calculation on the reverse shock emission in the high-$\\sigma$ regime.  We only discussed one type of the medium, i.e., assuming a constant medium density, typically for the interstellar medium. In principle, the medium density can vary with distance from the central engine. In particular, a wind-type medium, characterized by the $n \\propto R^{-2}$ profile, has been widely discussed. Our MHD shock theory could be straightforwardly used to the wind case to study the reverse shock emission in combination with the previous pure hydrodynamical treatments (Chevalier \\& Li 2000; Wu et al. 2003; Kobayashi \\& Zhang 2003b; Kobayashi, \\Mesz~\\& Zhang 2004, see Fan et al. 2004a for a preliminary treatment)."
    },
    "0404/astro-ph0404189_arXiv.txt": {
        "abstract": "Recent observations of globular clusters (GCs) in intermediate-age (2--4 Gyr old), early-type merger remnants have provided the hitherto `missing link' between young merger remnants and `normal' elliptical galaxies in the form of a GC subsystem with colors and luminosities consistent with population synthesis model predictions for those ages and $\\sim$\\,solar metallicity. Here we present new, deep observations of the GC system of the intermediate-age merger remnant NGC~1316, using the {\\it ACS} camera aboard {\\it Hubble Space Telescope}, which allowed us to create luminosity functions (LFs) {\\it as a function of galactocentric radius}. We find that the inner 50\\% of the `red' GC system shows a clear turnover in its LF, at about 1 mag fainter than that of the `old' blue GCs. This constitutes direct, dynamical evidence that metal-rich GC populations formed during a gas-rich merger can evolve into the `red', metal-rich GC populations that are ubiquitous in `normal' giant ellipticals. ",
        "introduction": "Over the last decade, the importance of young massive star clusters as signposts to major star formation events has been recognized. In particular, the discovery of young globular clusters (GCs) in galactic mergers and young merger remnants using {\\it Hubble Space Telescope \\/} (HST) imaging (e.g., Schweizer 2002 and references therein) has generated a great deal of excitement. This is hardly surprising, as these clusters have provided previously unexpected opportunities including {\\it (i)\\/} the chance to study the formation and evolution of GCs in the local universe rather than trying to derive how they formed $\\sim$\\,14 Gyr ago; {\\it (ii)\\/} the possibility to age-date merger events (i.e., major star formation epochs) in galaxies from star clusters with simple, single-burst stellar populations.  In many cases, follow-up spectroscopy of these young clusters confirmed their nature as star cluster and their ages, and in one case even their high masses predicted from their colors and luminosities (Maraston et al.\\ 2004). The metallicities of the young clusters have been found to be near solar, as expected for clusters formed out of enriched gas in spiral disks. Such metal-rich clusters are now known in merger remnants at essentially all ages, ranging from birth (e.g., NGC~4038/4039, Whitmore et al.\\ 1999) through youth (several 10$^8$ yrs, e.g.\\ NGC 7252, Schweizer \\& Seitzer 1998) to middle age, when only faint ripples and loops reveal their merger history (e.g. NGC~1316, Goudfrooij et al.\\ 2001a).  It should also be emphasized that the cluster systems of young merger remnants were also found to contain fainter, redder objects with colors and luminosities consistent with those of old, metal-poor halo GCs that are very common in nearby galaxies such as the Milky Way. It is commonly believed that these fainter GCs in merger remnants probably belonged to the (now merged) progenitor spirals. This bimodality in star cluster systems of young merger remnants is rather interesting in the context of formation scenarios for {\\it normal\\/} elliptical galaxies, since the latter are also well known to generally host GC systems with bimodal color distributions % (Kundu \\& Whitmore 2001; Larsen et al.\\ 2001 to name a few). Deep {\\it HST\\/} imaging of `normal' ellipticals shows that roughly half of their GCs are blue (metal-poor) and half are red (metal-rich, typically roughly solar or somewhat less). Spectroscopy with 10-m class telescopes revealed that both `blue' and `red' GC subpopulations of `normal' ellipticals are typically old ($\\ga$\\,8 Gyr, % Cohen, Blakeslee \\& C\\^ot\\'e 2003; Puzia et al.\\ 2004). A natural interpretation of these data is that the metal-rich GCs in `normal' ellipticals formed in gas-rich mergers at $z \\ga 1$, and that the formation process of galaxies with significant populations of metal-rich GCs was similar to that in galaxy mergers observed today.  One might think ``now don't get carried away'' at this point, and indeed it is important to examine this scenario in detail. If correct, one {\\it should\\/} be able to {\\it (i)\\/} find ellipticals with second-generation GCs of intermediate age (i.e., 2\\,--\\,5 Gyr) and {\\it (ii)\\/} study the evolution of second-generation GC systems from young through intermediate to old ages, and compare their properties with theoretical predictions. These two issues are addressed in the remainder of this paper. ",
        "conclusions": ""
    },
    "0404/astro-ph0404230_arXiv.txt": {
        "abstract": "s{The r--process involves neutron-rich nuclei far off stability for which no experimental cross sections are known. Therefore, one has to rely on theory. The difficulties in the predictions are briefly addressed. To investigate the impact of altered rates, a comparison of r--process production in hot bubble models with largely varied rates is shown. Due to the (n,$\\gamma$)-($\\gamma$,n) equilibrium established at the onset of the r-process, only late-time neutron captures are important which mainly modify the abundances around the third r--process peak.} ",
        "introduction": "Nucleosynthesis of elements beyond the iron peak requires reactions with neutrons due to the high Coulomb barriers which prevent charged particle reactions. Except for the relatively underabundant proton-rich p--nuclei, two processes have been identified for the production of intermediate and heavy nuclei: the slow neutron-capture process (s--process) and the rapid neutron-capture process (r--process). With neutron number densities around 10$^8$ cm$^{-3}$ and low effective neutron energies of around 30 keV, the s--process synthesizes nuclei along the line of stability as the neutron captures are generally slower than all beta-decays encountered along its path (with the exception of several branching points where the two timescales become similar). Approximately half of the intermediate and heavy elements are created in the much faster r--process with neutron number densities exceeding 10$^{22}$ cm$^{-3}$, effective neutron energies around 100 keV, and much shorter process times of up to a few seconds. These conditions point to an explosive site but the actual site has yet to be identified. The long favored idea of a high-entropy bubble in the neutrino wind ejected from a type II supernova shows persistent problems in explaining production across the full mass range of r--nuclei. Furthermore, there are indications that there must be two distinct sites ejecting r--process material at different frequencies (see other contributions in this volume). In consequence, most r--process investigations focus on simplified, parameterized models which allow to study the required conditions and their sensitivities to nuclear inputs.  Due to the high neutron densities the r--process synthesizes very neutron-rich nuclei far off stability which subsequently decay to stability when the process ceases due to lack of neutrons or low temperatures. This raises the question whether we can predict reactions far off stability sufficiently well to make statements about r--process conditions. In the following two main topics are briefly addressed: The difficulties in predicting neutron captures far off stability, and the impact of neutron captures on the resulting r--process abundances. ",
        "conclusions": "The simple comparison shown above for the hot bubble model has to be interpreted cautiously. Despite the fact that there might be considerable uncertainties in the theoretical rates far off stability changing {\\it all} rates in a range of 4 orders of magnitude seems unrealistic. Even if new effects (like pygmy resonances [7] or overestimated cross sections [3]) might change the rates by factors of 10 for extremely neutron-rich nuclei, late-time captures will not include such nuclei but will occur closer to stability. Moreover, for reproducing the solar r--process pattern it is necessary to superpose a number of components with different entropies. Thus, effects of rates altered on a large scale, as shown above, can be compensated by a scaling in entropy and a different weight distribution. Thirdly, a more realistic seed abundance distribution might also decrease the difference in heavy element production between the different cases. Higher entropies realistically start with seed abundances in the $A\\approx 110$ region and require less neutrons to form more heavy elements. However, this was not taken into account here to purely show the influence of altered neutron captures.  Despite the above caveats the main conclusions are consistent with other studies [5,8]. Components with high entropy freeze out slower and late-time neutron captures can modify the final abundance distribution mainly in the region $A>140$. Therefore, emphasis has to be put on improving the prediction of nuclear cross sections and astrophysical reaction rates in that mass region."
    },
    "0404/astro-ph0404006_arXiv.txt": {
        "abstract": "The observed infall of galaxies into the Virgo Cluster puts strong constraints on the mass of the cluster.  A non-parametric fully non-linear description of the infall can be made with orbit reconstructions based on Numerical Action Methods.  The mass of the cluster is determined to be $1.2 \\times 10^{15} M_{\\odot}$.  The mass-to-light ratio for the cluster is found to be seven times higher than the mean ratio found across the region within $V=3000$ km/s. ",
        "introduction": "The general problem we will consider is the reconstruction of orbits that galaxies might have followed to arrive at their currently observed positions on the sky and in redshift.  On this occasion we will focus on the particularly interesting circumstances associated with infall into the Virgo Cluster. The program involves three distinct components.  First, we require as complete a map as possible of the angular positions and velocities of galaxies.  It will be assumed that there is some correlation between the distribution of galaxies in the catalog and the actual distribution of mass, which is the parameter that really interests us.  However it is to be anticipated from the outset that the relationship between what we see -- galaxies of various types -- and the distribution of mass may be complex (Tully 2003, 2004).  It is also to be appreciated that large scale tides may be dynamically significant so one needs a map of the distribution of galaxies that extends well beyond the region of immediate focus.  In the present instance, our interest is in the environs of the Virgo Cluster at a distance of 16.8 Mpc.  Our map of the distribution of galaxies is based on a sample of 3151 galaxies within $\\sim 40$ Mpc.  The sample is complete to $0.1 L^{\\star}$ within 25 Mpc at high galactic latitudes. Selection function corrections are made as a function of distance and `fake' sources are added at low latitudes to account for missing sources in the zone of obscuration.  The second critical component is a catalog of accurate distances to galaxies. Distance measures, $d$, allow a separation of observed velocities, $V_{\\rm obs}$, and peculiar velocities, $V_{\\rm pec}$, since $V_{\\rm pec}=V_{\\rm obs}-{\\rm H}_0 d$. Here, H$_0$ is the Hubble Constant which is taken to be 80 km/s/Mpc. For purposes of separating $V_{\\rm pec}$ from $V_{\\rm obs}$ it is only required that distances and the Hubble Constant be on the same scale; that is, with averaging that takes into account large scale flows: H$_0=<V_{\\rm obs}/d>$.  If distances are, say, systematically measured too large then a self-consistent value of H$_0$ will be too low.  Derived values of $V_{\\rm pec}$ are independent of the distance scale zero-point. Another factor with regard to distances to appreciate is that errors are a percentage of the distance, 10\\% to 20\\% depending on the methodology, with the consequence that errors in the derived $V_{\\rm pec}$ grow linearly with distance.  At the Virgo Cluster a 10\\% error corresponds to $\\pm140$ km/s in $V_{\\rm pec}$. This uncertainty is tolerable for the Virgo infall problem but for more distant clusters the situation would be unsatisfactory.  The distance catalog itself is a synthesis of our own observations (Pierce and Tully, in preparation) that exploit the luminosity--linewidth method (Tully and Fisher 1977) and material from the literature.  That literature material provides distances from a variety of methodologies: luminosity--linewidth (Mathewson et al. 1992, Lu et al. 1993, Tully and Pierce 2000), Cepheid (Freedman et al. 2001), Tip of the Giant Branch (Karachentsev et al. 2003), Planetary Nebula Luminosity Function (Jacoby et al. 1990), and Surface Brightness Fluctuation (Tonry et al. 2001).  This latter important source is distinguished separately in the ensuing discussion.  In total we have distances to almost 900 galaxies from the luminosity--linewidth and other assorted methods and 292 distances from the Surface Brightness Fluctuation method.  The third ingredient is the theoretical machinery to convert information on the amplitude of the peculiar velocities of galaxies into a map of the mass distribution.  In linear theory there is a direct relation between $V_{\\rm pec}$ and matter density fluctuations.  The relationship is more complex in the vicinity of collapsed structures, such as the Virgo Cluster. We make use of Numerical Action Methods which we have referred to as `Least Action'.  The numerical techniques have been discussed by Peebles (1989,1995), Shaya et al. (1995), and Phelps (2002).  The procedure allows for orbit reconstructions in highly nonlinear regimes, though ambiguity arises if the orbits are complex, especially if dynamical friction or orbital energy exchange is important.  These are not serious issues for the Virgo infall problem. ",
        "conclusions": "Similar analyses can be made of dozens of other lines-of-sight through the infall domain around the Virgo Cluster.  If there is a distance measurement that confirms that a target is inside the zero-velocity surface around the cluster then the line-of-sight component of its velocity puts a constraint on the cluster mass.  In each case, the cluster mass must be at least enough to provide an Action infall solution.  We determine that a minimum but sufficient Virgo Cluster mass for the models is $1.2 \\times 10^{15} M_{\\odot}$.  This mass is $~50\\%$ larger than the mass determined from application of the Virial theorem (Tully and Shaya 1984). However, the Virial radius for the cluster is $\\sim 0.8$ Mpc while the mass measured by infall is on a scale $\\ge 1.8$ Mpc.  The infall analysis, which is based on very simple physics, gives an estimate of the {\\it global} mass of the cluster.  It can be anticipated that this mass will exceed the masses determined over more limited scales by gravitational lensing, X-ray, or Virial studies.  Unfortunately, the analysis is not easily duplicated on other clusters, at least not so cleanly.  The Virgo Cluster is the only environment {\\it near} enough that distance measurements distinguish between infall and expansion regimes, and {\\it massive} enough to have an extensive, well populated infall domain.  It provides a single good case.  What we find, though, is that whereas overall the Action models find preference for $M/L \\sim 125 M_{\\odot}/L_{\\odot}$ and $\\Omega_m \\sim 0.2$, in the Virgo Cluster we need $M/L \\sim 900 M_{\\odot}/L_{\\odot}$."
    },
    "0404/astro-ph0404012_arXiv.txt": {
        "abstract": "We present a sequence of I-band images obtained at the Venezuela 1m Schmidt telescope during the outburst of the nebula recently discovered by J.W. McNeil in the Orion L1630 molecular cloud. We derive photometry spanning the pre-outburst state and the brightening itself, a unique record including 14 epochs and spanning a time scale of $\\sim 5$ years. We constrain the beginning of the outburst at some time between Oct. 28 and Nov. 15, 2003. The light curve of the object at the vertex of the nebula, the likely exciting source of the outburst, reveals that it has brightened $\\sim 5$ magnitudes in about 4 months. The time scale for the nebula to develop is consistent with the light travel time, indicating that we are observing light from the central source scattered by the ambient cloud into the line of sight. We also show recent FLWO optical spectroscopy of the exciting source and of the nearby HH 22. The spectrum of the source is highly reddened; in contrast, the spectrum of HH 22 shows a shock spectrum superimposed on a continuum, most likely due to reflected light from the exciting source reaching the HH object through a much less reddened path. The blue portion of this spectrum is consistent with an early B spectral type, similar to the early outburst spectrum of the FU Ori variable V1057 Cyg; we estimate a luminosity of ${\\rm L \\sim 219 L_\\sun}$. The eruptive behavior of the McNeil nebula source, its spectroscopic characteristics and luminosity, suggest we may be witnessing an FU Ori event on its way to maximum. Further monitoring of this object will decide whether it qualifies as a member of this rare class of objects. ",
        "introduction": "Last February 9, 2004, the discovery of a new nebula, roughly 12' south of the reflection nebula NGC 2068 (M78) in the L1630 molecular cloud of the Orion star-forming complex was announced by \\cite{mcn04}. The region surrounding the newly revealed object contains a number of pre-main sequence stars, as well as Herbig-Haro objects HH 22 and HH 23 \\citep{eim97}. The IRAS source 05436-0007 \\citep{cla91} is located at the vertex of the new nebula, but was not identified by him as a young object. The 2MASS images show an extremely red source at this position (J05461313-0006048), but no object is visible at optical wavelengths in Palomar Observatory Sky Survey plates. More recently, dust continuum imaging of this region by \\citet{lmz99} revealed several sources at 350 and $1300\\mu$m. They suggest that their source LMZ 12, spatially coincident with IRAS 05436-0007, is a Class 0 source with $L_{bol}\\simeq 2.7 L_\\sun$, and the probable exciting source of HH 23, located $\\sim 2\\farcm 5$ north of the new nebula.  This new nebula was made widely known by B. Reipurth's announcement in the Star Formation Newsletter No. 136\\footnote{http://www.ifa.hawaii.edu/$\\sim$reipurth/newsletter.htm}. In a follow up study \\citet{rea04} indicate that the object brightened by $\\sim 3$ magnitudes in the near infrared, and that its optical and K-band spectra show a certain resemblance with EXors. On the other hand, the object may be undergoing an FU Ori type eruption, which would imply a larger variation in brightness and a longer period at maximum light (Herbig 1977). To discriminate between these possibilities both pre-outburst and post-outburst observations are needed.  We report here optical images, photometry and spectroscopy of the McNeil Nebula, previous to its brightening, and during the outburst itself. These observations span $\\sim$ 5 years, with a particularly good coverage of the time when the outburst started. The observations were obtained during a sensitive optical variability study spanning most of the Orion OB1 Association (Brice\\~no et al. 2001). ",
        "conclusions": "The behavior of the McNeil nebula is reminiscent of the early evolution of the outbursts of FU Ori and V1057 Cyg (Herbig 1977). The optical (red) rise of McNeil, $\\sim 5$ mag in about a third of a year, is not much different than the early rise of these two objects, prototypical of the FUor class; moreover, the rise time of another FUor, V1515 Cyg, was much longer, indicating that the rate of brightening is not a universal quality among these objects. Like the McNeil object, V1057 Cyg initially showed a P Cygni profile on strong H$\\alpha$ emission (see inset in Figure 4), evolving into stronger (blue-shifted) absorption later on. Strong absorption in the upper Balmer lines was also present in spectra of V1057 Cyg, classified as B3 by \\citep{wel71}. However, the lack of He I lines lead Herbig to assign an approximate type of early A (Herbig 1977).  Herbig also noted the presence of wide emission at the Ca II resonance lines, which lead to a very weak Ca II 3933 {\\AA} absorption; this is also consistent with our optical spectra.  The eruptive behavior suggests that this may be an FU Ori object on its way to maximum, or possibly a member of the so-called class of EXors (Herbig 1977), which exhibit similar or smaller increases in optical brightness but last shorter periods of time. Continued monitoring should better help distinguish between these possibilities. Given the small number of objects in the FUor and EXor classes, it may be that there exists a continuum of outburst behavior which will be filled in by the discovery of additional objects. We note that at its recent brightness of ${\\rm I_c} = 14.4$, assuming ${\\rm A_I} = 7.2$, a distance of 400 pc, and adopting an A0 spectral type as a compromise between the range of spectral types we inferred before, the system luminosity is $L \\sim 219 {\\rm L_\\sun}$ (and would be higher if the intrinsic spectrum is earlier); this large luminosity is more similar to FUors than EXors.  The eruptive mechanism of FUors, and likely also that of EXors, is thought to be rapid outbursts in disk accretion (Hartmann \\& Kenyon 1996), possibly driven by the pile up of material in a disk due to rapid protostellar envelope infall (Kenyon \\& Hartmann 1991). The estimated bolometric luminosity  $L_{bol}\\simeq 2.7 L_\\sun$ and a SED like that of a Class 0 source \\citep{lmz99} is consistent with the pre-outburst object being a low-mass protostar. This object looks more similar to FU Oris, which are often highly embedded (Hartmann \\& Kenyon 1996) than EXors, which seem to be much less extincted.  This very interesting source warrants further monitoring to help determine its true nature."
    },
    "0404/astro-ph0404538_arXiv.txt": {
        "abstract": "The combination of ever more precise radio timing data and serendipitous  {\\it HST} data has confirmed that the outer companion to PSR B1620$-$26 is a planet. Here we summarize the observational situation, including preliminary new timing solutions and the implications of the measured system parameters. We detail the proposed formation scenarios, discussing the advantages and problems of each for explaining the origin of the triple, and we speculate on some of the implications for planet formation in the early universe. Future data on this system will provide additional constraints on fundamental modes of planet formation. We predict that many more exchanged planets will be discovered orbiting recycled pulsars in globular clusters as the sensitivity and duration of radio timing increases. Strong observational tests of some of the alternative formation models should be possible with additional data. ",
        "introduction": "The long series of radio timing data of the triple pulsar PSR~B1620$-$26 (Lyne et al.\\ 1988) has now confirmed that the second companion (Backer et al.\\ 1993), is a $1-3 \\, M_J$ substellar object---a planet---in a low eccentricity, wide circumbinary orbit about the inner pulsar--white dwarf binary. The outer orbit is significantly inclined to the inner orbital plane, and HST photometric data confirms that the white dwarf is young, and is a proper motion member of the cluster (Thorsett et al.\\ 1999; Sigurdsson et al.\\ 2003).  The mass of the object is very well constrained. The timing data puts a firm {\\it lower} mass limit of a jupiter mass; as more data has come in, the {\\it upper bound} on the mass has steadily shrunk. When timing first hinted at the presence of the third object in the system, the mass was only weakly constrained (c.f.\\ Michel 1995), although various arguments suggested that a low mass solution might be preferred (c.f.\\ Thorsett et al.\\ 1993, Sigurdsson 1993, Phinney 1993, Rasio 1994, Joshi and Rasio 1997). By 1999 Thorsett et al., using measurements of the precession of the inner orbit produced by the tidal field of the outer object, had effectively excluded stellar mass companions, and the main issue was whether the object was properly a brown dwarf, formed through star formation processes, or whether it was a planet. When the assumption is made that the pulsar spin-period derivative is negligibly small (which is true if the magnetic field of the pulsar is a typical $3\\times10^8$~G, Thorsett et al.\\ 1993, 1999), then preliminary analysis of the most recent timing data finds a unique Keplerian solution, with $m_3\\sin i\\sim1.7M_J$, $e\\sim0.13$, and orbital period $P_b\\sim68$~yrs. (A detailed analysis, relaxing the assumption on the spin-period derivative and including orbital perturbation measurements to constrain $\\sin i$ is in preparation.)  The serendipitous observations of the white dwarf companion in deep multi-epoch HST imaging of M4 (Sigurdsson et al.\\ 2003, Richer et al.\\ 2002, 2003, Bassa et al.\\ 2003) provide constraints on the white dwarf mass and hence the inner binary inclination. These results, which confirm the prediction from the radio data, strongly favor a $1-3 \\, M_J$ mass second companion.  We must ask ``is it really a planet?'' By definition, an object of this mass, orbiting a star, would be a jovian planet if observed in the solar neighbourhood. By inference, it is a gas giant. It seems very unlikely that a rocky or icy planet of such mass could form in M4 in any circumstances.  A serious issue raised by the existence of this object is whether it formed through the canonical ``bottom-up'' core accretion process, or whether it formed through ``top-down'' collapse of cold gas that had become secularly unstable under its own gravity (c.f.\\ Boss 2002, 1997). The former process would imply the existence of lower-mass planets that had failed to grow to the point where nebular gas accretes to the planet, and hence we would predict the presence or rocky or icy terrestrial planets in M4 specifically and in globular clusters more generally. In view of the observed correlation between the incidence rate of observable ``hot'' and ``warm'' jovians and host star metallicity in the solar neighbourhood (c.f.\\ Fischer and Valenti 2003), the possible prevalence of ``normal'' jovians around pop II stars would be of particular interest, independent of the formation mechanism. Clearly the possibility of low mass planet formation around pop II stars has significant anthropic implications and for considerations of the prospect for extraterrestrial life. Conversely, if the planet formed through top down collapse, we can further ask whether it formed in a cold protoplanetary disk, or as an independent quasi-spheroidal collapse---essentially as a low mass or failed brown dwarf. In principle, the different scenarios are testable. The possibility that the difficult of growing a massive icy core in low metallicity protoplanetary disks might be partially offset by the onset of gas accretion at a lower core mass is particularly intriguing (c.f.\\ Rice and Armitage 2003).  In addition to the very precise radio timing measurements of the spin frequency and its various derivatives and the observed time variation of some of the orbital parameters, a number of other observables of the system combine to provide key hints as to the formation and dynamical history of the system. These almost completely constrain its possible past. The data allow us to invert the dynamical history of the system with remarkable confidence to recover the initial conditions with relatively little ambiguity:  \\begin{itemize} \\item{} Almost as soon as the system was detected, the white dwarf orbital eccentricity was noted to be anomalously large (McKenna and Lyne 1988). If the pulsar were spun-up through conservative mass transfer from the progenitor of the white dwarf (c.f.\\ Rappaport et al.\\ 1995, Phinney and Kulkarni 1994), we would expect the white dwarf eccentricity to have been several orders of magnitude smaller after the white dwarf detached. It is generally agreed that the current eccentricity was dynamically induced after mass-transfer was complete. The process by which this occurred has been somewhat contentious, with a number of less than satisfactory processes proposed. The cause of the white dwarf orbital eccentricity is now almost certainly confirmed to be the Kozai mechanism (Ford et al.\\ 2000), in which angular momentum is exchanged from the inclination of the planet's orbit to the eccentricity of the white dwarf's orbit. The current orbits are known to be significantly non-coplanar (relative inclination of about $40^{\\circ}$), and the inclination must have been significantly larger still before angular momentum exchange took place ($\\sim 70-80^{\\circ}$), for the Kozai mechanism to have induced the observed white dwarf orbital eccentricity. \\item{} The inferred current and original high relative inclination between the inner binary plane and the orbital plane of the planet is highly significant and a major clue to the possible formation paths for this system. \\item{} The planet orbit is now thought to have relatively {\\it low} orbital eccentricity. This is a major constraint on formation scenarios, since most mechanisms that lead to high orbital inclination naturally also lead to high orbital eccentricity. Exchange processes that leave a planet in a stable circumbinary orbit tend to naturally lead to moderate eccentricities and high inclination, since high eccentricity post-exchange orbits are generally dynamically unstable, and low eccentricity orbits are improbable due to the small available phase space at low eccentricity. The current low planet orbital eccentricity may then be due to circularization of the initially moderately eccentric planet orbit during adiabatic mass loss of the white dwarf progenitor envelope, during which the system mass went from an initial $\\sim 2.3 \\msun$ to the current $\\sim 1.8 \\msun$. The planet's semi-major axis increased in proportion at that time.  This sequence requires the planet be in place before the white dwarf progenitor evolved off the main-sequence.  To a good approximation, during the RGB evolution of the white dwarf progenitor, we can treat the planet as a test particle in a Keplerian orbit about a point mass potential with a central mass equal to the total mass of the neutron star and its companion. The loss of the envelope mass is slow, and the planet orbital evolution is adiabatic; its energy changes in response to the loss of mass from the center, but the other integrals of motion are invariant. Eccentricity in general is an orbital parameter, not an invariant, however, for a Keplerian orbit, $1 - e^{2}_{p} \\propto  E_p J^2_p$; the eccentricity is a function of the integrals only, and is invariant to isotropic slow mass loss. The current eccentricity, $e_{fin} = 0.2 \\pm 0.1$, from current fits to the timing data, is somewhat lower than expected from exchange models (which predict $e \\sim 0.3-0.7$). It is possible this system had an unusually low initial post-exchange eccentricity, but this then raises the same fine tuning problems present in other scenarios.  Alternatively we can postulate a post-exchange planet eccentricity of about $0.5$ and ask whether circularization could have taken place without substantial decrease in inclination. Given a current semi-major axis of $\\sim 20 AU$, the post-exchange, pre-RGB phase, semi-major axis must have been $\\sim 16 AU$, with a corresponding periastron of $\\sim 8 AU$, assuming a median post-exchange eccentricity of $0.5$. As the white dwarf progenitor evolved, its orbit initially circularized, and then expanded as conservative mass transfer took place, until it reached its current orbit; the time scale for such a mass transfer phase and spin-up is $\\gtorder 10^8$ years. With a periastron of $\\sim 8 AU$, compared with a white dwarf progenitor final orbit of $1 AU$, we find that substantial tidal circularization of the planets orbit could have occurred at late stages of the mass transfer phase, yet the time scales are long enough that complete circularization is precluded. Following Verbunt and Phinney (1995), we estimate $\\ln \\Delta e \\sim -1$, or $\\delta e \\sim 0.3$, consistent with an initial post-exchange orbital eccentricity of $\\sim 0.3-0.5$, in the range predicted, and consistent with the currently observed eccentricity.  Inclination changes during this phase should have been negligible; unless there was significant planar outflow during mass transfer, in which case coupling of the planet to the excretion ring during plane crossings could in principle be a concern. The system here of course is a triple, not a point mass secondary interacting with an extended massive central star, but the additional lever arm of the giant rotating about the inner system center-of-mass will in general somewhat enhance the tidal torque, making this scenarios more plausible. The inferred eccentricity of the planet's orbit, before the evolution of the pulsar companion to white dwarf, is then about $0.5$, exactly in the range predicted for an exchange scenario; and is not sensitive to the exact value of the current eccentricity. Allowing the neutron star to accrete $\\sim 0.1 \\msun$ during mass transfer, with a correspondingly larger final total central mass, does not significantly change the predicted initial eccentricity; if there was slow isotropic mass loss, some additional circularization due to tidal interaction with the wind is conceivable. Mass accretion by the planet is negligible for all scenarios; mass loss due to LMXB ablation of the planet is also negligble for plausible LMXB phase luminosities.  If the planet eccentricity was initially low, negligible tidal circularisation took place, if the planet eccentricity was initially substantially higher, then there was opportunity for significant tidal circularisation during the mass transfer phase; if the mass transfer phase was extended. Since this is a long period, low mass binary pulsar, we expect sub-Eddington mass transfer with mass transfer time scales of $O(10^8)$ yrs or longer, depending on the core mass of the main sequence progenitor post-exchange and the eccentricity of the main sequence progenitor orbit before circularization of the inner orbit and onset of mass transfer (cf Burderi et al 1996, Webbink et al 1983).   \\item{} The HST data reveals the white dwarf is young. The main sequence progenitor was presumably as old as the cluster, 12.7 Gyrs, but the white dwarf emerged from the mass transfer phase just under 0.5 Gyrs ago. This is consistent with a single exchange, where the white dwarf progenitor was acquired by the neutron star some 1--2~Gyr ago, and the recoil induced by the super-elastic exchange ejected the system to the outer parts of the cluster, where stellar densities are low and interaction time scales are long. Therefore the system has probably been mostly dynamically isolated since the exchange. \\item{} The projected pulsar position is somewhat outside cluster core. The true position of the pulsar is of course probably somewhat further from the core than we see in projection. Since the system is significantly more massive than the main sequence turnoff, or indeed any likely other stellar component in M4, it tends to sink to the core through dynamical friction. The current characteristic time-scale for the orbit of the triple to sink back into the cluster core through dynamical friction is about a Gyr, and its current position is consistent with an initial orbit with an apocenter beyond the half mass radius and total dynamical lifetime of 2-3 Gyrs, which is consistent with a single exchange origin, combined with ejection from the core. \\end{itemize} Given these data, and the current pulsar kinematics, we can construct a canonical formation scenario for the system. ",
        "conclusions": ""
    },
    "0404/astro-ph0404362_arXiv.txt": {
        "abstract": "We present surface brightness fluctuations (SBFs) in the near--IR for 191 Magellanic star clusters available in the Second Incremental and All Sky Data releases of the Two Micron All Sky Survey (2MASS), and compare them with SBFs of Fornax Cluster galaxies and with predictions from stellar population models as well. We also construct color--magnitude diagrams (CMDs) for these clusters using the 2MASS Point Source Catalog (PSC). Our goals are twofold. First, to provide an empirical calibration of near--IR SBFs, given that existing stellar population synthesis models are particularly discrepant in the near--IR. Second, whereas most previous SBF studies have focused on old, metal rich populations, this is the first application to a system with such a wide range of ages ($\\sim$ 10$^6$ to more than 10$^{10}$ yr, i.e., 4 orders of magnitude), at the same time that the clusters have a very narrow range of metallicities (Z $\\sim$ 0.0006 -- 0.01, ie., 1 order of magnitude only). Since stellar population synthesis models predict a more complex sensitivity of SBFs to metallicity and age in the near--IR than in the optical, this analysis offers a unique way of disentangling the effects of age and metallicity.   We find a satisfactory agreement between models and data. We also confirm that near--IR fluctuations and fluctuation colors are mostly driven by age in the Magellanic cluster populations, and that in this respect they constitute a sequence in which the Fornax Cluster galaxies fit adequately. Fluctuations are powered by red supergiants with high--mass precursors in young populations, and by intermediate--mass stars populating the asymptotic giant branch in intermediate--age populations. For old populations, the trend with age of both fluctuation magnitudes and colors can be explained straightforwardly by evolution in the structure and morphology of the red giant branch. Moreover, fluctuation colors display a tendency to redden with age that can be fit by a straight line. For the star clusters only, (\\barH\\ - \\barKs) = (0.21$\\pm$0.03)Log(age/yr) $-$ (1.29$\\pm$0.21); once galaxies are included, (\\barH\\ - \\barKs) = (0.20$\\pm$0.02)Log(age/yr) $-$ (1.25$\\pm$0.16). Finally, we use for the first time % a Poissonian approach to establish the error bars of fluctuation measurements, instead of the customary Monte Carlo simulations. ",
        "introduction": "\\label{intro}  While the mean surface brightness of a galaxy is independent of distance, the variance about the mean decreases with distance --- i.e., given the same angular resolution, more distant galaxies appear smoother. This is the principle behind surface brightness fluctuation measurements \\citep[SBFs;][]{tonr88,blak01}, one of the most powerful methods to determine cosmological distances \\citep[e.g.,][]{tonr97,liu01,jens03}. SBFs arise from Poisson fluctuations in the number of stars within a resolution element, and they are measured through the observed ratio of the variance to the mean surface brightness of a galaxy; that is, the ratio (denoted \\barL) of the second to the first moment of the stellar luminosity function, scaled by the inverse of 4$\\pi d^2$, where $d$ is the distance. SBF measurements are expressed in \\barm\\ and \\barM, which are, respectively, the apparent and absolute magnitudes of \\barL.  SBF magnitudes, however, depend not only on galaxy distances, but also on the age and metallicity of stars. Therefore, SBFs also offer a unique possibility to investigate unresolved stellar populations.  For example, as a luminosity--weighted mean, \\barM\\ is much more sensitive to giant stars than integrated colors \\citep{wort93a,ajha94}. For the same reason, \\barM\\ is relatively insensitive to differences in the IMF for intermediate--age and old systems.  We engaged in this work with the aim of providing an empirical calibration of near--IR SBFs, specifically for the study of unresolved stellar populations. The near--IR is very favorable for SBF measurements, from the point of view of improved signal (the light of intermediate and old populations is dominated by the asymptotic giant branch, AGB, and the red giant branch, RGB), reduced dust extinction and, last but not least, the model prediction that near--IR SBFs might help break the age--metallicity degeneracy \\citep{wort93b}. However, there is the very important disadvantage that existing stellar population synthesis models are particularly discrepant in the near--IR spectral region \\citep{char96,liu00,blak01}. The disagreement is as high as $\\sim$ 0.2 mag in $(V\\!-\\!K)$, compared to $\\sim$ 0.05 mag in $(B\\!-\\!V)$. The ill--determined contribution of asymptotic giant branch (AGB) stars to the integrated light may be the most important source of this problem \\citep{ferr95}. Such an uncertainty is bound to compromise the calibration of \\barM. {\\it An empirical calibration of near--IR SBFs is therefore essential}. ",
        "conclusions": "\\label{disc}      This study has shown that in MC star clusters, most of which have roughly the same relatively low metallicity, near--IR fluctuation magnitudes and colors are driven by age.  Our result is not unexpected. In their classical study, \\citet{sear80} demonstrated that the properties of the MC clusters' integrated light {\\it in the optical wavelengths} are determined by their red giants. They also inferred that the sequence from class I to class III is one of age, and insensitive to abundance. Regarding the sequence of the older clusters, classes IV through VII, \\citet{sear80} posited that it was both sensitive to increasing age {\\it and} decreasing metallicity. Given that (a) both near--IR wavelengths and SBFs are more sensitive than integrated optical light to the red giant stars in these clusters, and (b) the star cluster metallicity stays nearly the same for classes Pre through V, it is not surprising that we have also found a sequence of age, slightly modulated by abundance in the case of the two oldest SWB classes.  It is true that the MC star clusters are mostly either too young (7 out of the 8 MC superclusters are younger than most of the Fornax galaxies) and/or too metal--poor to be relevant to the galaxies. However they seem to outline a trend with age that includes the galaxies, as is shown most clearly by Figure \\ref{flvslage}. For this reason, even if star cluster populations might be of limited direct value for the modelling of old ellipticals and spheroids, they are important for the calibration of stellar population synthesis models.  Regarding the agreement between the data and the models in their present state, we find that it is very good qualitatively, but that it could be improved in the details. For example, in principle, we could have read off the metallicity of the superclusters from Figures \\ref{flvslage} and \\ref{flcolvslage}, given their ages. However, the metallicities we would have inferred are not consistent in all cases with the ones that correspond to their SWB class. Also, the models, and in particular those with the highest metallicity, cannot reproduce the very red fluctuation colors exhibited by a few of the Fornax galaxies. Conversely, models predict redder fluctuation colors than those of intermediate--age clusters, and they underestimate the contribution of bright stars to the integrated luminosity in these same clusters. Arguably, the models work best for old, metal poor populations. This is probably not a coincidence, but is due to the fact that models have tried to match their features for the longest time. Moreover, old populations evolve more slowly. On the MC supercluster data side, the oldest clusters are also the most massive, and therefore have the smallest stochastic errors. The rapidly evolving young populations are harder to match, as are the intermediate--age ones, with their poorly, albeit increasingly better, understood asymptotic giant branches.  We plan to continue this work in various directions, e.g., improve the calibration of the models with the SBF data, compare to other models, and compute SBFs of the highest metallicity clusters in our Galaxy and in M~31. We will also investigate the relationship between \\barM$_{K_s}$ and ($V - I$) color in the MC star clusters. \\citet{liu02} discovered a linear dependence of \\barM$_{K_s}$ with ($V - I$) in a sample of 26 ellipticals, S0s, and spiral bulges, which might be tracing late bursts of star formation in these systems; a study of the MC star clusters, which probe a range twice as large in \\barM$_{K_s}$ and three times larger in color, is likely to throw light into the origin of this trend."
    },
    "0404/cond-mat0404480_arXiv.txt": {
        "abstract": "Condensed matter analogs of the cosmological environment have raised the hope that laboratory experiments can be done to test theoretical ideas in cosmology. I will describe Unruh's sonic analog of a black hole (``dumbhole'') that can be used to test Hawking radiation, and some recent proposals on how one might be able to create a dumbhole in the lab. In this context, I also discuss an experiment already done on the Helium-3 AB system by the Lancaster group. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404404_arXiv.txt": {
        "abstract": "Although variations in elemental abundance ratios in the Milky Way Galaxy certainly exist, details remain uncertain, particularly in the inner Galaxy, where stars and \\ion{H}{2} regions in the Galactic plane are obscured optically. In this paper we revisit two previously studied, inner Galaxy \\ion{H}{2} regions: G333.6$-$0.2 and W43. We observed three new positions in G333.6$-$0.2 with the Kuiper Airborne Observatory and reobserved the central position with the Infrared Space Observatory's Long Wavelength Spectrometer in far-infrared lines of S$^{++}$, N$^{++}$, N$^{+}$, and O$^{++}$. We also added the N$^+$ lines at 122 and 205 \\micron\\ to the suite of lines measured in W43 by Simpson et al. (1995). The measured electron densities range from $\\sim 40$ to over 4000 cm$^{-3}$ in a single \\ion{H}{2} region, indicating that abundance analyses must consider density variations, since the critical densities of the observed lines range from 40 to 9000 cm$^{-3}$. We propose a method to handle density variations and make new estimates of the S/H and N/H abundance ratios. We find that our sulfur abundance estimates for G333.6$-$0.2 and W43 agree with the S/H abundance ratios expected for the S/H abundance gradient previously reported by Simpson et al., with the S/H values revised to be smaller owing to changes in collisional excitation cross sections. The estimated N/H, S/H, and N/S ratios are the most reliable because of their small corrections for unseen ionization states ($\\lesssim 10$\\%). The estimated N/S ratios for the two sources are smaller than what would be calculated from the N/H and S/H ratios in our previous paper. We compute models of the two \\ion{H}{2} regions to estimate corrections for the other unseen ionization states. We find, with large uncertainties, that oxygen does not have a high abundance, with the result that the N/O ratio is as high ($\\sim~0.35$) as previously reported. The reasons for the uncertainty in the ionization corrections for oxygen are both the non-uniqueness of the \\ion{H}{2} region models and the sensitivity of these models to different input atomic data and stellar atmosphere models. We discuss these predictions and conclude that only a few of the latest models adequately reproduce \\ion{H}{2} region observations, including the well-known, relatively-large observed Ne$^{++}$/O$^{++}$ ratios in low- and moderate-excitation \\ion{H}{2} regions. ",
        "introduction": "The presence of radial abundance gradients in the plane of the Milky Way is now an established fact and is seen in both stars and gaseous nebulae (see, e.g., Henry \\& Worthey 1999; Rolleston et al. 2000; and references in both papers). The cause is generally thought to be due to the formation history and subsequent evolution of the Galaxy; thus the observed gradients are a major tool for understanding this history (see, e.g., Hou, Prantzos, \\& Boissier 2000; Chiappini, Matteucci, \\& Romano 2001; Chiappini, Romano, \\& Matteucci 2003, and references therein). The basic idea is that the inner Galaxy formed before the outer Galaxy and the higher molecular gas density in the inner Galaxy produces a higher star formation rate. The result is a greater return to the interstellar medium (ISM) in the inner Galaxy of both ``primary'' alpha elements from massive star supernovae and ``secondary'' elements like nitrogen. Secondary nitrogen is produced by CNO burning of already existing carbon and oxygen in intermediate-mass stars and subsequently returned to the ISM through mass loss. However, both the wide range of uncertain input parameters to chemical evolution models (Pagel 2001) and the uncertain details of the abundance variation of each element, primary and secondary, contribute to our less-than-complete understanding of the formation and evolution of the Milky Way. In this paper we try to improve our knowledge of the abundances in inner Galaxy \\ion{H}{2} regions and also test the methods that have been used to date to determine the abundance ratios from low-excitation, optically obscured \\ion{H}{2} regions that are needed for abundance gradient calculations.  \\ion{H}{2} regions in the inner Galaxy (galactocentric radius $R_G \\lesssim 6$ kpc, where $R_\\odot = 8$ kpc) are not accessible to optical observers because of interstellar extinction. Consequently, the abundances must be determined from far-infrared (FIR) lines, such as those of N$^{+}$ at 122 and 205 \\micron, N$^{++}$ at 57 \\micron, O$^{++}$ at 52 and 88 \\micron, Ne$^{+}$ at 12.8 \\micron, Ne$^{++}$ at 15.5 and 36 \\micron, S$^{++}$ at 18.7 and 33.5 \\micron, and S$^{3+}$ at 10.5 \\micron. Large surveys include those of Simpson et al. (1995a, hereafter SCREH) and Afflerbach et al. (1997), who observed the FIR lines (wavelengths $>$ 17 \\micron) with the Kuiper Airborne Observatory (KAO) in 18 inner Galaxy \\ion{H}{2} regions, and Mart\\'{\\i}n-Hern\\'{a}ndez et al. (2002a; 2003), who observed 13 inner Galaxy \\ion{H}{2} regions with the Infrared Space Observatory (ISO). The KAO observations included FIR measurements of the lines with wavelength $> 18$ \\micron, whereas the ISO observations included the mid-infrared lines, measured with ISO's Short Wavelength Spectrometer, SWS, but not the [\\ion{N}{2}] 205 \\micron\\ line.  Two of the abundance gradients that are most needed for galactic chemical evolution studies are the O/H and N/H ratios, representing primary and secondary elements, respectively. In optically obscured \\ion{H}{2} regions the only ionization state of oxygen available for study is O$^{++}$, which means that corrections are required for the unseen O$^+$ ions. The situation is better for nitrogen, as there are [\\ion{N}{2}] lines as well as [\\ion{N}{3}] lines in the FIR wavelength range. Unfortunately, the [\\ion{N}{2}] 122 \\micron\\ line is difficult to measure for airborne observations because of atmospheric absorption, and the [\\ion{N}{2}] 205 \\micron\\ line is difficult because of poor detector sensitivity. In addition, the [\\ion{N}{2}] lines are difficult to interpret because they are collisionally de-excited at much lower densities than the lines from the other ions. Consequently, the papers to date have used only the [\\ion{N}{3}] and [\\ion{O}{3}] lines to estimate N/O and the secondary/primary abundance ratio in optically obscured \\ion{H}{2} regions (SCREH; Afflerbach et al. 1997; Mart\\'{\\i}n-Hern\\'{a}ndez et al. 2002a). This is particularly troubling since a majority of the nitrogen and an even larger fraction of the oxygen are singly ionized in the high-metallicity \\ion{H}{2} regions of the inner Galaxy (note that in the outer Galaxy, the lower metallicity \\ion{H}{2} regions can be much more highly excited and thus the N$^{++}$/O$^{++}$ ratio is often a good approximation to N/O, Rudolph et al. 1997). On the other hand, a majority of the primary element sulfur is doubly ionized in \\ion{H}{2} regions. Here we devise a method of analyzing the density-sensitive [\\ion{N}{2}] lines to more accurately estimate nitrogen abundances with respect to sulfur without the need for large corrections for unseen ionization states. Using the observed N$^{++}$/N$^{+}$ ratio and calculated \\ion{H}{2} region models, we then estimate the ionization correction factors (ICFs) for oxygen.  Two of the most luminous \\ion{H}{2} regions in the Galaxy are G333.6$-$0.2 at $R_G \\sim 5.5$ kpc (Wilson et al. 1970) or 5.6 kpc (Caswell \\& Haynes 1987) and W43 (G30.8$-$0.0) at $R_G \\sim 4.2$ kpc (Reifenstein et al. 1970) for $R_\\odot = 8.0$~kpc. The radio fluxes from these objects require ionizing stellar luminosities of $10^{50}$ and $2.9\\times10^{50}$ photons s$^{-1}$, respectively (see Rubin 1968b). Both \\ion{H}{2} regions are located in the Scutum-Crux spiral arm (see Vallee 2002) in the Milky Way molecular ring. Both are excited by clusters of massive stars, observed at near-infrared (NIR) wavelengths. The cluster exciting G333.6$-$0.2 contains stars that are very young, embedded in dust (Blum et al. 2002), in keeping with the youthful appearance of this extremely dense, compact \\ion{H}{2} region. On the other hand, the cluster exciting W43 must be older and more evolved because it contains a Wolf-Rayet star in addition to O giants or supergiants (Blum, Damineli, \\& Conti 1999; Cotera \\& Simpson 1997), although there are still numerous stars forming in the associated molecular clouds (Motte, Schilke, \\& Lis 2003). This is also consistent with the more mature appearance of W43, where the ionized gas is well separated from the exciting star cluster.  In this study, we report measurements of the [\\ion{N}{2}] lines in the three positions observed in W43 by SCREH and numerous lines in three additional positions in G333.6$-$0.2 north of the center position observed by Colgan et al. (1993) with the KAO; we also report measurements of the central position in G333.6$-$0.2 with ISO's Long Wavelength Spectrometer (LWS). In Section 2 we describe the observations, in Section 3 we discuss analysis techniques that are used to determine abundances in both low- and high-density \\ion{H}{2} regions, in Section 4 we model the two \\ion{H}{2} regions to determine the ICFs needed for those elements without measured abundances of the singly ionized atoms, and in Section 5 we discuss models in general and give our conclusions. By measuring both N$^{+}$ and N$^{++}$, we hope to avoid the need for large (and uncertain) ICFs for nitrogen and thus improve the reliability of the secondary/primary abundance ratio measurement in the Milky Way molecular-ring \\ion{H}{2} regions. ",
        "conclusions": "\\subsection{Stellar Atmosphere Models}  As was discussed above, many of the ionization correction factors are critically dependent on the choice of stellar atmosphere model. This can be demonstrated either by plotting ratios of the numbers of photons able to ionize various ions as a function of \\teff\\ or by plotting ionization fractions or ICFs calculated from models; here we do the latter because ICFs are more readily applicable to abundance computation. This is illustrated in Figure 5, where we plot the ICF ratios $<$N$^{++}$/N$>$/$<$O$^{++}$/O$>$, $<$O$^{++}$/O$>$/$<$S$^{++}$/S$>$, and $<$O$^{++}$/O$>$ versus the $<$N$^{++}$/N$>$/$<$N$^{+}$/N$>$ ratio, which indicates excitation without the necessity for any abundance assumption. The models from Rubin (1985) that are plotted in Fig. 5 all have total nucleon density $= 10^3$ cm$^{-3}$, ionizing luminosities of $10^{49}$ to $10^{50}$ photons s$^{-1}$, and Orion Nebula metallicities; the models from SCREH had similar densities and ionizing luminosities but metallicities ranging from 1 to 2 times Orion nebula metallicities. The models from both Rubin (1985) and SCREH were computed with the versions of NEBULA available at the time and both used LTE stellar atmosphere models from Kurucz (1979). In addition to the models described in \\S 4, constant nucleon density ($ = 10^3$ cm$^{-3}$) spherical models with ionizing luminosities of $10^{49}$ photons s$^{-1}$ were computed and plotted. These recent NEBULA models used stellar atmosphere models from either Pauldrach et al. (2001), Lanz \\& Hubeny (2003), or contemporary\\footnote{\\url{http://kurucz.harvard.edu/}} LTE stellar atmosphere models described by Kurucz (1992). The stellar atmosphere models had ``Solar'' abundances and the nebular models had 1.5 times Orion Nebula abundances (Simpson et al. 1998) such that the model abundance ratios (all times $10^{-4}$) are C/H = 3.75, N/H = 1.02, O/H = 6.0, Ne/H = 1.5, S/H = 0.11, and Ar/H = 0.04. Because they are widely used, we also plot ICFs from the models of Stasi\\'nska \\& Schaerer (1997). For these models both the stellar and nebular abundances were ``Solar''. We note that some of the elemental abundances considered to be ``Solar'' are changing with time; for example the ``Solar'' O/H ratio used by Stasi\\'nska \\& Schaerer (1997) was O/H$ = 8.51 \\times 10^{-4}$ (Anders \\& Grevesse 1989), whereas currently, the ``Solar'' O/H ratio is thought to be $5.45 \\times 10^{-4}$ (Holweger 2001). Mart\\'{\\i}n-Hern\\'andez et al. (2002b), Smith et al. (2002), and Morisset (2004) discuss the effects of stellar abundances on \\ion{H}{2} ionization. Details of the physics used, though, are much more important than the exact ``Solar'' abundances, as we will see.  Unless an \\ion{H}{2} region can be observed at both optical and infrared wavelengths, observers cannot measure the abundances of N, O, Ne, S, and Ar without making corrections for unseen ionization states for some elements. For optically obscured \\ion{H}{2} regions like W43 and G333.6$-$0.2, the most significant missing ion is O$^+$, since oxygen is the most abundant heavy element and since these low excitation \\ion{H}{2} regions are believed to have most of their oxygen as O$^+$. Rubin et al. (1994) and SCREH assumed that they knew the O/S ratio and used the observed O$^{++}$/S$^{++}$ ratio to estimate the ICF for their measured N$^{++}$/O$^{++}$ ratios, since they did not have any other line pair that could be used to indicate excitation. In this paper we use our observed N$^{++}$/N$^{+}$ ratios to estimate \\teff; however, we see from Fig. 5a how dependent the N$^{++}$/O$^{++}$ ICF is on the stellar atmosphere model.  Mart\\'{\\i}n-Hern\\'{a}ndez et al. (2002a) took advantage of their ISO SWS measurements of the Ne$^{++}$/Ne$^{+}$ ratios to estimate ICFs without the need for assuming {\\it a priori} some abundance ratio; unfortunately, the large beam size difference between the SWS and LWS instruments means that their observed Ne$^{++}$/Ne$^{+}$ ratios may not always be appropriate for estimating the excitation of O$^{++}$ and N$^{++}$. Moreover, Ne$^{++}$ requires ionizing photons $> 41$ eV, and it is at these high energies that the stellar atmosphere models are especially divergent, as we now discuss.  An important test of the validity of stellar atmosphere models of hot stars is whether \\ion{H}{2} region models produced with these atmospheres predict line fluxes that agree with observations. Except for a few rare instances, this is the only way that the EUV fluxes of the models can be tested, since interstellar absorption by hydrogen predominantly prevents the EUV fluxes from hot O stars from reaching the Earth. Here we discuss the stellar fluxes that can doubly ionize neon (photon energy $ > 41$ eV). This has been called the [\\ion{Ne}{3}] problem (e.g., Sellmaier et al. 1996) because it has been observed that the Ne$^{++}$/O$^{++}$ ratio is relatively constant over a large range of \\ion{H}{2} region excitation (SCREH; Stasi\\'nska \\& Schaerer 1997 and references therein; Kennicutt et al. 2003), contrary to predictions of models that used LTE stellar atmospheres (e.g., SCREH). Proposed solutions have involved using non-LTE stellar atmospheres, usually with winds (Rubin, Kunze, \\& Yamamoto 1995; Sellmaier et al. 1996; Stasi\\'nska \\& Schaerer 1997). Figure 6 shows the observations of SCREH along with the predictions of the models used in Fig. 5 and the G333.6$-$0.2 and W43 models described in \\S4. Models calculated with black bodies for the stellar atmosphere spectrum are also plotted since Morisset et al. (2004) found that such models give a reasonable fit to the mid-IR ISO observations. Like the new nebular models plotted in Fig. 5, these models also have spherical geometry, constant nucleon density of $10^3$~cm$^{-3}$, $10^{49}$ ionizing photons, and 1.5 times the Orion Nebula abundances  listed in Table 3.  Surprisingly, in Fig. 6 only the models of Stasi\\'nska \\& Schaerer (1997) reproduce the observations over the total range of observed O$^{++}$/S$^{++}$, although if Ne/O is as large as 0.25 in all \\ion{H}{2} regions as it is in the Orion Nebula (Simpson et al. 1998; Table~3), these \\ion{H}{2} region models produce too high a Ne$^{++}$/O$^{++}$ ratio because of the very large EUV fluxes in Schaerer \\& de Koter's (1997) stellar atmosphere models. Moreover, Schaerer \\& de Koter's models have too much EUV flux for reproducing other observations (e.g., Smith et al. 2002). Earlier, Sellmaier et al. (1996) had demonstrated that the [\\ion{Ne}{3}] problem was supposedly solved when they obtained a good fit to the data by using non-LTE atmospheres with winds computed with Pauldrach's code as it then existed. However, our \\ion{H}{2} region models with non-LTE stellar atmospheres with winds from Pauldrach et al. (2001) taken from both A. Pauldrach's web site and from the stellar atmosphere models produced by Smith et al. (2002) using Pauldrach's WM-basic code predict much lower Ne$^{++}$/O$^{++}$ than observed for Dwarf atmospheres with \\teff\\ $< 40$ kK and for Supergiant atmospheres with \\teff\\ $< 35$ kK. Our \\ion{H}{2} region model produced with a 35~kK supergiant atmosphere from F. Sellmaier (private communication in 1995) lies above the Pauldrach et al. Supergiant line but not as high as the model with the 35~kK supergiant atmosphere in Sellmaier et al.'s (1996) Figure 2, even though we are now assuming Ne/O = 0.25 (Simpson et al. 1998) instead of 0.2025. Part of the difference of the current nebular model predictions with previous models is the higher O$^{++}$ ionization relative to Ne$^{++}$ and S$^{++}$ obtained with the Opacity Project cross sections described in Section 4. The model computed with Sellmaier's (1995) atmosphere has higher ionization than the model computed with Pauldrach et al.'s (2001) S-35  atmosphere because it used Orion Nebula abundances (as in Sellmaier et al. 1996) instead of the higher metallicity (1.5 times the Orion Nebula abundances of Table 3) used for the rest of the models plotted here. The result, though, is a return of the mismatch between the stellar atmosphere models and the EUV fluxes of real stars, with possible additional discrepancies for the ionization, recombination, and collisional excitation cross sections for oxygen, neon, and sulfur. Moreover, we see that all the models computed without winds do a poorer job of reproducing the observations than both the models computed with non-LTE atmospheres with winds and the models computed with simple black bodies for stellar atmospheres (for this reason we give low weight to the T model for W43 in Table 5). This probably means that the codes producing the stellar atmosphere models used for computation of \\ion{H}{2} region models should include the effects of stellar winds.  \\subsection{Abundance Ratios}  At the low excitations seen in inner Galaxy \\ion{H}{2} regions, the predicted flux of the O$^{++}$ lines is extremely sensitive to the stellar atmosphere, but if the models are accurate, the O/H ratio of both G333.6$-$0.2 and W43 could be as low as $3 - 5 \\times 10^{-4}$, substantially lower than would be predicted from the observed Orion Nebula abundance and an O/H abundance gradient of $-0.06$ dex kpc$^{-1}$ (Henry \\& Worthey 1999). This low O/H ratio is hard to understand because of the low $T_e$ measured in inner Galaxy \\ion{H}{2} regions from radio recombination lines. There would need to be some source of systematic error for the O/H ratio to be higher than these estimates: i.e., what would be required is that the excitation of oxygen in the two \\ion{H}{2} regions is lower than estimated in this paper from He$^+$/H$^+$ recombination lines and N$^+$ and N$^{++}$ forbidden lines. Possible reasons that the G333.6$-$0.2 and W43 excitation might be overestimated are as follows: (1) The abundance of N$^+$ is underestimated because the [\\ion{N}{2}] line fluxes are undermeasured owing to telluric absorption at 122 \\micron\\ or diffraction at 205 \\micron. This is a particular problem for W43, where the two measurements of the [\\ion{N}{2}] 122 \\micron\\ line in different apertures can be used to produce very different electron densities and N$^{++}$/N$^{+}$ abundance ratios (0.6 to 3.2 in Table 5). (2) The estimated N$^{++}$/N$^{+}$ ratio is too high because some of the estimated N$^+$ is missing owing to inadequate correction for density variations. (3) The estimated N$^{++}$/N$^{+}$ ratio is too high because of incorrect collisional excitation cross sections for the pertinent energy levels. (4) The models predict too much O$^{++}$/O versus N$^{++}$/N$^{+}$. The spread of O$^{++}$/O ratios in Fig. 5c shows that any excessive O$^{++}$/O is due to model atmospheres and not due to the $\\sim 30$\\% higher oxygen ionization resulting from the Opacity Project cross sections. (5) The ionization of helium might be overestimated if the He/H abundance ratio is actually greater than the assumed 0.10. We conclude that determinations of the O/H abundance ratio in low excitation \\ion{H}{2} regions from FIR data have large uncertainties at this time, with the largest contributor to the uncertainty being the choice of stellar atmosphere models, with some possible contribution from the atomic data used in the nebular models.  We have demonstrated that the abundances of nitrogen and sulfur can be obtained from FIR observations and the ratio of N/S can be obtained when both ionization states of N$^{+}$ and N$^{++}$ are measured. (Simpson et al. 1998, Mart\\'{\\i}n-Hern\\'{a}ndez et al. 2002a and 2003, and Giveon et al. 2002 also show that the abundances of Ne/H and Ar/H, as well as S/H, can be determined from mid-IR observations with little bias owing to excitation or abundance.) Compared to SCREH, the S/H abundance ratio is higher in G333.6$-$0.2, owing to the new data at the additional North positions, but lower in W43 because of the higher collisional excitation cross sections for sulfur. The N/H ratio is about 50\\% larger in G333.6$-$0.2, owing to the new data at the North positions, but a factor of 2 smaller in W43, owing to the smaller ICF for N$^+$. However, since the estimated O/H ratio is also much smaller in W43, now that we are estimating the ionization from N$^{++}$/N$^+$ and not from O$^{++}$/S$^{++}$ with the assumption of a constant O/S ratio as SCREH did, the estimated N/O ratios agree with SCREH for both G333.6$-$0.2 and W43. With this new analysis, the S/H and N/O ratios are also consistent with SCREH's measurements.  \\subsection{Conclusions Regarding Systematic Errors}  There are several possible sources of systematic error that must be considered when measuring abundances in obscured, low-excitation \\ion{H}{2} regions, that is, all the \\ion{H}{2} regions in the inner Galaxy. The impact of these systematic errors is that there is a very large uncertainty in the total metallicity and important abundance ratios like N/O. These sources of systematic error are the following:  (1) The most abundant heavy element ion in low-excitation \\ion{H}{2} regions, O$^+$, has no bright infrared lines. The consequence is that the abundance of this critically important element must be estimated from the abundance of an ionization state with fractional abundance $< 0.5$.  (2) Nitrogen, the secondary element which one needs for studies of galaxy chemical evolution, has two important ionization states: N$^+$ as well as N$^{++}$. Fortunately, both species have measurable FIR lines, but {\\it both} must be measured to determine the excitation.  (3) Moreover, both FIR N$^+$ lines have greatly different values of $N_{crit}$ from those of S$^{++}$ or N$^{++}$. Thus one needs to be aware of and compensate for the ever-present density variations.  (4) The average density could be quite different in the various ionization zones, as shown by the different electron densities derived from $N_e$-diagnostic line pairs: low ionization (N$^+$), intermediate ionization (S$^{++}$), and high ionization (O$^{++}$).  (5) ICFs estimated from models have major uncertainties owing to uncertainties in the stellar atmosphere models and possibly the atomic data used in the \\ion{H}{2} region models.  It is clear that the abundances of primary (oxygen, neon, or sulfur) and secondary (nitrogen) elements will need to be measured in many more inner Galaxy \\ion{H}{2} regions before we understand the chemical evolution of the Milky Way. In the near future, measurements of N$^{++}$, N$^+$, O$^{++}$, Ne$^{++}$, and S$^{++}$ will be possible from the Stratospheric Observatory for Infrared Astronomy, SOFIA. We recommend that future observers measure excitation-sensitive lines, such as Ne$^+$ and Ne$^{++}$, as well as both N$^+$ and N$^{++}$ to determine the \\ion{H}{2} region ionization state needed to estimate the ICF for O$^+$. It is important that a sufficient number of lines be measured in each \\ion{H}{2} region (and extended \\ion{H}{2} regions mapped) so that detailed models of the \\ion{H}{2} region can be made. Discrepancies between the predicted line fluxes from the models and the observations can then be used to indicate EUV energy regions where the stellar atmosphere models may need revision."
    },
    "0404/astro-ph0404318_arXiv.txt": {
        "abstract": "We use numerical simulations of a cosmological volume to study the X-ray ionisation and heating of the intergalactic medium by an early population of accreting black holes. By considering theoretical limits on the accretion rate and observational constraints from the X-ray background and faint X-ray source counts, we find that the maximum value of the optical depth to Thompson scattering that can be produced using these models is \\taue$\\simeq 0.17$, in agreement with previous semianalytic results.  The redshifted soft X-ray background produced by these early sources is important in producing a fully ionised atomic hydrogen in the low density intergalactic medium before stellar reionisation at redshift $z \\sim 6-7$. As a result stellar reionisation is characterised by an almost instantaneous ``overlap phase'' of \\HII regions. The background also produces a second \\GII reionisation at about redshift three and maintains the temperature of the intergalactic medium at about 10,000 K even at low redshifts.  If the spectral energy distribution of these sources has a non-negligible high energy power-law component, the luminosity in the soft X-ray band of the ``typical'' galaxies hosting intermediate-mass accreting black holes is maximum at $z \\sim 15$ and is about one or two orders of magnitude below the sensitivity limit of the Chandra deep field. We find that about a thousand of these sources may be present per square arcmin of the sky, producing potentially detectable fluctuations. We also estimate that a few rare objects, not present in our small simulated volume, could be luminous enough to be visible in the Chandra deep field.  XEUS and Constellation-X satellites will be able to detect more of these sources that, if radio loud, could be used to study the 21 cm forest in absorption.  A signature of an early X-ray preionisation is the production of secondary CMB anisotropies on small angular scales ($<1$ arcmin). We find that in these models the power spectrum of temperature fluctuations increases with decreasing angular scale ($\\Delta T \\sim 16 \\mu$K at $\\sim 1$ arcsec scales), while for stellar reionisation scenarios the power decreases on smaller scales.  We also show that the redshifted 21 cm radiation from neutral hydrogen can be marginally detected in emission at redshifts $7<z<12$. At a redshift of about $z \\sim 30$ a stronger and narrower (in redshift space) signal in absorption against the CMB, that is peculiar to these models, could be detectable. ",
        "introduction": "The WMAP satellite measured an optical depth to Thompson scattering of the IGM \\taue$\\simeq 0.17 \\pm 0.04$ \\citep{Kogut:03}. The result, implies a much earlier start of reionisation with respect to the redshift when reionisation was completed (at $z_{\\rm rei} \\sim 6$, as estimated from the absorption spectra of high redshift quasars). This has also been interpreted as one of the stronger pieces of evidence for the importance of zero-metallicity (\\pop3) stars in the early universe \\citep[\\eg,][]{Cen:03b, WyitheL:03, Somerville:03, CiardiFW:03, Sokasian:03, ChiuFO:03} and of the formation of the first small mass galaxies (with $M_{dm} \\sim 10^6-10^8$ M$_\\odot$), a topic that is still under debate \\citep[\\eg,][]{HaimanRL:97, RicottiGS:01, RicottiGSb:02, Machacek:03}.  In a previous paper (paper~I), \\cite{RicottiOI:03} have argued that, even assuming the most favorable properties for \\pop3, it is difficult or impossible to produce the measured value of \\taue using UV radiation from \\pop3 stars. The main problems with this scenario are the negative feedback from ionising radiation, mechanical feedback from SN explosions and the contamination of high density regions by detritus from the same stars which produce the ionising radiation.  Taking into account the metal enrichment from SNe and pair instability SNe the \\pop3 epoch turns out to be so short lived that \\pop3 stars will never be able to complete the reionisation of the IGM.  But if a significant fraction of \\pop3 stars, instead of exploding as SNe, implode into black holes (BHs), then the negative feedback on star formation is much reduced and the UV radiation produced by thermonuclear reactions in the first stars might be important for reionisation. In \\cite{RicottiO:03} (paper~IIa) we have shown that in such a scenario, the X-ray background produced by accretion onto stellar mass seed black holes would be more effective in producing the large \\taue measured by WMAP than the UV radiation from the parent stars. \\cite{MadauR:03} have also investigated a scenario in which accretion on intermediate-mass black holes produces the large optical depth to Thompson scattering measured by WMAP. In their paper they focus is on the physics of black hole mergers and accretion onto seed black holes. Their paper is complementary to our studies that, instead, focus on radiative transfer processes.  Before the WMAP measurement of a large \\taue, the effect of X-ray preionisation has been investigated by \\cite{Oh:00} and \\cite{Venkatesan:01}. The values of \\taue that they found were smaller since the \\HI ionisation was produced by secondary photoelectrons that can ionise the gas to a maximum of 10\\% ionisation fraction. We find larger values of \\taue, consistent with WMAP, because of the additional ionisations produced by the redshifted soft X-ray background.  \\begin{table*} \\centering \\caption{Input parameters for hydrodynamic simulations with radiative transfer.\\label{tab:1}} \\begin{tabular*}{15. cm}[]{ll|cccccccccc} \\# & {RUN} & {$N_{\\rm box}$} & {$L_{\\rm box}$} & {Mass Res.} & {Res.} & {$g_\\nu$} & {$\\epsilon_*$}&{$\\epsilon_{\\rm UV}$\\fesc} &{$\\epsilon_{\\rm qso}$} & $z_{\\rm off}$ & $F_{\\rm IMF}$ \\\\ {} & {} & {} & {$h^{-1}$ Mpc} & {$h^{-1}$ M$_\\odot$} & {$h^{-1}$ pc} & {} & {} & {} & {} & {} & {}\\\\ \\hline 1 & M-PIS & 128 & 1.0  & $3.94\\times 10^4$ & 488 & Pop~III & 0.1 & $2\\times 10^{-3}$ & $2\\times 10^{-3}$ & 8 & 7\\\\ 2 & M-SN1 & 128 & 1.0  & $3.94\\times 10^4$ & 488 & Pop~III & 0.1 & $3\\times 10^{-4}$ & $2\\times 10^{-3}$ & 10 & 1\\\\ 3 & M-SN2 & 128 & 1.0  & $3.94\\times 10^4$ & 488 & Pop~III & 0.1 & $3\\times 10^{-5}$ & $2\\times 10^{-3}$ & 11 & 0.1\\\\ 4 & M-BH & 128 & 1.0  & $3.94\\times 10^4$ & 488 & Pop~III & 0.1 & $3\\times 10^{-5}$  & see Eq.~\\ref{eq:X3} & 13 & 0.1\\\\ \\end{tabular*} \\flushleft {Parameter description. {\\em Numerical parameters:} $N_{\\rm box}^3$ is the number of grid cells, $L_{\\rm box}$ is the box size in comoving \\Mpc. {\\em Physical parameters:} $g_\\nu$ is the normalised SED (\\pop3). The meaning of the other parameters is explained in the text (\\S~\\ref{ssec:code}).} \\end{table*}  In this paper we complement the semianalytic results presented in paper~IIa by using hydrodynamic cosmological simulations that include a recipe for star formation and 3D radiative transfer for hydrogen and helium ionising radiation. An approximate solution of the radiative transfer equations \\citep{GnedinA:01} is used to speed up the calculations and make the coupling of the radiative transfer and hydrodynamic equations computationally feasible. In addition, radiative transfer for the optically thin X-ray radiation and \\H2 dissociating radiation is solved exactly. Line radiative transfer in the \\H2 Lyman-Werner bands ($11.3 < h_P\\nu <13.6$ eV) is also solved exactly for the volume averaged component of the radiation field. Using this method we are able to simulate local and global radiative feedback effects of galaxy formation on cosmological scales.  The results of the cosmological simulations are used to predict the distinctive observational signatures of the X-ray preionisation scenario compared to stellar reionisation models.  We find that, if this scenario is correct, a new population of X-ray sources that do not have optical counterparts may be detected in the Chandra deep field. Interestingly, these sources, that might be high redshift AGNs but also galaxies hosting a large number of bright Ultra-luminous X-ray sources (ULXs), might have been already observed \\citep{Koekemoer:03}.  Simulations of a cosmological volume allow us to construct maps for the secondary CMB anisotropies produced after recombination by the first sources of ionising radiation. We also use the simulation outputs and the semianalytic models presented in paper~IIa to calculate the redshifted 21 cm radiation from the IGM prior to reionisation in emission or absorption against the CMB. In the X-ray preionisation scenario the IGM is only partially ionised at redshift $z>z_{\\rm rei} \\simeq 6-7$.  Therefore, the 21cm signal from the IGM at $z<12$ can be detected even though it is rather weak. Note that, in stellar reionisation scenarios that can produce the optical depth to Thompson scattering observed by WMAP, the expected 21cm signal is instead undetectable.  Finally we estimate the additional high energy background due to the postulated population of high redshift X-ray sources and its subsequent signatures [$(10 \\pm 5)$\\% in the 2-50 keV bands].  This paper is organised as follows: in \\S~\\ref{sec:res_cs} we show the results of cosmological simulations of X-ray preionisation by mass accretion on seed BHs. In \\S~\\ref{sec:lum} we estimate the number of detectable X-ray point sources at high redshift. In \\S~\\ref{sec:21cm} we calculate the redshifted 21cm signal for one of our simulations in absorption/emission against the CMB. In \\S~\\ref{sec:cmb} we compute the amplitude of the power spectrum of CMB secondary anisotropies on scales of a few arcmin produced by X-ray preionisation and stellar reionisation scenarios. We summarise the results in \\S~\\ref{sec:conc} and we discuss the observational signatures of the X-ray preionisation scenario when compared to stellar reionisation scenarios. ",
        "conclusions": "\\label{sec:conc}  The present paper is the third in a series (see paper~I and paper~IIa) devoted to the study of physically plausible reionisation scenarios. The models have optical depths to Thompson scattering \\taue$\\sim 0.17$, as measured by WMAP \\citep{Kogut:03}, and are consistent with observations of the IGM optical depth to \\lya and \\lyb photons toward high redshift quasars \\citep{Fan:03, Songaila:04}. We have studied models of reionisation by stars (paper I) or mini-quasars (paper IIa) considering realistic physical scenarios and observational constraints. Our main conclusion is that it is very difficult for any of these models to produce optical depths in excess of 0.17. If this large optical depth is produced by ultraviolet radiation from stellar sources (\\ie, \\pop3 stars with a top-heavy IMF), we find that zero-metallicity stars must be the dominant mode of star formation up to redshift $z\\sim 10$. If this scenario is correct, we need to understand the reasons for the very inefficient mixing of metal enriched gas from SNe with the gas in which star formation takes place. At the moment this inefficient mixing is not reproduced by numerical simulations. A feasible alternative scenario requires that most \\pop3 stars must implode into black holes without exploding as SNe or have subluminous SN explosions with a large amount of metal fall-back onto the compact remnant in order to reduce their metal yields. A scenario in which a large fraction of \\pop3 stars end their lives as pair-instability supernovae is not compatible with the large optical depth to Thompson scattering measured by WMAP.  In this paper we use cosmological simulations to study a reionisation scenario in which standard reionisation by \\popII stars is preceded by partial ionisation and reheating at early times by an X-ray background (\\ie, ``X-ray preionisation'' models studied in paper~IIa using semianalytic simulations). In these models the ionisation rate from the secondary radiation produced by accretion on compact remnants from the first stars dominates over the primary ultraviolet radiation emitted by the stars during their lifetime. The most appealing aspect of these models is their insensitivity to the duration of the epoch of \\pop3 stars domination. In fact, the relative importance of \\pop3 stars with respect to \\popII stars in practice has to be treated as a free parameter due to the large uncertainties in modelling the complex physical processes that regulate metal production and mixing.  In order to reproduce WMAP results we find that a fraction of about $10^{-4}$ of all the baryons in the Universe needs to be accreted by compact objects before redshift 10.  This fraction is comparable in mass to the observed mass density of supermassive BHs in the galactic nuclei today.  This does not pose a strong constraint on the models because a sizeable fraction of these early intermediate-mass BHs is expected to be removed from the galaxies during the last phases of the merger or to reside in the interstellar medium of galaxies, being hardly detectable \\cite{MadauR:01}. Perhaps a small fraction of these intermediate-mass BHs may be accreting and contribute to the observed population of ULX \\citep{Agol:02,Krolik:04}.  The stronger constraint on the model is posed by the observed soft X-ray background. Assuming that the early mini-quasars population produces an optical depth to Thompson scattering \\taue$\\sim 0.17$, their contribution to the background in the 5-50 keV bands is 5-20\\%. Future X-ray missions may be able to detect these sources if they exist. The predicted fluxes of the most rare objects is at the limit of the detection of the Chandra deep fields and about 1000 objects with flux $10^{-18}$ erg cm\\mm s\\m should be present per 5.3 arcmin$^2$.  Note that these calculations assume a spectral energy distribution of the sources that has a non-negligible high energy power-law component (\\cf, \\fig~4 in paper~IIa), in agreement with observations of QSOs and ULX in the local universe.  If the source spectra are dominated by a multicolour disk thermal component their contribution to both the X-ray background in the 2-50 keV bands and the faint source counts would be negligible.  The redshifted X-ray background also has interesting consequences for the reionisation history of He and the thermal history of the IGM at redshift $z \\sim 3$, independently of the assumed spectra of the sources. In this paper we confirm the results of paper~IIa in which we found that \\GII is almost fully reionised for the first time at redshift $z \\sim 17$ and afterwards slowly recombines before experiencing a second reionisation at redshift $z \\sim 3$ produced by the redshifted X-ray background. The heating rate from the background radiation keeps the temperature of the IGM at about 10,000 K, in rough agreement with observations of the line widths of the \\lya forest at $z \\sim 3-4$. We also emphasise that the redshifted X-ray background is important in producing a fully ionised atomic hydrogen in the low density intergalactic medium before stellar reionisation at redshift $z \\sim 6-7$. As a result stellar reionisation is characterised by an almost instantaneous ``overlap phase'' of \\HII regions.  The patchy topology of reionisation produced by stellar sources contrast with the spatially homogeneous partial ionisation by X-rays. This produces distinctive signatures on temperature/polarisation CMB anisotropies and on the redshifted 21cm emission/absorption from neutral hydrogen at high redshift. \\begin{enumerate} \\item The power spectrum of the EE polarisation is sensitive to the visibility function $g(z)$, defined in \\eq~(\\ref{eq:vis}). The Plank satellite should be able to distinguish between the visibility function produced by an early X-ray partial ionisation (\\cf, \\fig~\\ref{fig:srei5} and \\fig~8 in paper~IIa) or the one expected for reionisation by stellar sources (\\cf, \\fig~4 in paper~I). \\item On small angular scales ($< 1 arcmin$ or $l>10^4$) the secondary anisotropies produced by non-linear structures during the early reionisation epochs dominate over other contributions, offering a unique opportunity to study the first episode of structure formation.  A remarkable feature can be observed: while in a model without early episode of X-ray emission the power spectrum falls off at small angular scales approximately as $l^{-3/4}$, in X-ray preionisation models the power instead grows roughly as $l^{3/4}$. The temperature anisotropies are of the order of $\\Delta T \\sim 16 \\mu$K at $\\sim 1$ arcsec scales. But it is extremely difficult to detect anisotropies on these scales, partly because of foreground sources contamination \\citep{Fomalont:93,Church:97}. In models in which reionisation is produced by UV from stellar sources the secondary anisotropies produced by the nonlinear Ostriker-Vishniac effect (\\ie, assuming uniform ionisation fraction) have less power than in the full calculation that includes the ``patchy reionisation'' signal (the signals are correlated). In X-ray preionisation models the nonlinear Ostriker-Vishniac effect and ``patchy reionisation'' signals are instead slightly anti-correlated. \\item The redshifted 21cm emission/absorption from neutral hydrogen is another powerful probe of the ionisation state of the IGM at high redshift.  In X-ray preionisation models a sharp absorption feature at $z\\sim 25-30$ falls just outside the broadcast TV bands. This feature will appear as a roughly 100 mK ``absorption line'' in the spectrum of the Galactic foreground, and may be observable with LOFAR. The partial ionisation of the IGM in the X-ray preionisation scenarios offers better opportunities to observe the fluctuations of the redshifted 21cm line in emission at redshifts lower than 10, where the signal is easier to measure. \\end{enumerate}  Finally indirect signatures of X-ray preionisation are related to the discovery of massive BHs in the nuclei of dwarf galaxies and/or the identification of intermediate mass BHs in the ISM of galaxies (\\eg, ULX).  If \\pop3 stars are instead responsible for the large optical depth measured by WMAP future observations with JWST and ground based near-infrared \\lya surveys using gravitational lenses may soon be able to observe these objects \\citep{Pello:04, RicottiH:04} and probe the ionisation state of the IGM.   \\subsection*{ACKNOWLEDGEMENTS} MR is supported by a PPARC theory grant. NG was partially supported by by NSF grant AST-0134373 and by National Computational Science Alliance under grant MCA03S023 and utilised IBM P690 array at the National Center for Supercomputing Applications. Research conducted in cooperation with Silicon Graphics/Cray Research utilising the Origin 3800 supercomputer (COSMOS) at DAMTP, Cambridge.  COSMOS is a UK-CCC facility which is supported by HEFCE and PPARC. MR thanks Martin Haehnelt and the European Community Research and Training Network ``The Physics of the Intergalactic Medium'' for support.  The authors would like to thank Andrea Ferrara, Martin Haehnelt, Piero Madau and Martin Rees for stimulating discussions."
    },
    "0404/astro-ph0404068_arXiv.txt": {
        "abstract": "We present theoretical models of black hole (BH) populations in young stellar environments, such as starbursts and young star clusters. Using a population synthesis approach we compute the formation rates and characteristic properties of single and binary BHs for various representative ages and choices of parameters. We find that most of the BHs (typically $80\\%$ for an initial 50\\% binary fraction) are single, but with many originating from primordial binaries (which either merged into a single massive star or were disrupted following a supernova explosion). A smaller but significant fraction (typically $20\\%$) of the BHs remain in binary systems. Main-sequence stars are the most frequent BH companions, but massive BH--BH binaries are the next most numerous group. The most massive BHs found in our simulations reach $\\sim 80\\,M_\\odot$, and are formed through mergers of massive binary components. If formed in a dense star cluster such a massive stellar BH may become the seed for growth to a more massive (``intermediate-mass'') BH. Although we do not include dynamical interactions, our results provide realistic initial conditions for $N$-body simulations of dense star clusters (e.g., globular clusters) including primordial BHs. ",
        "introduction": "Until recently it was believed that BHs existed in two separate mass ranges: stellar-mass BHs, with masses $\\sim 10\\,M_\\odot$, and supermassive BHs, with masses $\\sim 10^6 - 10^9\\,M_\\odot$ (see, e.g., Fryer \\& Kalogera 2001; Peterson 2003; Tegmark 2002). However, over the last few years, evidence has been mounting for the existence of {\\em intermediate-mass\\/} BHs (Miller \\& Colbert 2004). Although the observations remain controversial a variety of plausible formation scenarios for intermediate-mass BHs (IMBHs) have been proposed (e.g., van der Marel 2003). Primordial formation of massive BHs in the early universe has been discussed for many years (e.g., Carr 1993; Khlopov, Rubin \\& Sakharov 2002). Later on, when the first very massive metal-free (Pop~III) stars form, one naturally expects that, given the lack of significant wind mass loss, these stars may collapse to form IMBH remnants with masses up to $\\sim 10^2 - 10^3\\, M_\\odot$ (e.g., Heger et al.\\ 2003). Some of these IMBHs may be in binaries either formed through captures in dense environments (Wyithe \\& Loeb 2003) or remaining from primordial Pop~III binaries (Belczynski, Bulik \\& Rudak 2004a). Dynamical formation processes for IMBHs in Pop~II and Pop~I star clusters have also been the subject of many recent studies. The two most promising scenarios involve runaway collisions and mergers of massive main-sequence stars (Portegies Zwart \\& McMillan 2002; G{\\\"u}rkan, Freitag, \\& Rasio 2004) and successive mergers of stellar-mass BH (Miller \\& Hamilton 2002). In this paper, we focus on the populations of stellar BHs forming out of the most massive stars in Pop~II or Pop~I young stellar environments.  Stellar BHs form through core collapse of massive stars with masses $\\ga 20\\,M_\\odot$. If $N$ stars form with a standard Kroupa (Kroupa \\& Weidner 2003) initial mass function (IMF) between $M_{\\rm min} = 0.08\\, M_\\odot$ and $M_{\\rm max} = 150\\, M_\\odot$, we expect $N_{\\rm BH} \\simeq 5 \\times 10^{-4}\\, N$ black holes to have formed after $\\sim 10^7\\,$yr, when all stars massive enough to produce a BH have evolved. The goal of this paper is to characterize this initial population of BHs. For a population of {\\em single stars\\/}, this is merely a question of characterizing the metallicity-dependent relation between progenitor mass and final BH mass (Fryer \\& Kalogera 2001; Heger et al.\\ 2003). However, most stars, and, especially, most massive stars are expected to form in binary systems (Duquennoy \\& Mayor 1991). The BH formation processes can be affected significantly by the evolution of the progenitor star in a binary, especially if the initial orbital separation is $\\la 500\\,R_\\odot$ and the star can overflow its Roche lobe as it evolves. In addition, many BHs will retain binary companions, which can drastically affect their later evolution, their detectability as X-ray or gravitational-wave sources, and the way they interact with their environment.  The study we present in this paper is based on a population synthesis approach, in which a large number of single and binary stars are evolved according to parametrized evolutionary prescriptions. Stellar evolution is followed starting from a population of zero-age main-sequence (ZAMS) stars with specified metallicity. Given enough time the stars evolve to form remnants, either white dwarfs (not relevant for the massive stars considered here), neutron stars, or BHs. Both observations and more detailed theoretical calculations are used to constrain and parametrize the uncertain stages of evolution (e.g., supernova explosions and common envelope phases). A number of population synthesis codes of varying levels of sophistication are currently being used to study many astrophysical problems (e.g., Fryer, Woosley, \\& Hartmann 1999; Nelemans, Yungelson \\& Portegies-Zwart 2001; Hurley, Tout \\& Pols 2002; Pfahl, Rappaport, \\& Podsiadlowski 2003; Belczynski, Kalogera \\& Bulik 2002, hereinafter BKB02). Here we use the most recent version of the {\\tt StarTrack} population synthesis code (BKB02; Belczynski et al. 2004b) to study BH formation in young metal-rich (Pop~I) and metal-poor (Pop~II) populations. This work is complementary to the previous study of the oldest binary BHs, descendants of metal-free Pop~III stars, by Belczynski et al.\\ (2004a).  Our results can be applied directly to the modeling of young starburst populations, with ages $\\sim 10^7-10^8\\,$yr. For those systems, we have derived complete synthetic sample catalogues of all BHs, single and in binaries, that have evolved from the initial burst of star formation. Our results can also be applied to the modeling of young star clusters (such as the super star clusters so prevalent in starbursts; see, e.g., de Grijs et al. 2003), although we do not take into account the effects of dynamical interactions in dense cluster cores (cf.\\ G\\\"urkan, Freitag \\& Rasio 2003; Ivanova et al.\\ 2004).  Characterizing the properties of a primordial BH population is also very important for many theoretical studies of old star clusters. Indeed many of these studies attempt to bypass the first $\\sim 10^7$yr of massive star evolution and use initial conditions that already contain BHs. For example, $N$-body simulations of globular clusters containing primordial BHs have been performed starting with a small number of identical single BHs of mass $10\\,M_\\odot$ (e.g., Fregeau et al.\\ 2002; Portegies Zwart \\& McMillan 2000). One goal of our work is to provide more realistic initial conditions for those studies. The further dynamical evolution of BHs in dense star clusters could play a key role in many problems of great current interest, such as understanding the formation of IMBHs and ultra-luminous X-ray sources (Miller \\& Hamilton 2002) and predicting the merger rate of BH binaries detectable by gravitational-wave detectors (Portegies Zwart \\& McMillan 2000).  Our paper is organized as follows. In \\S\\,2 we describe the evolutionary model used in the population synthesis calculations, including details about BH formation. In \\S\\,3 we present the results of our calculations. First we present the results for our reference model, then we follow with the description of a number of alternative models. Finally, in \\S\\,4 we summarize our main results. ",
        "conclusions": "We have calculated the evolution of the massive stars found in young stellar environments, such as starbursts. The final products of massive star evolution, single and binary BHs, were then studied. All our results were discussed taking into account the many model uncertainties. A number of alternative calculations with varied initial conditions and evolutionary parameters were performed and presented. We also supplied the necessary data to normalize our results to any given total mass of a starburst galaxy or star cluster, with arbitrary choice of initial binary fraction. The calibrated results may then be used as part of initial conditions for realistic $N$-body simulations of dense stellar clusters that include primordial BHs.  Soon after the initial star formation burst, most BHs are found as single objects, although a significant fraction of BHs are also found in binaries. A number of single BHs are formed as the end product of binary evolution, either through binary disruption following a SN explosion or through a merger following a dynamically unstable RLOF episode. The most common binary BHs are  BH--MS and BH--BH systems. The period distribution of binaries containing BHs is usually bimodal, with a majority of systems in the long period peak ($P \\sim 10^4 - 10^6$\\ days). These wide binaries would be ``soft'' if placed in a dense stellar environment (e.g., a globular cluster) and they would then be disrupted following any strong interaction with another passing star or binary, thereby further enhancing the population of single BHs. The remaining, short-period BH binaries ($P \\sim 1-100$\\ days) would instead undergo hardening and evolve, over many relaxation times, to produce a population of very compact binaries that could eventually merge through GR emission.  The typical BH masses are found to be within the range $7-25\\,\\msun$, both for single and binary BHs. However, the single BHs formed through binary mergers can reach masses as high as $\\sim 80 \\msun$. Since most mergers are assumed in our models to be accompanied by significant mass loss, BHs formed through binary evolution (without any dynamical interactions) could in principle reach even higher masses, up to $\\sim 100 \\msun$ (in the absence of significant merger-induced mass loss). This result has many important implications. First, some ultra-luminous X-ray sources might be explained by a $\\sim 100 \\msun$ {\\em stellar\\/} BH accreting from a lower-mass companion. This would require the capture of a new companion (most likely through an exchange interaction with another binary), but no dynamics would be involved in the BH formation (cf.\\ Kalogera, King, \\& Rasio 2004). Second, these most massive stellar BHs may act as seeds for the formation of true IMBHs (with masses $\\gtrsim 1000 \\msun$) that could reside at the centers of some dense star clusters (Gebhardt, Rich, \\& Ho 2002; Gerssen et al.\\ 2002; Miller \\& Hamilton 2002). Third, a broader mass range for the tightest BH--BH binaries (possibly undergoing further hardening through dynamical interactions in a dense star cluster) will modify predictions for the gravitational-wave signals detectable by laser-interferometer instruments such as LIGO and VIRGO (Flanagan \\& Hughes 1998)."
    },
    "0404/astro-ph0404542_arXiv.txt": {
        "abstract": "We report on X-ray spectroscopic observations with {\\it XMM/Newton} of the ultracompact, double white dwarf binary, GP Com. With the Reflection Grating Spectrometers (RGS) we detect the L$\\alpha$ and L$\\beta$ lines of hydrogen-like nitrogen (N VII) and neon (Ne X), as well as the helium-like triplets (N VI and Ne IX) of these same elements. All the emission lines are unresolved. These are the first detections of X-ray emission lines from a double-degenerate, AM CVn system. We detect the resonance (r) and intercombination (i) lines of the N VI triplet, but not the forbidden (f) line. The implied line ratios for N VI, $R = f/i < 0.3$, and $G = (f + i)/r \\approx 1$, combined with the strong resonance line are consistent with formation in a dense, collision-dominated plasma. Both the RGS and EPIC/MOS spectra are well fit by emission from an optically thin thermal plasma with an emission measure (EM) $\\propto ( kT / 6.5 \\; {\\rm keV})^{0.8}$ (model {\\it cevmkl} in XSPEC). Helium, nitrogen, oxygen and neon are required to adequately model the spectrum, however, the inclusion of sulphur and iron further improves the fit, suggesting these elements may also be present at low abundance. We confirm in the X-rays the underabundance of both carbon and oxygen relative to nitrogen, first deduced from optical spectroscopy by Marsh et al. The average X-ray luminosity of $\\approx 3 \\times 10^{30}$ ergs s$^{-1}$ implies a mass accretion rate $\\dot m \\approx 9 \\times 10^{-13} \\; M_{\\odot}\\; {\\rm yr}^{-1}$. The implied temperature and density of the emitting plasma, combined with the presence of narrow emission lines and the low $\\dot m$ value, are consistent with production of the X-ray emission in an optically thin boundary layer just above the surface of the white dwarf. ",
        "introduction": "The AM CVn stars are among the most compact binary systems known. Their optical spectra are dominated by broad helium emission lines originating in an accretion disk formed by mass transfer from a degenerate helium dwarf onto the primary white dwarf. Their orbital periods range from about 10 minutes to 1 hour.  They are natural laboratories for the study of poorly understood binary evolution processes, such as common envelope evolution.  Moreover, the absence of hydrogen allows for the study of accretion disks dominated by helium. They may also be a significant channel for the production of Type Ia supernovae and neutron stars via accretion induced collapse (for a review see Warner 1995).  GP Com is one of the better studied, nearby AM CVn systems. It has a 46.5 minute orbital period (Nather Robinson \\& Stover 1981), resides at a distance of $\\approx 70$ pc (Thorstensen 2003), and has been studied extensively in the optical and UV. Recent studies by Marsh (1999) and Morales-Rueda et al.(2003) have used optical and UV+optical spectroscopy to probe the dynamics of the system in great detail. In addition to helium, nitrogen emission lines have been seen in the UV (Lambert \\& Slovak 1981; Marsh et al. 1995), and optical (Marsh, Horne \\& Rosen 1991). These observations identified a nitrogen overabundance, relative to carbon and oxygen, as well as an underabundance of heavy elements (for example, Ca, Si and Fe). The low metallicity may be consistent with its suggested halo origin (see Giclas, Burnham \\& Thomas 1961). The evolution scenario favored by Marsh et al. (1991) posits an initially low metallicity for the progenitor stars. The CNO elements are subsequently produced within the primary star and then mixed throughout the secondary during a common envelope phase. The nitrogen abundance is then increased by CNO-cycle hydrogen burning within the secondary (ie. the mass donor). Finally, a number of ultracompact systems with neutron star primaries also show apparent neon enrichment (see Juett \\& Chakrabarty 2003; Schulz et al. 2001).  Accretion should make such systems soft X-ray sources, and indeed, several have been detected in the X-ray band, including the prototype AM CVn, CR Boo, and GP Com (van Teeseling \\& Verbunt 1994; Ulla 1995; and Eracleous, Halpern \\& Patterson 1991). The X-ray flux from accreting, non-magnetic white dwarfs is thought to originate in a boundary layer that shocks and decelerates the Keplerian flow, allowing it to eventually settle on the white dwarf. Early work suggested that the boundary layer should be optically thick--radiating in the soft X-ray band and extreme UV--at mass accretion rates above $\\dot m \\approx 1.6 \\times 10^{-10} M_{\\odot} \\; {\\rm yr}^{-1}$ (Pringle \\& Savonije 1979). At lower rates the boundary layer should become more radially extended and optically thin, producing a harder X-ray spectrum. Aspects of this basic scenario have been confirmed by more recent calculations (see, for example, Patterson \\& Raymond 1985; Narayan \\& Popham 1993; Popham \\& Narayan 1995).  To date, X-ray observations of most AM CVn systems have been made at relatively low signal to noise levels and with modest spectral resolution. For example, ROSAT observations of GP Com found the source at a flux level $\\approx 1.2 \\times 10^{-11}$ ergs cm$^{-2}$ s$^{-1}$ in the 0.2 - 2 keV band, and confirmed that it is variable. These observations also found that the 0.2 - 2 keV spectrum cannot be described by a single temperature component (van Teeseling \\& Verbunt 1994), but detailed modelling was problematic due to the relatively poor statistics and restricted bandpass.  In this Letter we summarize results of recent, high signal to noise X-ray spectral measurements of GP Com with {\\it XMM/Newton}.  These data represent perhaps the best X-ray spectral measurements of an AM CVn system to date, and also show the first X-ray emission lines from such a system. ",
        "conclusions": "The nitrogen lines in the X-ray band provide additional confirmation of the overabundance of nitrogen relative to carbon and oxygen deduced from the optical (Marsh et al. 1991). The high nitrogen abundance likely reflects the operation of CNO-cycle thermonuclear processing in the secondary (Marsh et al. 1991).  We find a nitrogen to helium ratio, by number, of $\\approx 4.8 \\times 10^{-3}$, which is moderately higher than the value derived by Marsh et al. (1991) of $2.7 \\times 10^{-3}$ (see their Table 4). Our derived oxygen to helium number ratio is, however, about a factor of 10 higher than that deduced from the optical. Although systematic modelling effects could be present, these results suggest that the X-ray emitting gas may have higher oxygen abundance than the accretion disk (which produces the optical emission lines). If the accretor were an oxygen-rich white dwarf, and some of its matter is mixed with the accreted material in the turbulent boundary layer, then this could perhaps increase the oxygen abundance relative to the donor material.  The neon lines provide the first definitive detection of an element heavier than oxygen in GP Com. Our derived neon mass fraction ($\\approx 0.0037$) is a factor of 2 larger than the solar value. Moreover, modelling of the optical spectra also suggests a neon enrichment of about 1.3 compared to solar abundances (Marsh et al. 1991). Neon is a by-product of helium burning under a wide range of conditions (see, for example, Clayton 1983; Iben \\& Truran 1983), however, the low carbon abundance is likely incompatible with helium burning having ocurred in the secondary. Therefore, this avenue for the neon production appears to be a dead-end.  Again, mixing of neon from the accretor (which likely has undergone helium burning) into the donor material in the boundary layer could perhaps increase the neon abundance.  It appears that gravitational settling and fractionation of neon (in the form of $^{22}Ne$) are also not relevant, since nitrogen is depleted in the production of this isotope, and we clearly measure a large nitrogen abundance (see, for example, Bildsten \\& Hall 2001).  Neon enhancements have been suggested for several ultracompact neutron star binaries; 4U 1543-624 and 2S 0918-549 (Juett \\& Chakrabarty 2003), and 4U 1627-67 (Schulz et al. 2001). For 4U 1543-624, Juett \\& Chakrabarty (2003) deduced Ne/O $\\approx 1.5$, by number, from studies of absorption edges in {\\it Chandra} high resolution spectra, however, we note that these measurements can be influenced by the unknown ionization structure in the circumbinary material.  We find a similar value of Ne/O = 1.4 (by number) for GP Com. Yungelson, Nelemans \\& van den Heuvel (2002) suggest fractionation of neon in the white dwarf donor as a possible explanation for the neon enhancements in the neutron star systems. This scenario again requires a helium burning donor and thus seems not to work in the case of GP Com.  The X-ray measurements provide a means to estimate the mass accretion rate, $\\dot m$, onto the white dwarf primary in GP Com. Assuming matter falls from the inner Lagrange point onto the primary, and that half the gravitational potential energy is dissipated in the accretion disk, the X-ray luminosity, $L_x$ can be approximated as (see, for example, Nelemans, Yungelson \\& Portegies Zwart 2004), \\begin{equation} L_x = \\frac{1}{2}\\frac{G M_1 \\dot m}{R} \\left (1 - \\frac{R}{R_{L_1}} \\right ) \\; , \\end{equation} where $M_1$, $R$, $R_{L_1}$, and $\\dot m$ are the primary mass, primary radius, the distance from the inner Lagrange point to the center of the accretor, and the mass accretion rate, respectively. Solving for $\\dot m$, and inserting the measured X-ray luminosity for $L_x$ gives, \\begin{equation} \\dot m = 3.48 \\times 10^{-13} \\frac{R_9}{(M_1/M_{\\odot})} \\left ( 1 - \\frac{R}{R_{L_1}} \\right )^{-1} \\; \\; M_{\\odot} \\; {\\rm yr}^{-1} \\; , \\end{equation} where $R_9$ is the radius of the primary in units of $10^9$ cm, and $R/R_{L_1} \\approx 0.2$ for a system like GP Com. For a primary mass, $M_1 = 0.5 M_{\\odot}$ (Morales-Rueda et al. 2003) and $R_9 = 1$ (Zapolsky \\& Salpeter (1969) we obtain $\\dot m \\approx 8.7 \\times 10^{-13} \\; M_{\\odot} \\; {\\rm yr}^{-1}$. If the accretor is a massive O/Ne/Mg white dwarf (suggested by the possibility of oxygen and neon mixing in the boundary layer), then the derived accretion rate drops by about a factor of 6.  This rate is substantially lower than the critical rate above which one expects an optically thick boundary layer (Pringle \\& Savonije 1979; Narayan \\& Popham 1993). In this regime the plasma should have a maximum temperature of $\\approx 10^8$ K, which is similar to the peak temperature of $6.5$ keV ($7.5 \\times 10^7$ K) deduced from our spectral modelling.  The presence of {\\it narrow} emission lines also supports spectral formation in a boundary layer. From the width of the N VII Ly$\\alpha$ line, we obtain an upper limit on the velocity dispersion of the emitting gas of $\\approx 250$ km s$^{-1}$ ($3\\sigma$) (This includes bulk and thermal motions).  This indicates an origin close to the white dwarf and not, for example, further out in the accretion disk, which would have much higher Keplerian velocities.  If the white dwarf accretor was synchronized with the orbit, then its rotational velocity would be $\\approx 23$ km s$^{-1}$. To have a rotational speed equal to our 250 km s$^{-1}$ limit, a $10^{4}$ km white dwarf would require a spin period of $\\approx 250$ seconds. Indeed, an upper limit to the Doppler (ie. thermal motion) width for the nitrogen line indicates a value close to t250 km s$^{-1}$, which supports a slowly rotating, or synchronized white dwarf. The narrowness of the X-ray lines also appears consistent with the narrow ``central spike'' component of the helium lines (Marsh 1999; Morales-Rueda et al. 2003), which almost certainly originate on the accreting white dwarf.  The XMM/Newton measurements provide a detailed look at the physics of the boundary layer in GP Com. Detailed comparisons with theoretical boundary layer models could provide interesting constraints on their temperature, density and rotational profiles. Deeper X-ray spectra would provide better temperature and density constraints from the helium-like triplets, and likely more lines of less abundant species could be detected."
    },
    "0404/astro-ph0404297_arXiv.txt": {
        "abstract": "Using adaptive optics on Keck and the VLT in the H- and K-bands, we have begun a project to probe the dynamics and star formation around AGN on scales of 0.1\\arcsec. The stellar content of the nucleus is traced through the 2.29\\micron\\ CO\\,2-0 and 1.62\\micron\\ CO\\,6-3 absorption bandheads. These features are directly spatially resolved, allowing us to measure the extent and distribution of the nuclear star forming region. The dynamics are traced through the 2.12\\micron\\ H$_2$ 1-0\\,S(1) and 1.64\\micron\\ [Fe{\\sc ii}] emission lines, as well as stellar absorption features. Matching disk models to the rotation curves at various position angles allows us to determine the mass of the stellar and gas components, and constrain the mass of the central black hole. In this contribution we summarise results for the two type~1 AGN Mkn\\,231 and NGC\\,7469. ",
        "introduction": "The ultraluminous infrared galaxy Mkn\\,231, which at 170\\,Mpc distance has $L_{\\rm bol}\\sim3\\times10^{12}$\\,L$_\\odot$, hosts an AGN that, with $M_{\\rm B} = -21.7$, is often classed as a QSO. However, it is now clear that a significant fraction of its luminosity infact originates in star formation, making Mkn\\,231 a key object for investigations into whether or not ULIRGs evolve into QSOs. The structural and kinematic properties of this and other late stage ULIRG mergers have been studied by \\cite{gen01} to investigate whether they might evolve later into (intermediate mass, $L_*$) ellipticals; and by \\cite{tac02} to test if, once rid of their gas and dust shells, they might be the progenitors of QSOs. As a sample, the ULIRGs appear to have elliptical-like properties: relaxed stellar populations with $r^{1/4}$ radial luminosity profiles, and velocity dispersions of $\\sim$180\\kms\\ with only moderate rotation. Based on numerical simulations of mergers between gas rich spirals, this is what one might expect since most ULIRGs, including Mkn\\,231, exhibit the huge tidal tails typical of such mergers. However, Mkn\\,231 presents a puzzle since its stellar velocity dispersion of 115\\kms\\ is rather small, yielding both a low bulge mass and a low M$_{\\rm BH}$, which results in a highly super-Eddington AGN luminosity. Adaptive optics observations on the Keck~II telescope, which are summarised here and described fully in \\cite{dav04b}, have resolved this problem.  \\begin{figure} \\centerline{\\psfig{file=mkn231.eps,width=13.5cm}} \\caption{Mkn\\,231. Left: spatial profiles of the continuum (upper) and CO absorption `flux' (lower) showing that the nuclear star forming region is resolved. Overplotted in the lower panel are $r^{1/4}$ (dotted) and exponential (dashed) profiles. Right: logarithm of the CO profile as a function of $r^{1/4}$, showing that the exponential (dashed) profile is a better match than the $r^{1/4}$ (dotted), and hence that the stars reside in a disk rather than a spheroid.} \\label{fig:mkn231} \\end{figure}  If the mean stellar type dominating the nuclear stellar H-band continuum does not vary too much, the 1.63\\micron\\ CO\\,6-3 bandhead absorption `flux' traces the spatial extent of the stars. Fig.~\\ref{fig:mkn231} (left) shows that this profile is different to the continuum, which has a strong narrow core originating in dust heated by the AGN. Distinguishing between the de~Vaucouleurs $r^{1/4}$ and exponential fits to the stellar profile is crucial for our understanding of the geometry of the star forming region. Hence in Fig~\\ref{fig:mkn231} (right) we plot the logarithm of the profiles as functions of $r^{1/4}$, on which scaling the de~Vaucouleurs profile appears as a straight line. It is then clear that an $r^{1/4}$ law does not match the data at larger radii, and that its $r_e$ is inconsistent with the scales on which the star formation is seen. On the other hand, an exponential profile with $r_e\\sim0.2$\\arcsec\\ does match the data, indicating that the stars lie in a disk rather than a spheroid.  A nearly face-on ($i$=$10^\\circ$) molecular disk is already known to exist in Mkn\\,231. We need to consider whether our result that the stars lie in the same disk is consistent with the stellar kinematics of \\cite{tac02}, who found $\\sigma=115$\\kms\\ and $V_{\\rm rot}/\\sigma=0.2$. Using the properties of the molecular disk at $r=0.6$\\arcsec\\ ($V_{\\rm rot}\\sin{i}=60\\kms$, M$_{\\rm dyn}=12.7\\times10^9$\\,M$_\\odot$, scale height 23\\,pc; \\cite[Downes \\& Solomon 1998]{dow98}) gives a mean velocity dispersion perpendicular to the disk plane of 80\\kms. Putting in also an exponential profile as above, accounting for seeing, and repeating the observations of \\cite{tac02} yields $\\sigma=107$\\kms\\ and an apparent $V_{\\rm rot}/\\sigma=0.3$. This means that a nearly face-on stellar disk can -- in the right circumstances -- masquerade as a spheroid.  The fraction of H-band light due to stars is found from the equivalent width of the CO\\,6-3 bandhead, $W_{\\rm CO}$. The stellar luminosity one derives can then be used in starburst models, yielding a minimum mass (occuring when late type supergiants dominate the continuum, at an age of $\\sim$10\\,Myr) of $1.3\\times10^9$\\,M$_\\odot$ out to $r=0.6$\\arcsec. The upper limit to the age and mass is set by M$_{\\rm dyn}$. At 0.6\\arcsec, we find M$_{\\rm dyn}=6.7\\times10^9$\\,M$_\\odot$ (half of that derived by \\cite{dow98}, since we measure $V_{\\rm rot}\\sin{i}=40\\kms$). Accounting for the gas mass leaves at most $4.3\\times10^9$\\,M$_\\odot$ of stars. Hence, from starburst models, the maximum age of the stars is 120\\,Myr. This remarkably young age is supported by other observations of PAH features (\\cite[Rigopoulou \\etal\\ 1999]{rig99}), mid infrared emission lines (\\cite[Genzel \\etal\\ 1998]{gen98}), the infrared spectral energy distribution (\\cite[Verma 1999]{ver99}), and the radio continuum (\\cite[Carilli \\etal\\ 1998]{car98}).  Our results show that the stars we see lie in a disk rather than a spheroid, and so $\\sigma$ cannot be used to estimate M$_{\\rm BH}$. Being so young, it is likely that the stars formed {\\em in situ} in the nearly face-on gas disk, which is itself a product of the merger that created Mkn\\,231. ",
        "conclusions": "We have presented H- and K-band adaptive optics observations which clearly resolve the nuclear star forming regions in the centres of 2 type~1 AGN. Constraining starburst models using the fraction of stellar light determined from $W_{\\rm CO}$ and the dynamics measured from the H$_2$ 1-0\\,S(1) line (and, where possible, stellar features), indicates that the nuclear star forming regions are extremely young and constitute a significant fraction of the total mass on these scales. In Mkn\\,231 it is likely that the stars have formed in the molecular disk of effective radius 160--200\\,pc which has resulted from the merger of gas rich spirals; in NGC\\,7469 the star cluster lies inside a molecular ring of radius 65\\,pc."
    },
    "0404/astro-ph0404178_arXiv.txt": {
        "abstract": "Seyfert narrow line regions (NLR) have emission line ratios which are remarkably uniform, displaying only $\\sim$0.5  dex variation between different galaxies. Existing models have been unable to explain  these observations without the introduction of ad hoc assumptions, geometrical restrictions or new parameters. Here we introduce a new model: dusty radiation pressure dominated photoionization, which provides a natural self-regulating characteristic leading to an invariance of the spectrum over a very wide range ($>100$) of ionization parameter. The dusty model is able to reproduce both the range and the absolute value of the observational line ratios not only in the standard optical diagnostic diagrams but also in UV diagnostic plots, providing an explanation to the problem in NLR observations. ",
        "introduction": "The emission lines of active galaxies have often been used in conjunction with models to constrain the physical and ionizati on structure of the emitting regions.  In particular ratio diagrams or line diagnostic diagrams prove to be an excellent visual aid in interpreting the emission line data. For example, the line diagnostic diagrams by \\cite{VO87} are capable of distinguishing three different groups of emission line galaxies: those excited by starbursts and two excited by an active nucleus - the Seyfert narrow line regions (NLRs) and the low ionization nuclear emission-line regions (LINERs). These diagrams are additionally interesting in that they show that the emission from Narrow Line Regions (NLR) is remarkably uniform, with only $\\sim$0.5 dex variation between Seyferts and less within individual galaxies. This uniformity of the spectral properties has since been confirmed in much larger samples (eg.\\ \\cite{VeronCetty2000}).  The standard paradigm proposes that the NLR are excited by photons originating at or near a compact nuclear source (see, eg.\\ \\cite{Ost89}) having a smooth featureless power-law, or broken power-law EUV ionizing spectrum. Within this model, the clustering of the observed line ratios within such a restricted domain of parameter space presents a problem, as it requires an approximately constant ionization parameter\\footnote{The ionization parameter is a measure of the number of ionizing photons against the hydrogen density ($U = S_{\\star}/\\mathrm{n}_{\\mathrm{H}}c$).} of $U\\sim10^{-2}$, whereas $U$ should be free to take on any value.  Modellers have been therefore forced to make the arbitrary (and possibly unphysical) assumption that the gas density in the ionized clouds must fall exactly as the inverse square of the distance from the nucleus.  In order to account for this failings of the previous standard models we have proposed a new paradigm for the photoionization of the NLR clouds, that of dusty, radiation pressure dominated photoionization (\\cite{DG02}). ",
        "conclusions": "First introduced in \\cite{DG02}, the dusty, radiation pressure dominated photoionization model, through the stagnation of the ionization parameter at large values, provides a simple explanation for the small variation of observed Seyfert NLR ratios. This stagnation is due to the effects of dust opacity and radiation pressure upon dust and is characteristic to these models. The significant point is that the dusty model is able to do this over both optical and UV ratios, without depending upon large variations in other parameters such as density or metallicity. These results not only provide an explanation for what has not been a fully understood observation for years but also provide ways in which to understand further the processes involved in the NLR and extended NLR of AGN."
    },
    "0404/astro-ph0404452_arXiv.txt": {
        "abstract": "Interstellar turbulence has implications for the dispersal and mixing of the elements, cloud chemistry, cosmic ray scattering, and radio wave propagation through the ionized medium. This review discusses the observations and theory of these effects. Metallicity fluctuations are summarized, and the theory of turbulent transport of passive tracers is reviewed.  Modeling methods, turbulent concentration of dust grains, and the turbulent washout of radial abundance gradients are discussed.  Interstellar chemistry is affected by turbulent transport of various species between environments with different physical properties and by turbulent heating in shocks, vortical dissipation regions, and local regions of enhanced ambipolar diffusion. Cosmic rays are scattered and accelerated in turbulent magnetic waves and shocks, and they generate turbulence on the scale of their gyroradii. Radio wave scintillation is an important diagnostic for small scale turbulence in the ionized medium, giving information about the power spectrum and amplitude of fluctuations. The theory of diffraction and refraction is reviewed, as are the main observations and scintillation regions. ",
        "introduction": "One of the most important developments in the field of interstellar gas dynamics during the last half-century was the renewed perception that most processes and structures are strongly affected by turbulence. This is a paradigm shift unparalleled in many other fields of astronomy, comparable perhaps to the discovery of extrasolar planets and cosmological structure at high redshift.  Interstellar turbulence and its implications were commonly discussed in the 1950s, but without the range of observations that are available today, the theory of the interstellar medium (ISM) drifted toward increasingly detailed models based largely on a preference for thermal and dynamical equilibria. Standing apart were two subfields that continued to include turbulence as the primary driver of all observations: radio scintillation and cosmic ray transport.  The connections between these subfields and the prevailing picture of the ISM were mostly ignored because of the large discrepancy in scales. Consequently, two other small-scale processes got lost in the equilibrium paradigm of the large-scale models: mixing of the elements and molecular chemistry. In this review, we discuss these four subfields in some detail. All have a well-established observational base going back several decades, but because of the complexity of turbulence and the infancy of the relevant theory, all of them are now in a state of rapid evolution.  Our previous review {\\it Interstellar Turbulence I} (this volume), emphasized observations in the dense neutral ISM and discussed in detail the various theoretical approaches to this field. ",
        "conclusions": "The turbulence that is observed directly on resolvable scales in the ISM also has important effects on very small scales, down to the level of atomic diffusion, mixing, and viscosity, and continuing far below to the thermal ion gyroradius.  The resolvable turbulence in the neutral medium was reviewed in the first part of this series ({\\it Interstellar Turbulence I}), along with general theory and simulations. The smaller scale implications were reviewed here.  Turbulence helps to disperse the elements made in supernovae and other stellar sources by stretching and folding the contaminated gas until the gradient length approaches the collision mean free path. Then atomic diffusion does the final step of interatomic mixing. This mixing and homogenization may occur partly at the source, in the shock fronts and contact discontinuities around the expanding flow following dynamical instabilities, and it may occur partly in the ambient ISM after the expansion subsides.  Observations suggest that the dispersion in elemental abundances among stars inside clusters and field stars of the same age, and the abundance differences between HII regions and diffuse clouds, are only a few percent up to perhaps 30\\%. Observations of clusters also suggest that most of the mixing occurs within several times $10^8$ years. This homogenization has to occur between the injection scale of supernovae and superbubble explosions and the star formation scale containing on the order of a solar mass. The corresponding range of spatial scales is $\\sim100$ for gas at the ambient density. Homogenization cannot be significant on scales much larger than a superbubble, because then it would diminish the galactic radial abundance gradient over a Hubble time.  These gradients enter the problem in another way too because turbulence mixes the gas over galactic radius to increase the elemental dispersion locally at the same time as it mixes this gas locally and leads to homogenization.  While the distribution of elemental abundances appears to be fairly narrow, with a dispersion on the order of 10\\%, it may also have a fat tail caused by occasional odd stars and diffuse clouds with very different abundances. This tail is reminiscent of other fat tails in the distribution functions for turbulent media, and these other tails seem to originate with intermittency. In the case of elemental dispersion, this means that the homogenization process is spotty so some regions survive for long periods of time with very little mixing.  Several methods for studying turbulent mixing have been employed. Because of the wide range of scales involved, these studies often employ simplifying assumptions that appear to capture the essential physics. They include artificial stochastic velocity fields and closure methods using moment equations. Direct numerical simulations of ISM elemental dispersion have only just begun.  The theory suggests that a turbulent medium mixes passive scalars like elemental abundances faster than a Gaussian velocity field because of the long-range correlations that are associated with turbulence.  Turbulence affects interstellar chemistry by mixing regions with different properties, heating the gas intermittently on the viscous scale, and enhancing ion-neutral collisions in regions with strong magnetic field gradients. Turbulent mixing is more far-reaching than diffusive mixing because of long-range correlated motions in a turbulent flow. Such mixing spreads out each chemical species over a large radial range inside a cloud. Turbulent heating promotes temperature-sensitive reactions inside otherwise cold clouds. Applications to the formation of CH$^+$ and OH in diffuse clouds look promising.  Chemical reaction networks that include these turbulent processes are only beginning to understand some of the implications, and direct simulations of chemistry in a turbulent medium are limited to only a few studies so far.  The scattering and acceleration of cosmic rays depends strongly on the presence of turbulence in the ISM. The scale for this turbulence is the gyroradius, which is less than the collisional mean free path for most cosmic rays, which have energies less than 1 GeV.  Thus, scattering relies entirely on magnetic irregularities in collisionless plasma turbulence. Cosmic ray acceleration is by two processes, both of which involve turbulence: The first-order Fermi mechanism accelerates cosmic rays by cycling them through shock fronts where they pick up a relative velocity kick comparable to the shock speed at each passage.  This cycling occurs because magnetic irregularities on each side of the front scatter the out-streaming cosmic rays back into the front.   The second-order Fermi mechanism accelerates cosmic rays through turbulent diffusion. Each collision with a randomly moving magnetic irregularity turns the cosmic ray around with a reflection speed in the moving frame that is comparable to the incident speed.  Over time, the result is a transfer of energy from the turbulence to the cosmic rays.  Cosmic ray scattering occurs in several ways. Fast moving particles can interact weakly but resonantly with numerous magnetic irregularities, and successive interactions of this type can randomly change the pitch angles of their helical motions along the field. This eventually leads to a complete reversal in the direction of motion. Cosmic rays can also interact strongly with large magnetic irregularities, as would occur in shock fronts and at the edges of dense cloud complexes. These interactions change the particle directions significantly each time.  The nature of cosmic ray diffusion in both momentum and space is not understood well because the structure and strength of the important magnetic irregularities are not observed directly.  If MHD turbulence is highly anisotropic on the scale of cosmic ray scattering, with transverse irregularities much stronger than parallel as suggested by theory, then resonant scattering processes during motions along the mean field can be very weak. The structure of magnetic waves below the mean free path is also unclear, as the usual fast and slow magnetosonic modes do not exist.  Cosmic rays can also generate turbulence by streaming instabilities following particle-wave resonances or by fire-hose and mirror instabilities that operate even without resonances. The expected anisotropy of Alfv\\'en wave turbulence diminishes the first of these mechanisms significantly, however.  The others are problematic because they require strong cosmic ray pressure anisotropies.  As a result, the impact of cosmic rays on turbulence is currently not understood.  Radio wave scintillation is indirect evidence for interstellar plasma turbulence. These radio observations span a very wide range of spatial scales through a combination of diffraction and refraction effects.  The scales are mostly below the collision mean free path, and they are far below the limits of angular resolution. The first of these limits makes it difficult to understand the origin of the density irregularities. Unlike the turbulence that scatters cosmic rays, which requires only magnetic irregularities and no density structure, the turbulence that causes scintillation requires small-scale density irregularities in the ionized medium.  The associated magnetic irregularities are not observed, and the connection to cosmic ray scattering is unclear, even though the length scales are about the same. The origin of density structures below the collision mean free path is unknown. Atomic diffusion should smooth them out on a sound crossing time unless magnetic field irregularities hold them in place. In that case they could be the result of slight temperature variations, with the cooler regions having higher electron densities, all divided up and mixed together by transverse magnetic motions. The nature of these motions below the mean free path is unclear, because, as mentioned above, the usual fast and slow compressional modes do not exist, nor do the usual thermal and pressure-regulated processes.  The second of the two limits on spatial scale imply that the geometrical properties of scintillation turbulence are difficult to observe. The scintillation is clearly anisotropic, but whether it is in sheets or filaments, for example, is unknown.  The importance of scintillation observations for studies of ISM turbulence is that they give the power spectrum of electron density fluctuations fairly accurately.  This is usually close to the Kolmogorov spectrum of incompressible turbulence. Rarely are the spectra so steep that the medium can be interpreted as a superposition of sharp edges, like shock fronts. One recent observation with fairly high precision obtained the relatively shallow Iroshnikov-Kraichnan power spectrum, leading to the conjecture that the slope varies from region to region. As discussed in {\\it Interstellar Turbulence I}, this variation may arise from a variation in the relative strength of the magnetic field compared to the turbulent motions, with the Iroshnikov-Kraichnan spectrum present in regions of relatively strong fields.  Observations suggest that scintillation arises in a distributed fashion from the ambient ionized medium and also from discrete high-density places like the edges of local bubbles, ionized molecular clouds, and HII regions.  These discrete regions are likely to be highly turbulent and they also have a juxtaposition of hot and cool gas, which is necessary for isentropic mixing and electron density structure.  There are evidently many uncertainties in the nature of ISM turbulence on small scales, even though the evidence for this turbulence is pervasive. Part of the problem is that none of the features of this turbulence have been observed directly: not the densities, magnetic fields, temperatures, or motions. Still, the density irregularities are revealed indirectly through scintillation, the magnetic field irregularities through cosmic ray scattering, the temperature fluctuations through chemistry, and the motions through elemental and chemical mixing. An additional problem is that many small scale effects of ISM turbulence rely on details of the theory that are independent of the usual scaling relations, such as viscous heating and elemental diffusion, which arise at the bottom of the cascade in the neutral medium.  Turbulence in the ionized medium is also below the collisional mean free path, where pressure and thermal effects are relatively unimportant. Moreover, the small scale ISM processes are often strongly dependent on the large scale processes, such as turbulent shock formation, energy and metal injection, and galactic-scale gradients.  This means that direct simulations of small scale turbulence are impossible without simplifying assumptions about the large-scale medium -- assumptions that require more knowledge about ISM turbulence on the large scale than is presently available.  While the observations and theory of ISM turbulence have come a long way from the first efforts in the 1950s, the details of this new information have led to a growing awareness that the complete problem is far too large to solve any time soon.  We rely on future generations of astronomers and physicists to continue this work, and hope that they find this field as intriguing and challenging as we do today."
    },
    "0404/astro-ph0404514_arXiv.txt": {
        "abstract": "We present an updated set of constraints for the progenitor of PSR\\,J0737-3039 and for the natal kicks imparted to pulsar~B taking into account both the evolutionary and kinematic history of the double neutron star. For this purpose, we use recently reported scintillation velocity measurements to trace the motion of the system in the Galaxy backwards in time as a function of the unknown orientation $\\Omega$ of the systemic velocity projected on plane of the sky as well as the unknown radial velocity $V_r$. The absolute limits on the orbital separation and the mass of pulsar~B's helium star progenitor just before its supernova explosion are $1.2\\,R_\\odot \\la A_0 \\la 1.7\\,R_\\odot$ and $2.1\\,M_\\odot \\la M_0 \\la 4.7\\,M_\\odot$. The kick velocity is constrained to be between 60\\,km\\,s$^{-1}$ and 1660\\,km\\,s$^{-1}$ and to be misaligned from the pre-SN orbital angular axis (which could be associated with pulsar~B's spin axis) by least $25^\\circ$ . We also derive probability distribution functions for the kick velocity imparted to pulsar~B and for the misalignment angle between pulsar~A's spin and the post-supernova orbital angular momentum for both isotropic and polar kicks. The most probable values of both quantities depend sensitively on the unknown radial velocity. In particular, tilt angles lower than $30^\\circ$--$50^\\circ$ tend to be favored for current radial velocities of less than $\\simeq 500$\\,km\\,s$^{-1}$ in absolute value, while tilt angles higher than $120^\\circ$ tend to be associated with radial velocities in excess of $\\simeq 1000$\\,km\\,s$^{-1}$ in absolute value. ",
        "introduction": "The recent discovery of the strongly relativistic binary pulsar (Burgay et al. 2003) which is also the first eclipsing double pulsar system found in our Galaxy (Lyne et al. 2004) has resparked the interest in the evolutionary history and formation of double neutron star (DNS) systems (Willems \\& Kalogera 2004; Dewi \\& van den Heuvel 2004; Willems, Kalogera \\& Henninger 2004). According to the standard formation channel for DNS binaries (e.g. Bhattacharya \\& van den Heuvel 1991), their progenitors evolve through a high-mass X-ray binary phase where the first-formed NS accretes matter from the wind of a high-mass companion. When the latter fills its Roche lobe, the extreme mass ratio triggers the formation of a common envelope which extracts orbital energy and angular momentum and causes the NS and donor core to spiral-in towards each other. If the inspiral can be stopped before the components coalesce, a tight binary is formed consisting of the first-formed NS and a helium star companion. After it exhausts its central helium supply, the helium star expands and a second, this time stable, mass-transfer phase is initiated which spins the NS up to millisecond periods. At the end of the mass-transfer phase the helium star's core explodes and forms the second NS.  In this paper, we present constraints on the properties of the progenitor of PSR\\,J0737-3039 just before the supernova (SN) explosion that gives birth to pulsar~B. In view of the recent scintillation velocity measurements, which were first presented at the meeting (Ransom et al., these proceedings) and updated shortly thereafter (Ransom et al. 2004), we update our results presented at the conference and account for the system's full kinematic history since the time of its formation. We furthermore present probability distribution functions (PDFs) for the kick velocity imparted to pulsar~B and, in reply to the call for theorists' predictions on PSR\\,J0737-3039, we predict the most likely misalignment angle between pulsar~A's spin and the post-SN orbital angular momentum as a function of $V_r$. ",
        "conclusions": "We have used the scintillation velocity measurements of Ransom et al. (2004) to derive the most up-to-date constraints on the progenitor of PSR\\,J0737-3039 and on the kick imparted to pulsar~B at birth as a function of the unknown orientation $\\Omega$ of the systemic velocity projected on plane of the sky and the unknown radial velocity $V_r$. We also derived the most probable pulsar kick velocity and spin tilt for both isotropic and polar kicks as a function of $V_r$. Once $\\Omega$ is measured in the coming year, it will be straightforward to use the results presented here to further constrain the natal kicks and the spin-tilt predictions."
    },
    "0404/astro-ph0404208_arXiv.txt": {
        "abstract": "s{BL Lac objects are an elusive and rare class of active galactic nuclei. For years their evolutionary behavior has appeared inconsistent with the trend observed in the population of AGN at large. The so-called ``negative'' evolution implies that BL Lacs were either less or fainter in the past. This effect is stronger for BL Lacs selected in X-ray surveys. We have investigated if one of the selection criteria, namely the flat-radio spectrum (imposed on the Radio-selected but not on the X-ray-selected samples), might explain the different evolutionary trend. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/hep-th0404253_arXiv.txt": {
        "abstract": "We study particle production in the tachyon condensation process as described by different effective actions for the  tachyon. By making use of invariant operators, we are able to obtain exact results for the density of produced particles, which is shown to depend strongly on the specific action. In particular, the rate of particle production remains finite only for one of the actions considered, hence confirming results previously appeared in the literature. ",
        "introduction": "\\label{intro} One of the most intriguing features of String Theory is the tachyon condensation process and its possible description in terms of a tachyon effective theory. Finding an effective action for the tachyon field is a difficult task. Nevertheless, in certain situations (like the decay of non-BPS D-branes) some aspects of the string dynamics can be described by an effective field theory action involving only the tachyon and massless modes (see for example Refs.~\\cite{Lambert:2002hk,Lambert:2001fa}). \\par The form of the effective action for the tachyon depends on the choice of the region in field space where it can be valid. One such a region corresponds to the neighbourhood of the perturbative string vacuum (in which the tachyon $T=0$) where one can reconstruct the tachyon effective action from string $S$-matrix. In this case, we expect the effective action to be of the form \\be \\label{Tperturb} S=-\\int \\d t\\,\\d^p x\\, \\left( {1\\over 2}\\,\\eta^{\\mu\\nu}\\,\\partial_{\\mu}T\\,\\partial_{\\nu}T -{\\mu^2\\over 2}\\,T^2+g\\,T^4+\\ldots\\right) \\ . \\ee where $\\eta_{\\mu\\nu}$ is the Minkowski metric (Greek indices run from $0$ to $p$), $\\mu$ and $g$ are constants and the dots represent higher order terms. \\par A second possibility is trying to reconstruct the effective action near some exact conformal points. One such a conformal point is a time-dependent background that represents an exact boundary conformal field theory~\\cite{Sen:2002an,Sen:2002in,Sen:2002nu,Gutperle:2003xf}, \\be T=f_0\\,e^{\\mu\\,x^0}+{\\bar f}_0\\,e^{-\\mu\\,x^0} \\ . \\ee A special case is then the ``rolling tachyon'' background \\be \\label{rol} T=f_0\\,e^{\\mu\\,x^0} \\ , \\ee where $\\mu^2=1$ or $1/2$ in the bosonic and supersymmetric case respectively, which would be described by the homogeneous Lagrangian~\\cite{Kutasov:2003er} \\be \\label{kutt} L=-\\frac{1}{1+\\frac{T^2}{2}}\\, \\sqrt{1+\\frac{T^2}{2}-\\left(\\partial_0 T\\right)^2} \\ . \\ee If the above admits a straightforward covariant generalization, one obtains \\be \\label{tachDBIx} S=-\\int \\d t\\, \\d^p x\\, {1\\over 1+{T^2 \\over 2}}\\, \\sqrt{1+{T^2 \\over 2} +\\eta^{\\mu\\nu}\\,\\partial_{\\mu}T\\,\\partial_{\\nu}T} \\ , \\ee which has just the form of a ``DBI'' (Dirac-Born-Infeld) action \\be \\label{tachDBI} S=-\\int \\d t\\,\\d^p x\\, V(\\tilde{T})\\, \\sqrt{1+\\eta^{\\mu\\nu}\\,\\partial_{\\mu}\\tilde{T}\\, \\partial_{\\nu}\\tilde{T}} \\ , \\ee where $\\tilde T=\\sqrt{2}\\,\\sinh\\left(T/\\sqrt{2}\\right)$ and the potential is given by \\be V(\\tilde T)=\\left(\\cosh\\frac{\\tilde T}{\\sqrt{2}}\\right)^{-1} \\ . \\label{Vp} \\ee Even if the rolling tachyon (\\ref{rol}) is an exact solution of the equation of motion that arises from the action (\\ref{tachDBIx}), it is not completely clear whether such an action should be used for the description of field theory space around the perturbative vacuum $T=0$. In particular, in Refs.~\\cite{Berkooz:2002je,Felder:2002sv,Frolov:2002rr}, the stability of the the classical time-dependent solution of Eq.~(\\ref{tachDBIx}) was analyzed, showing that fluctuations become large in a very short time interval and hence they might significantly change the classical evolution. One can also interpret the huge growth of fluctuations as the creation of a large number of fluctuation modes in the time-dependent background of the classical solution \\cite{Kluson:2003rd}. We can then expect that at some time near the beginning of the tachyon condensation the density of the number of particles created reaches the string density and the linearized approximation, in which fluctuations are assumed small, breaks down. \\par Finally, one can also consider as a starting point a tachyon effective action of the form \\be \\label{Tbos} S=-\\int \\d t\\,\\d^p x\\, \\frac{1}{1+T}\\sqrt{1+T+ \\frac{\\eta^{\\mu\\nu}\\,\\partial_{\\mu}T\\,\\partial_{\\nu}T}{T}} \\ , \\ee which arises when one considers the tachyon of a D-brane in bosonic String Theory and was derived in Ref.~\\cite{Kluson2003} following the proposal of Ref.~\\cite{Kutasov:2003er}. \\par We will analyze the production of particles for the three actions reviewed above by making use of Lewis' invariant operators~\\cite{Lewis}. This technique yields exact results and will allow us to go beyond the usual linear approximation in all the cases under consideration. Our analysis follows the same line as that of Ref.~\\cite{Klusonlast}. In particular, we shall restrict our study of particle production to the fluctuation modes with initial momentum $k^2>\\mu^2$ at infinite past. Modes with $k^2<\\mu^2$ exponentially grow even at the beginning of the tachyon condensation and this may significantly change the classical evolution. The analysis of these modes is a challenging task beyond the scope of this paper and even for the modes with a positive frequency $\\Omega_k ^2=k^2-\\mu^2$ at far past we will get some interesting results. ",
        "conclusions": "\\setcounter{equation}{0} We have computed the density of particles produced in the early stages of the tachyon condensation. The use of invariant operators allows us to obtain exact results and to show that the production rate strongly depends on the form of the action. The DBI action~(\\ref{tachDBIx}) and the Bosonic Action~(\\ref{Tbos}) produce divergent densities in the early stages of the tachyon condensation, whereas the Perturbative Effective Action~(\\ref{Tperturb}) seems to yield a regular behavior when the linearized approximation is satisfied. This suggests that such an action should be regarded as the most suitable candidate to describe the process of tachyon condensation. On the other hand one should note that backreaction effects have been neglected in our approach. Any conclusions therefore need to be confirmed by an analysis that takes into account the backreaction, and our results must be considered as preliminary until a more thorough (and very likely numerical) analysis is performed. This is left for future research beyond the scope of this paper.    \\label{conc}"
    },
    "0404/astro-ph0404044_arXiv.txt": {
        "abstract": "{The spectroscopic analysis of 117 serendipitous sources in the HELLAS2XMM 1df (1 degree field) survey is described. Of these, 106 sources, of which 86\\% have a spectroscopic redshift, are used to evaluate the fraction of X-ray absorbed (log N$_H$$>$22) Active Galactic Nuclei (AGN) in the 2--10 keV flux range 0.8--20$\\times$10$^{-14}$ erg cm$^{-2}$ s$^{-1}$. This fraction turns out lower than what is predicted by two well known Cosmic X-Ray Background synthesis models, and the discrepancy is significant at the 99.999\\% level. This result consolidates the findings recently obtained by other authors. In the flux interval explored, the data are consistent with an intrinsic distribution of the absorbing columns (flat per decade above log$N_H>$21) independent of luminosity and redshift, together with an AGN luminosity function evolving purely in luminosity. It is shown that, on the other hand, extrapolation to lower fluxes fails to reproduce the results inferred from the Chandra Deep Field North survey. It is found that about 40\\% of the high luminosity sources in our sample have best fit logN$_H$$>$22, and the surface density of these X--ray obscured QSOs can then be estimated at about 48 per square degree, at the flux limit of $\\sim10^{-14}$ erg cm$^{-2}$ s$^{-1}$ of the HELLAS2XMM 1df survey.  As a side issue, 5 or 6 out of 60 sources, that is about 10\\%, identified with broad line AGN, turn out to be affected by logN$_H$$>$22 absorption. ",
        "introduction": "After the success of ROSAT (Hasinger et al. 1998) in resolving about 75\\% of the X-ray background (XRB) in the 0.5--2 keV band into sources largely associated with Active Galactic Nuclei (AGN), the satellites Chandra and XMM-Newton achieved a similar result, up to at least 85\\% of the XRB, in the 2-10 keV band (Mushotzky et al. 2000, Giacconi et al. 2001, 2002, Hasinger et al. 2001; Alexander et al. 2003; see also Moretti et al. 2003 and references therein). The combination of the results in the two bands provides also the observational support for the intuition by Setti \\& Woltjer (1989) that the XRB could be explained by a dominant contribution of AGN, affected by photoelectric obscuration in different proportions over a wide range of gas columns N$_H$. This suggestion led to several attempts, all formally successful, to synthesize the XRB starting from somewhat different assumptions about the AGN Luminosity Function (LF) and its cosmological evolution, and N$_H$ distributions (e.g. Comastri et al. 1995, Gilli et al. 2001, Wilman \\& Fabian 1999, Miyaji et al. 2000, Ueda et al. 2003). In this context, an important issue, which is being explored with increasingly more detailed X--ray spectral analysis and spectroscopic identification of the optical counterparts, is the fraction of sources with different intrinsic N$_H$ as a function of their flux. This approach provides very strong constraints, especially when accompanied by the study of the LF performed using the same data (e.g. Ueda et al. 2003). The present limits are set by progressively poorer statistics in the X-ray spectra and in the optical spectroscopic identification as one goes to fainter sources. Thus, while a treatment as just outlined, based on the full ensemble of sources utilized by Fiore et al. (2003, hereafter Paper IV), is deferred to La Franca et al. (in prep.), this paper aims to exclusively present the information on the fraction of sources affected by different levels of X-ray obscuration, down to a limit in F(2-10 keV) of about 10$^{-14}$ erg cm$^{-2}$ s$^{-1}$ (corresponding to about 35\\% of the XRB), extracted from the HELLAS2XMM 1df sample. This sample comprises 117 sources, 93 of them (80\\%) with a spectroscopic redshift available.  The spectral counts extraction is described in Sect. 2, their best fit analysis in Sect. 3, the synthesis of the results in Sect. 4. Sect. 5 is devoted to a discussion of the results compared to XRB synthesis models, Sect. 6 to the conclusions. ",
        "conclusions": "Starting from 117 sources from the HELLAS2XMM 1df survey (Paper IV), after the exclusion of 1  source with N$_H$ unconstrained and 9 sources with unknown z and R$<$23, the spectroscopic analysis of the remaining 107 X-ray spectra (86\\% with spectroscopic redshift) led to the following main result. The fraction of the 106 sources with logN$_H$$>$22 in the flux interval 0.8--20$\\times$10$^{-14}$ erg cm$^{-2}$ s$^{-1}$  is inconsistent, at the 99.999\\% confidence level, with the predictions of two well known XRB synthesis models, one by Comastri et al. (1995), the other by Gilli et al. (2001, their model B). This result consolidates the discrepancy also found by other authors in this flux interval, as mentioned in the previous section.  As an exercise for the Comastri et al. (1995) model, leaving unchanged all other assumptions, the adoption of a simple and different intrinsic distribution of the source percentage per decade of N$_H$, from log$N_H$=20 to 25 (which is consistent with the results from the present sample, see Fig. 4), leads to a much better agreement down to 10$^{-14}$ erg cm$^{-2}$ s$^{-1}$, but fails to reproduce the much larger percentage of absorbed sources, derived from the CDFN survey, in the 10$^{-15}$ to 10$^{-14}$ erg cm$^{-2}$ s$^{-1}$  flux interval.  A study encompassing a much wider flux range, with a sufficiently large sample of objects (such as the one used in Paper IV), that should tackle simultaneously the problems of the shape and evolution of the LF, of the N$_H$ distribution as a function of luminosity and cosmic epoch and eventually the XRB synthesis (with an approach akin to that followed by Ueda et al. 2003), goes beyond the scope of the present paper, and will be the subject of La Franca et al. (in prep.).  An important result, which basically confirms what was found in Paper IV, is that in our sample at least 28\\%, most likely about 40\\% of AGN with logL$_{2-10keV}>44$ (that is of the QSO) are obscured in X-rays (log$N_H>22$). This fraction can be translated, taking into account the sky coverage, into a surface density of highly obscured QSO of $\\sim48$ deg$^{-2}$, at the flux limit of $\\sim10^{-14}$ erg cm$^{-2}$ s$^{-1}$ of the HELLAS2XMM 1df survey.  As a side issue, note that in the sample studied, while a value of the parameter log(X/O) much greater than unity confirms itself to be strongly indicative of high obscuration in high luminosity AGN, as shown in Paper IV, there are 5 or 6 out of 60, that is about 10\\% of sources, with logN$_H$$>$22, that are optically classified as AGN1, in agreement with a previous finding by Page et al. (2003; see also Brusa et al. 2003, Akiyama et al. 2003). Notably they are all concentrated at logL$_{2-10keV}>44$. Among various possibilities, it is pointed out that variability may be one of the causes of this inconsistency."
    },
    "0404/astro-ph0404014.txt": {
        "abstract": "{We report on two photometric monitoring campaigns of Very Low Mass (VLM) objects in the young open cluster around $\\sigma$\\,Orionis. Our targets were pre-selected with multi-filter photometry in a field of 0.36\\,sqdeg. For 23 of these objects, spanning a mass range from 0.03 to 0.7\\,$M_{\\odot}$, we detect periodic variability. Of these, 16 exhibit low-level variability, with amplitudes of less than 0.2\\,mag in the I-band, which is mostly well-approximated by a sine wave. These periodicities are probably caused by photospheric spots co-rotating with the objects. In contrast, the remaining variable targets show high-level variability with amplitudes ranging from 0.25 to 1.1\\,mag, consisting of a periodic light variation onto which short-term fluctuations are superimposed. This variability pattern is very similar to the photometric behaviour of solar-mass, classical T Tauri stars. Low-resolution spectra of a few of these objects reveal strong H$\\alpha$ and Ca-triplet emission, indicative of ongoing accretion processes. This suggests that 5-7\\% of our targets still possess a circumstellar disk. In combination with previous results for younger objects, this translates into a disk lifetime of 3-4\\,Myr, significantly shorter than for solar mass stars. The highly variable objects rotate on average slower than the low-amplitude variables, which is expected in terms of a disk-locking scenario. There is a trend towards faster rotation with decreasing mass, which might be caused by shortening of the disk lifetimes or attenuation of magnetic fields. ",
        "introduction": "\\label{intro}  Open clusters are ideal environments to study stellar properties and evolution, because they contain a homogeneous population of objects with known distance, metallicity, and age. Recent deep surveys have unveiled the population of several open clusters far down into the substellar regime. Examples are the surveys by Moraux et al. (\\cite{mbs03}), Pinfield et al. (\\cite{phj00}) and Zapatero Osorio et al. (\\cite{zrm99}, \\cite{zrm97}) in the \\object{Pleiades}, and the work of Barrado y Navascu\\'es et al. in \\object{IC2391} (\\cite{bsb01}) and \\object{$\\alpha$\\,Per} (\\cite{bbs02}). In two cases, namely in the clusters around \\object{$\\sigma$\\,Ori} (Zapatero Osorio et al. \\cite{zbm00}) and the Orion \\object{Trapezium Cluster} (Lucas \\& Roche \\cite{lr00}), even isolated planetary mass objects were found. Thanks to these surveys, large samples of Very Low Mass (VLM) objects are known today, among which we conveniently will subsume all objects with masses below $0.4\\,M_\\odot$, including very low mass stars, brown dwarfs and free-floating planetary mass objects. All objects with masses below this $0.4\\,M_\\odot$ limit are thought to be fully convective (Chabrier \\& Baraffe \\cite{cb00}), giving a physical motivation for this definition of VLM objects.  Wide-field photometric monitoring is a powerful tool to investigate properties of large object samples: If an object exhibits asymmetrically distributed surface features, e.g. magnetically induced star spots, its flux will be modulated with the rotation period. Hence, period search in the lightcurve allows the determination of the (projection-free) rotation period. The amplitude of the periodicity, in turn, depends on the properties of the star spots, allowing conclusions about surface activity processes. Pointing to dense open cluster fields, such monitoring campaigns become very efficient, as one can register many objects contemporaneously within one field of view.  Compared to solar-mass stars, there is a significant lack of known rotation periods for VLM objects. Photometric monitoring studies delivered a small number of periods for evolved VLM stars and Brown Dwarfs (Bailer-Jones \\& Mundt \\cite{bm99}, \\cite{bm01}, Tinney \\& Tolley \\cite{tt99}, Mart\\'{\\i}n et al. \\cite{mzl01}, Clarke et al. \\cite{ctc02}, Gelino et al. \\cite{gmh02}). Only three periods are known for VLM objects with ages between 50 and 100\\,Myr (Mart\\'{\\i}n \\& Zapatero Osorio \\cite{mz97}, Terndrup et al. \\cite{tkp99}). Two periods for VLM members of $\\sigma$\\,Ori were published by Bailer-Jones \\& Mundt (\\cite{bm01}). Recently, Joergens et al. (\\cite{jfc03}) report on 5 measured periods for Brown Dwarfs and VLM stars in the very young ChaI star forming region. Our own study in the young open cluster IC4665 (age 36\\,Myr) delivered several rotation periods for VLM objects (Eisl{\\\"o}ffel \\& Scholz \\cite{es02}).  Observations of very young VLM objects are of special interest, because they can deliver clues about their formation process. Recent results suggest that objects in the substellar regime form similar to stars. Several authors detected typical T Tauri star phenomena on VLM objects, e.g. outflow processes (Fern\\'andez \\& Comer\\'on \\cite{fc01}) and mid-infrared excesses attributed to the presence of a circumstellar disk (Natta \\& Testi \\cite{nt01}, Apai et al. \\cite{aph02}, Jayawardhana et al. \\cite{jas03}). Moreover, L\\'opez Mart\\'{\\i} et al. (\\cite{les03}) and Liu et al. (\\cite{lnt03}) demonstrate a correlation between H$\\alpha$ emission and mid-infrared excess for Brown Dwarfs down to $0.02\\,M_\\odot$, which they interpret as an indication for ongoing accretion.  Solar-mass T Tauri stars show various types of photometric variability. Active accretion processes often manifest themselves by large amplitude variations showing several possible periods (e.g., Fern\\'andez \\& Eiroa \\cite{fe96}, Bouvier et al. \\cite{bck95}, Herbst et al. \\cite{hmw00}). If VLM objects undergo a T Tauri phase as well, they should exhibit a similar photometric behaviour. Hence, photometric monitoring can deliver an independent contribution to the ongoing debate about VLM object formation.  The $\\sigma$\\,Ori cluster is an appropriate target for such a study, because it is relatively nearby (350\\,pc, B\\'ejar et al. \\cite{bzr99}), has negligible extinction ($E_\\mathrm{B-V}=0.05$, B\\'ejar et al. \\cite{bzr99}) and an age of 3\\,Myr (Zapatero Osorio et al. \\cite{zbp02}). The extended work by B\\'ejar et al. and Zapatero Osorio et al. revealed a rich VLM population. We monitored this cluster in two photometric time series. Complementary observations, presented in Sect.\\,\\ref{sel}, identified the cluster members in the time series field. We report on the monitoring campaigns in Sect.\\,\\ref{mon} and the time series analysis in Sect.\\,\\ref{tsa}. The following sections describe the results of out lightcurve analysis: We first establish two origins for the observed variability (Sect.\\,\\ref{inter}). In Sect.\\,\\ref{rot}, we then concentrate on the investigation of the rotation periods. Section\\,\\ref{accr} contains the discussion of low-resolution spectra for highly variable objects. Finally, in Sect.\\,\\ref{conc}, we present our conclusions.  %__________________________________________________________________ ",
        "conclusions": "\\label{conc}  We present the first photometric variability study for VLM members of the young $\\sigma$\\,Ori cluster. With multi-filter photometry, 135 cluster member candidates were identified, including 90\\% VLM objects ($M<0.4\\,M_{\\odot}$). This preliminary member list still contains about 30\\% contaminating field stars and thus needs spectroscopic verification.  We monitored these targets in two I-band time series in January and December 2001. The first campaign, using the 2-m Tautenburg Schmidt telescope, covered 110 candidates, spanning masses from 0.02 to 0.5\\,$M_{\\odot}$. The second time series, obtained with the 1.23-m telescope on Calar Alto, registered 52 targets spanning a mass range from 0.02 to 1.4\\,$M_{\\odot}$. 27 Objects were observed in both runs. After reduction and relative calibration, we obtained differential lightcurves for all targets. We found that the difference image analysis package from G{\\\"o}ssl et al. (\\cite{gr02}) improves the precision of the Calar Alto lightcurves by several mmag.  Time series analysis was focused on the period search, but we included a test to detect irregular variations as well. We examined the variability of the targets based on the scattering of the lightcurves and found that the very young targets contain significantly more variable objects than field stars. The fraction of variable objects as well as the variability amplitude does not significantly change over the entire mass range.  We detected periodic variability for 14 CA and 13 TLS targets, with periods spanning from 4 to 240 hours. Four targets show a periodicity in both campaigns. The periodically variable candidates clearly fall into two groups: 16 of them show low-level variability which is in most cases well-approximated by a sine wave. We argue that this type of variability is most probably caused by magnetically induced photospheric spots co-rotating with the objects. We see evidence for the spot evolution on VLM objects, since most of these targets show the periodicity in only one campaign.  The lightcurves of the remaining objects have amplitudes $>$0.25\\,mag and show obvious deviations from the sine shape. Their photometric behaviour is very similar to those of classical T Tauri stars. The quasi-periodic oscillation is therefore most probably caused by hot spots formed by matter flow from an accretion disk. Thus, our results imply the existence of accretion disks around VLM stars and Brown Dwarfs. The fact that most of the highly variable objects show a near-infrared colour excess confirms this finding. Low-resolution spectroscopy of a subsample of these targets assures their cluster membership and their youth. We find strong emission features, in particular a dominating H$\\alpha$ emission line, characteristic for accreting objects. The extension of T Tauri behaviour to VLM stars and Brown Dwarfs strongly supports star-like formation scenarios for these objects.  Highly variable and thus accreting objects rotate on average more slowly than low-level variable targets. This is expected in terms of a disk-locking scenario, where magnetic coupling between star and disk removes angular momentum during the T Tauri phase. In agreement with previous publications, we find that strong accretors are rare in the $\\sigma$\\,Ori cluster, but they certainly exist among VLM stars as well as Brown Dwarfs. Comparison with studies of younger clusters and more massive stars give evidence that disk dissipation timescales are significantly decreased with decreasing object mass.  In agreement with previous publications, we find a tendency of faster rotation towards decreasing object masses, possibly caused by decreasing effectiveness of angular momentum removal. This could either be explained with shorter disk lifetimes or with weaker magnetic fields, as inferred from the low activity level of our targets."
    },
    "0404/astro-ph0404102_arXiv.txt": {
        "abstract": "Orbital period changes are an important diagnostic for understanding low mass X-ray binary (LMXB) accretion-induced angular momentum exchange and overall system evolution. We present our most recent results for the eclipse timing of the LMXB \\exo. Since its discovery in 1985 it has apparently undergone three distinct orbital period ``epochs\", each characterized by a different orbital period than the previous epoch. We outline the orbital period behavior for \\exo\\ over the past 18 years and discuss the implications of this behavior in light of current theoretical ideas for LMXB evolution. ",
        "introduction": "There are eight currently known low mass X-ray binary (LMXB) systems that undergo full or partial eclipses. Such systems are important for the study of evolution in LMXBs because the eclipse edges provide timing markers that make possible the systematic observation of orbital period changes. The best studied orbit in this group is that of \\exo\\ which has been X-ray active since 1985 \\citep{pwgg86}. This long timeline has allowed an unprecedented look at its orbital dynamics \\citep[see][and references therein]{whw+02}. The emerging picture for the orbital period behavior, however, is anything but the expected smooth variations in \\Porb\\ and \\Porbdot\\ based on theoretical calculations done to date \\citep[e.g., see][and references therein]{prp02}. Rather, the observed period changes are discontinuous across multiple distinct epochs, and large apparent changes in \\Porb\\ of the order seconds can be observed on timescales as short as one orbit. The magnitude of the observed changes in orbital period are much larger than expected from LMXB evolutionary theory. This is likely an indication that the observed variations in \\Porb\\ are short timescale effects of angular momentum redistribution in the system and are masking the underlying long term orbit evolution. The Rossi X-Ray Timing Explorer (RXTE) satellite has allowed us to monitor systematically the orbit of \\exo\\ in an effort to delineate and understand the process of angular momentum exchange and its effects on the \\exo\\ orbit. In this conference paper we report our recent progress. ",
        "conclusions": "The X-ray eclipse timings for \\exo\\ appear to show three distinct orbital periods in three successive epochs along with significant intrinsic period jitter in these epochs. Furthermore, inspection of Figure~\\ref{fig-exorxteusaomc} shows that during the early part of the RXTE epoch (until approximately March 2000) the period jitter is especially prominent. After this date the orbit period changes again by $\\sim$-6.7 ms nearly returning to the value characteristic of the EXOSAT$-$GINGA epoch. Such abruptly changing orbit periods are observed in a number of Algol binaries [see the discussion in \\citep{simo99} and references found therein] although with larger amplitude in ${\\Delta \\Porb}/{\\Porb}$. For Algol systems ${\\Delta \\Porb}/{\\Porb} \\sim 10^{-5}$ whereas in \\exo\\ we find ${\\Delta \\Porb}/{\\Porb} \\sim 10^{-6}$. \\citep{hall91} attributed the abrupt period changes in Algol systems to magnetic activity in a convective secondary star inducing changes in its quadrupole moment and altering the orbital angular momentum. If this explanation is applicable to the period changes in \\exo\\ \\citep[see also][]{hwc97} then around the March 2000 magnetic activity in the secondary was altered in some manner changing the secondary's influence on the system orbital angular momentum distribution. Before this date the orbit period jittered around a mean value but the jitter is reduced after this date. This may imply that magnetic activity in the secondary also was reduced, perhaps as a result of magnetic cycling similar to the 11-year cycle in the Sun. Small changes in radius of the secondary are predicted by the magnetic activity explanation of Algol orbit period changes put forward by \\citet{ap87} due to the changing magnetic pressure support for the outer layers of the secondary star. The data in Figure~\\ref{fig-exoecldurvsomc} \\begin{figure} \\includegraphics[height=4.5in,angle=270.0]{f3.ps} \\caption{ The durations of all RXTE observed eclipses plotted as a function of observed \\OMC\\ residual for \\exo. The points at increasing negative \\OMC\\ values are the same eclipses observed after the dramatic change in period near MJD 51710 in Figure~\\ref{fig-exorxteusaomc}. \\label{fig-exoecldurvsomc}} \\end{figure} suggest that the radius of the occulting star decreased across the MJD 51710 boundary. Totality duration is $\\sim$7 seconds shorter after MJD 51710 than before MJD 51710. However, a spherically symmetric decrease of 1-2\\% in the radius of the secondary star is unlikely and other non-spherical changes in the radius of the secondary must be considered as well.  The measurements shown in Figures~\\ref{fig-exoallomc},~\\ref{fig-exorxteusaomc}, and~\\ref{fig-exoecldurvsomc} give a magnitude and direction for the abrupt orbital period changes and the magnitude and direction of the abrupt change in the duration of X-ray totality. Any attempt to understand these data will have to pose some sort of almost instantaneous change in the binary system parameters with an accompanying redistribution of the orbital and spin angular momenta rather than a slow change in the system because of the effects of mass exchange.  \\begin{theacknowledgments} This research is supported by the Office of Naval Research, the NASA Astrophysical Data Program, and the NASA RXTE Guest Observer Program. \\end{theacknowledgments}"
    },
    "0404/astro-ph0404428_arXiv.txt": {
        "abstract": "We have observed the persistent but optically unidentified X-ray source X1908+075 with the PCA and HEXTE instruments on {\\it RXTE}. The binary nature of this source was established by \\citet{WRB} who found a 4.4-day orbital period in results from the {\\it RXTE} ASM.  We report the discovery of 605 s pulsations in the X-ray flux.  The Doppler delay curve is measured and provides a mass function of 6.1 $M_\\sun$ which is a lower limit to the mass of the binary companion of the neutron star.  The degree of attenuation of the low-energy end of the spectrum is found to be a strong function of orbital phase.  A simple model of absorption in a stellar wind from the companion star fits the orbital phase dependence reasonably well and limits the orbital inclination angle to the range $38\\arcdeg - 72\\arcdeg$.  These measured parameters lead to an orbital separation of $\\sim60 - 80$ lt-s, a mass for the companion star in the range 9-31 $M_\\sun$, and an upper limit to the size of the companion of $\\sim 22~R_\\sun$.  From our analysis we also infer a wind mass loss rate from the companion star of $\\gtrsim 1.3 \\times 10^{-6}~M_\\sun$ yr$^{-1}$ and, when the properties of the companion star and the effects of photoionization are considered, likely $\\gtrsim 4 \\times 10^{-6}~M_\\sun$ yr$^{-1}$. Such a high rate is inconsistent with the allowed masses and radii that we find for a main sequence or modestly evolved star unless the mass loss rate is enhanced in the binary system relative to that of an isolated star.  We discuss the possibility that the companion might be a Wolf-Rayet star that could evolve to become a black hole in $10^4$ to $10^5$ yr.  If so, this would be the first identified progenitor of a neutron star--black hole binary. ",
        "introduction": "X1908+075 is an optically unidentified, highly absorbed, and relatively faint X-ray source that appeared in surveys carried out with instruments on the {\\em Uhuru}, {\\em OSO 7}, {\\em Ariel 5}, {\\em HEAO-1}, and {\\em EXOSAT} satellites.  The early detections and position determinations are summarized by \\citet[hereafter WRB]{WRB}. They conclude that the position of X1908+075 is likely to be within the overlapping region of the error box of a source detected in an {\\em Einstein} IPC image and one of an array of {\\em HEAO 1} A-3 position ``diamonds'', and is thus known with an accuracy of $\\sim 1'$. Inspection of the POSS plates within the source error box reveals no optical counterpart down to magnitude 20.  This is consistent with the heavy optical extinction implied by the interstellar hydrogen column density of $\\sim 4 \\times 10^{22}$ atoms cm$^{-2}$ measured on the basis of the low-energy absorption in the X-ray spectrum \\citep[see][]{ss85}.  The intensity of X1908+075 has been monitored for the past eight years using the All-Sky Monitor (ASM) aboard the {\\it Rossi X-Ray Timing Explorer} ({\\it RXTE}), and has typically been in the range 2-12 mCrab in the 2-12 keV energy band.  WRB analyzed data from the first three years of operation of the ASM, and thereby discovered a 4.4 day periodicity in the X-ray intensity.  The periodic component of the intensity variations is clearly energy dependent in both strength and detailed form.  At most orbital phases, the variation is roughly sinusoidal while a relatively sharp dip forms the minimum in the 5-12 keV band.  These characteristics suggest that the modulation is produced by a varying amount of absorption along the line of sight as the source moves through the stellar wind of a massive companion star. The hard X-ray spectrum led WRB to also suggest that this source could be an X-ray pulsar.  The possibility that X1908+075 might be an X-ray pulsar led us to carry out a set of pointed observations with the PCA and HEXTE instruments on {\\it RXTE} in late 2000 and early 2001.  The data revealed the presence of strong X-ray pulsations at a period of 605 seconds.  We were also able to detect Doppler delays in the pulse arrival times.  However, because the number of independent high-quality pulse arrival times obtained from this data set was small, it proved difficult to unambiguously disentangle orbital effects from intrinsic changes in the pulse period.  The latter could, in principle, be quite large for a neutron star rotating with a period as long as 600~s \\citep[see, e.g.,][]{bild97,del01}.  Therefore, we obtained additional observations of the source with {\\it RXTE} during late 2002 and early 2003.  In this paper we report the results of our analysis of the {\\it RXTE} observations of X1908+075.  The pointed observations are described in \\S 2.  A pulse timing analysis is described in \\S 3.  We present an accurate pulse period and a determination of the spin-down of the neutron star due to accretion and magnetic torques.  The orbital Doppler delay curve is measured, thereby confirming the 4.4-day period found in the ASM X-ray light curve.  The resultant mass function is $6.1~M_\\sun$, indicating that the pulsar does indeed orbit a massive companion star.  The results of an orbital-phase-dependent spectral analysis are presented in \\S 4.  We detect a very pronounced modulation of the low energy attenuation as a function of orbital phase.  In \\S 5 we model this modulation by absorption in a spherically symmetric stellar wind, whereby we obtain constraints on the orbital inclination, on the properties of the companion star, and on the stellar wind.  We discuss the implications of our results in \\S 6, including the possibility that this system may be the progenitor of a neutron star--black hole binary. ",
        "conclusions": "We have found 605 s pulsations in the X-ray intensity of X1908+075. Doppler shifts of the pulse frequency, and changes in the intensity and low-energy attenuation that cyclically recur at the previously known orbital period of 4.4 days allow us to measure orbital parameters and to conclude that the system contains a highly magnetized neutron star orbiting in the wind of a massive companion star.  The observed variations in the shape of the spectrum and in the overall intensity indicate that the X radiation from the neutron star is both absorbed and scattered in the wind.  We have estimated the absorption on the basis of a crude spectral model that assumes that any attenuation is caused by photoelectric absorption in neutral (unionized) gas with solar element abundances.  In this model, we did not include the effects of the photoionization of the wind by the X-rays.  We can now use our estimates of the wind column density and the X-ray luminosity to roughly estimate the ionization parameter $\\xi = L_X/n r^2$ \\citep*{tts69} at a typical location in the binary system, where $n$ is the number density of atoms in the wind and $r$ is the distance from the X-ray source.  If we take the orbital radius $a$ as a characteristic distance within the binary system, i.e., we set $r = a$, and for a density use $n \\sim N_{H_{\\rm wind}}(\\phi = 0)/a$, we obtain $\\xi = L_X/(N_H a) \\sim$ 20--60 ergs cm s$^{-1}$. Under these conditions we would expect that most atoms would lose their outer electrons, but atoms of carbon, oxygen, etc. would retain at least their K-shell electrons \\citep{tts69,km82}.  \\citet{kk84} have estimated the equilibrium conditions of gases photoionized by continuum X-ray spectra.  They consider the condition of the gas as a function of a modified ionization parameter $\\Xi = L_X/(4\\pi c n r^2 kT)$ where $c$ is the speed of light, $k$ is Boltzmann's constant and $T$ is the gas temperature.  For a characteristic location in the X1908+075 system, we have $\\Xi \\sim 8 (T/10^5 {\\rm K})^{-1}$.  For each of the two spectral shapes that \\citet{kk84} consider in detail, one can roughly sketch the parameters of a self-consistent solution taking into account that the actual density will be higher than that estimated on the basis of the neutral matter absorption cross section. The parameters of the solutions in the two cases are roughly the same: $T \\sim 4 \\times 10^4$ K, $\\Xi \\sim 10$, the X-ray opacity in the $\\sim$2--6 keV band is reduced by a factor of $\\sim 2$ from the value expected for cold neutral gas of cosmic abundances, and $N_{H_{\\rm wind}}(\\phi = 0) \\sim 1 \\times 10^{23}$ cm$^{-2}$.  Since, in this regime, the temperature is a steep function of $\\Xi$, it is likely to vary over a wide range at different places in the system.  A more realistic estimate of the physical properties of the wind would need to self-consistently consider the actual source X-ray spectrum, the wind velocity, the possibility of non-solar abundances, the position-dependent degree of ionization in the wind, the effects of the ionization upon the wind acceleration, and the transfer of radiation in the system including both the effects of scattering and absorption.  This is beyond the scope of this paper.  Using our fitted column densities and wind speeds estimated from the formula of \\citet{vink00}, we find the mass loss rate in the wind from the companion star to be $\\gtrsim 1.3 \\times 10^{-6} M_\\sun$ yr$^{-1}$ for all the cases with acceptable values of $\\chi^2$ and with companion star radii in the range expected for main sequence stars. For those cases where the estimated mass loss rate is not more than a factor of order 3 greater than an empirical prediction based on the stellar mass, temperature, and luminosity, we find the rate must be $\\gtrsim 2 \\times 10^{-6} M_\\sun$ yr$^{-1}$.  The depression of the X-ray opacity of the gas from ionization implies that the wind mass loss rate must be higher than this latter rate, likely, as discussed above, by a factor of two and possibly by a larger factor.  As noted in \\S 5, over the allowed region of the $M_c - R_c$ plane that we have identified, such mass loss rates are significantly higher than those predicted using the empirically-based mass loss rate prescriptions of Vink et al.  However, one should note that these prescriptions do not take into account any of the effects of the star being in a binary system, including the presence of a critical potential lobe.  Furthermore, we used the effective temperatures computed by Podsiadlowski for stars that evolve without mass loss and without being affected by any of the phenomena that occur in a binary system.  Thus, we cannot exclude the possibility that the companion star is on the main sequence.   Another possibility is that the companion is a Wolf-Rayet (WR) star. A WR star could well have a mass that is consistent with the value that we find for the companion, but its radius would be much smaller than that of a main-sequence star of comparable mass.  It would have a prodigious wind \\citep{nug00} that could reach or surpass the mass loss rate that we infer (with assumptions) for X1908+075.  According to \\citet{nug00}, WR stars of either type WN or WC and mass in the range $\\sim 9$--15 M$_\\sun$ have wind mass loss rates between 4 and $20 \\times 10^{-6}$ M$_\\sun$ yr$^{-1}$.  Of course, the winds of WR stars, which have little or no hydrogen, would have significantly larger X-ray absorption cross sections per unit mass than winds from unevolved stars.  If there is little hydrogen in the wind, we may have {\\em overestimated} the wind mass loss rate, but such a composition would nonetheless clearly indicate the WR nature of the companion.  If the companion star is indeed a WR star, then the X1908+075 system has potentially important implications, in general, for binary stellar evolution, and for the formation, in particular, of neutron star--black hole (NS-BH) systems.  If it is a WR star, then it is likely the He or CO remnant core of a star that was originally much more massive, possibly of mass $\\sim 23-35~M_\\sun$ \\citep[e.g.,][]{hurl00}.  In this case we expect the companion to undergo core collapse in $10^4$ to $10^5$ years and leave behind a stellar-mass black hole \\citep{brwn01}.  Thus, there is a possibility that the X1908+075 system is the progenitor of a neutron star--black hole binary where the neutron star formed first.  This would be significant in at least two ways.  First, in spite of the fact that six NS-NS binaries have been discovered \\citep[see, e.g.,][]{bur03,chmp04}, no NS-BH binaries have yet been found.  There are theoretical arguments \\citep{spn04,ppr04} which suggest that the latter binaries should be at least a factor of 10 less populous than their NS-NS cousins \\citep[but see][]{vt03}. An identified progenitor would help theorists better estimate the current population of NS-BH binaries.  Second, if NS-BH binaries are formed, the issue of which collapsed star forms first is also important to our understanding of binary stellar evolution. The most direct way of producing these systems is for the BH to form first \\citep[from the more massive component of the system; see, e.g.,][]{pzve97,blwb00,fry01,nele01,lbw02,prh02}, but there are also channels where the NS forms first, as discussed above \\citep[see, e.g.,][]{vt03}.  Since this paper was originally written, \\citet{MG04} have reported the results of near-infrared observations of stars in or close to the error box of X1908+075, and have found a star whose JHK magnitudes and colors, and H and K band spectra, suggest an O or B supergiant at $d \\sim 7$ kpc.  They believe that this star is likely to be the counterpart of X1908+075.  If confirmed, this identification would rule against a WR companion star.  We expect to obtain observations of X1908+075 with the Chandra X-ray Observatory, and to thereby obtain an X-ray position sufficiently accurate to secure the optical identification."
    },
    "0404/astro-ph0404334_arXiv.txt": {
        "abstract": "The SDSS has opened a new era for the study of AGN spectroscopic properties and how these depend on luminosity and time.  In this presentation we review some of the current issues and problems in studies of high-redshift quasars and Narrow-Line Seyfert 1 galaxies. Investigations employing SDSS will, in some situations, still have to pay attention to selection biases and other fundamental limitations of the data. ",
        "introduction": "The existing literature and pre-Sloan Digital Sky Survey (SDSS) data pertaining to the emission line properties of AGNs have numerous limitations in terms of sample selection criteria and homogeneity. These issues are particularly evident in studies of high redshift quasars, where the focus is on $z > 4$. In this contribution we discuss some of the outstanding questions in this area and how these connect to our understanding of other AGNs, including Narrow-Line Seyfert 1 (NLS1) galaxies. SDSS will lead to significant improvements in several ways relevant to these topics, but it is important to recognize some of the concerns and potential problems that we will still encounter in using SDSS in studies of this type. ",
        "conclusions": ""
    },
    "0404/astro-ph0404116_arXiv.txt": {
        "abstract": "In the early days of methanol maser discoveries the $9_{2}-10_{1}\\ \\rm{A}^{+}$ transition at 23.1~GHz was found to exhibit maser characteristics in the northern star-forming region W3(OH), and probable maser emission in two other sources.  Attention subsequently turned to the 6.6-GHz $5_{1}-6_{0}\\ \\rm{A}^{+}$ methanol maser transition, which has proved a valuable tracer of early high-mass star formation.  We have undertaken a new search for 23.1-GHz methanol masers in 50 southern star formation regions using the Parkes radiotelescope.  The target sources all exhibit class~II methanol maser emission at 6.6~GHz, with 20 sources also displaying maser features in the 107.0-GHz $3_{1}-4_{0}\\ \\rm{A}^{+}$ methanol line.  Strong emission at 23.1~GHz in NGC~6334F was confirmed, but no emission was detected in the remaining sources.  Thus the 23.1-GHz methanol masers are rare.  A maser model in which methanol molecules are pumped to the second torsionally excited state by radiation from warm dust can account for class~II maser activity in all the transitions in which it is observed.  According to this model the 23.1-GHz maser is favoured by conditions representing low gas temperature, high external dust temperature, low gas density, and high column density of methanol;  the scarcity of this maser indicates that such combinations of conditions are uncommon.  We have undertaken new model calculations to examine the range of parameters compatible with the upper limits on 23.1-GHz emission from our survey.  Further constraints apply in sources with upper limits to maser emission at 107.0~GHz, and the combination of data for the two transitions delineates a narrow range of gas density and methanol abundance if the dust temperature is 175~K or greater.  While the results are subject to the uncertainties of the chosen model, they may be applicable to the majority of methanol maser sites in the vicinity of newborn high-mass stars, in which methanol masers other than the 6.6- and 12.1-GHz transitions are not detected. ",
        "introduction": "Methanol masers at 6.6 and 12.1~GHz have been detected in several hundred star-forming regions (e.g. Caswell \\etal\\ 1995a, 1995b), where they are often closely accompanied by OH maser emission and continuum emission from a nearby ultra-compact \\HII\\ region.  These two exceptionally strong maser lines are the archetypal class~II methanol masers.  Known only in star-forming regions, they probe the early stages in the development of a massive star, which is otherwise shielded from view by the enveloping cloud of dust and gas.  In some sources the strong 6.6-GHz $5_{1}-6_{0}\\ \\rm{A}^{+}$ and 12.1-GHz $2_{0}-3_{-1}\\ \\rm{E}$ masers are accompanied by class~II maser emission in up to 20 additional methanol transitions.  The latter are considerably weaker but otherwise display similar characteristics to the strong maser transitions: narrow lines at the same velocity, often displaying multiple components, sometimes variable, with small source size implying high brightness temperature where this has been investigated.  The class~I methanol masers, in contrast, are observed in a different set of methanol lines, and are found offset from other indicators of recent massive star formation.  The $9_{2}-10_{1}\\ \\rm{A}^{+}$ line at 23.1~GHz was the first methanol maser identified outside the class~I source Orion~KL (Wilson \\etal\\ 1984), and was the first discovered maser of what later became known as the class~II variety.  The maser nature of the emission towards W3(OH) was confirmed by VLA measurements (Menten \\etal\\ 1985, 1988a), in which several velocity components were identified as emanating from regions with small angular size, with brightness temperature $>10^7$~K for the strongest feature of intensity 10~Jy and velocity -43.2~\\kms.  Emission spectra with similar narrow features in single-dish observations of NGC~7538 at 0.5~Jy (Wilson \\etal\\ 1984) and NGC~6334F at 52~Jy (Menten \\& Batrla 1989) were also identified as probable masers.  Many class~II methanol maser sources were subsequently identified at 6.6~GHz, but no widespread surveys for the 23.1-GHz transition have been published to date, although preliminary results of a VLA survey were presented by Vargas \\etal\\ (2000).  We report here a new search for 23.1-GHz class~II methanol masers in southern star formation regions.  There is considerable evidence that class~II methanol masers at different frequencies may be spatially coincident.  VLBI observations of W3(OH) show that the strongest 6.6- and 12.1-GHz features coincide within a maser spot size (0.001~arcsec) and also coincide in velocity (Menten \\etal\\ 1992).  Furthermore, Menten \\etal\\ (1988b) found that individual 12.1-GHz maser components in W3(OH) are concentrated towards the most intense 23.1-GHz emission centres (observed with lower resolution).  Such observations suggest that the various maser transitions are excited simultaneously, and so provide constraints on maser pumping models.  To date, only three maser sources have been identified at 23.1~GHz, all of which also exhibit maser emission in the 107.0-GHz $3_{1}-4_{0}\\ \\rm{A}^{+}$ transition.  The high resolution observations of W3(OH) show correspondence between the emission regions at these two frequencies (Menten \\etal\\ 1988a, Sutton \\etal\\ 2001); however, the peak emission at 107.0~GHz stems from the northern edge of the ultra-compact \\HII\\ region where the corresponding 23.1-GHz maser spot is weak, while the southern edge produces the 23.1-GHz peak together with weaker emission at 107.0~GHz.  This means that the 23.1-GHz line carries new information on the conditions in maser sources, consistent with the two transitions exhibiting different sensitivities to physical conditions in the models.  The 23.1-GHz transition stems from levels which are higher in energy than those involved in most other methanol maser lines, and so weak maser emission is unlikely to be confused by overlapping thermal emission, as can happen for some millimetre methanol maser transitions.  The class~II methanol masers are pumped by infrared radiation from warm dust in the model of Sobolev \\& Deguchi (1994).  The pumping proceeds via the \\vt=2 and \\vt=1 torsionally excited states.  This model successfully accounts for the predominance of the 6.6- and 12.1-GHz masers (Sobolev, Cragg \\& Godfrey 1997a), and the existence of many other weaker class~II maser transitions, including the 23.1 and 107.0-GHz lines (Sobolev, Cragg \\& Godfrey 1997b).  The weaker maser transitions can be very sensitive to the model conditions, with calculated brightness temperatures changing by many orders of magnitude as the model parameters are varied.  Many transitions become masers simultaneously in the model.  This provides a probe of the physical conditions in the maser region.  On the one hand it permits best-fitting model conditions to be estimated via multi-transition analysis, when masers at several frequencies are observed in the same source.  Such studies have been undertaken for three sources in which data for many transitions are available (Cragg \\etal\\ 2001, Sutton \\etal\\ 2001).  On the other hand, upper limits on nondetected maser lines can also set constraints on the model parameters.  Surveys at different frequencies have detected the weaker maser transitions in only a few of the class~II methanol maser sources; for example, the 107.0-GHz maser is the most common of the weaker maser lines, but is detected in only 25 sources of more than 175 surveyed (Val'tts \\etal\\ 1995, Val'tts \\etal\\ 1999, Caswell \\etal\\ 2000, Minier \\& Booth 2002).  Thus the absence of the weaker methanol maser transitions in the majority of class~II methanol maser sources can be used to define the range of conditions under which the 6.6-GHz masers usually develop.   \\begin{figure} \\centerline{\\epsfxsize=8.5cm\\epsfbox{fig1.eps}} \\caption{23.1-GHz methanol maser spectrum detected in 351.417+0.645 (NGC~6334F), with three Gaussians fitted (dotted trace).} \\label{fig:ngc_2} \\end{figure}   \\begin{table} \\caption{Gaussian components fitted to detected 23.1-GHz maser in 351.417+0.645 (NGC~6334F)} \\label{tab:gfit}  \\centerline{ \\begin{tabular}{ccc} \\hline Flux Density  &   Velocity      &     FWHM \\\\ (Jy)     &    (\\kms)      &    (\\kms)  \\\\ \\hline \\\\ 32.77 &  --10.39  &   0.24\\\\ 6.35 &  --10.55   &  0.36\\\\ 3.85 &  --11.01    & 0.20\\\\ \\end{tabular}}  \\end{table} ",
        "conclusions": "A new search for 23.1-GHz methanol masers in 50 southern star formation regions has confirmed the previously detected maser emission in NGC~6334F, but found no new masers at this frequency.  The scarcity of 23.1~GHz masers in the target sources, which all exhibit class~II methanol maser emission at 6.6~GHz, provides new constraints for models of the maser pumping.  Further constraints apply in sources with upper limits to maser emission at 107.0~GHz.  Combining data for the two transitions sets limits on the physical conditions in sources representative of the majority of methanol maser sites in the vicinity of newborn high-mass stars, subject to the uncertainties of the applied model.  Conditions are likely to be near the limits of 6.6~GHz maser activity, with gas density $10^{6.5}-10^{8.5}$~\\cc\\ near the upper threshold, and/or dust temperature $100-150$~K near the lower threshold."
    },
    "0404/astro-ph0404266_arXiv.txt": {
        "abstract": "{ Asteroseismology is found to be a excellent tool for detecting diffusion-induced helium gradients inside main-sequence A stars. Models have been computed for 1.6 and 2.0 M$_{\\odot}$ stars with pure helium diffusion, at different ages, so that the helium gradient lies at different depths inside the star. The adiabatic oscillation frequencies have been analysed and compared with those of a model without diffusion. Clear signatures of the diffusion-induced helium gradient are found in the so-called ``second differences\" : these frequency differences present modulations due to the partial reflexion of the sound waves on the layer where the helium gradient occurs. A tentative application to the roAp star HD60435, which presents enough detected oscillation frequencies for the test to be possible, is very encouraging. The results suggest the presence of a helium gradient inside the star, which is consistent with the idea that the triggering of the oscillations is due to the hydrogen $\\kappa$-mechanism. ",
        "introduction": "It is now widely recognized that the abundance anomalies observed in peculiar A stars are basically due to element diffusion. Those for which the radiative acceleration is larger than gravity (like most of the metals) are pushed upwards while the others,such as helium, diffuse downwards (Michaud \\cite{michaud70}, Vauclair {\\&} Vauclair \\cite{vauclair82}). The resulting atmospheric abundances depend on the competition between element diffusion and macroscopic motions like convection, mass loss, rotation-induced mixing and other kinds of mixing processes. Among the peculiar A stars, those which have large organized magnetic fields (the so-called Ap stars) are more complex than the others (like Hg-Mn stars) : they show evidence of abundance spots or rings modulated by the magnetic structure. Among the coolest Ap stars (effective temperatures between 6700K and 8700K) some oscillate with periods of a few minutes (the roAp stars) while the others seem stable (noAp stars). In roAp stars, the amplitudes of the oscillation modes are modulated according to the rotation period, which is explained in the framework of the oblique pulsator model (Kurtz \\cite{kurtz90}).  The helium $\\kappa$-mechanism was invoked in the past as a possible way to trigger the oscillations : a well-adjusted competition between helium settling and mass loss could increase the helium abundance in stellar atmospheres as observed in helium-rich stars (Vauclair \\cite{vauclair75}, Vauclair, Dolez {\\&} Gough \\cite{vauclair91}). However, recent papers (Dziembowski {\\&} Goode \\cite{dziembowski96}, Balmforth et al. \\cite{balmforth01}) show that, even in this case, helium is not able to destabilize the star. On the other hand they find that, in some cases, the oscillations in A stars can be triggered by the hydrogen $\\kappa$-mechanism. Such a process may be enhanced by the settling of helium, which diffuses downwards in the absence of strong mass loss.  In the present paper we study the asteroseismic signature of a helium gradient inside main-sequence A-type stars. As mentioned by Gough (\\cite{gough90}), rapid variations of the sound velocity inside a star lead to partial reflections of the sound waves, which may clearly appear as frequency modulations in the so-called ``second differences\" : \\begin{equation} \\delta_2\\nu_{n, l} = \\nu_{(n+1), l}+\\nu_{(n-1), l}-2\\nu_{n, l} \\end{equation} which are computed for the same value of the azimutal number $l$. The modulation period of the oscillations is twice the ``acoustic depth\" of the region where the feature occurs (i.e. the time needed for the sound waves to travel between this region and the stellar surface).  Such an effect has been extensively studied in the literature for stellar models in which helium settling is neglected. In this case, the modulations in the frequencies are mostly due to the HeII ionisation zone and to the edge between the convective and radiative zones (Monteiro {\\&} Thompson \\cite {monteiro98}, Roxburgh {\\&} Vorontsov \\cite {roxburgh01}, Mazumdar {\\&} Antia \\cite {mazumdar01}, Miglio et al. \\cite {miglio03})  We show here that the presence of a diffusion-induced helium gradient leads to a kink in the sound velocity with a very clear signature in the oscillation frequencies : asteroseismic observations of A-type stars can test helium diffusion and lead to a precise value of the acoustic depth corresponding to the position of the helium gradient inside the star.  For the present study, we neglect the effects of magnetic fields (the stars are assumed spherically symetric) and we do not compute the diffusion of heavy elements. Complete computation of element diffusion including the radiative acceleration on metals is a demanding work, especially for these stars where heavy element diffusion occurs in the atmosphere, so that radiative transfer should be solved in detail in the optically thin regions. This is out of the scope of the present paper, where we are interested in the structural changes induced by helium settling below the stellar surface and its influence on the oscillation frequencies. It has recently been shown (e.g. Richard, Michaud {\\&} Richer \\cite{richard01}) that in case of pure diffusion iron should accumulate deeper inside the stars, at temperatures around 200 000 K, where it can create a small convective zone. As the results we obtain here are very encouraging, new computations including heavy element diffusion should be done in the near future to see whether such a feature may be checked with asteroseismology.  We apply this test to the only roAp star in which enough modes have been observed, namely HD 60435 (Matthews et al. \\cite{matthews87}). The results are extremely encouraging and show the importance of detecting as many modes as possible in rapidly oscillating Ap stars. ",
        "conclusions": "\\subsection{Seismic signatures of helium gradients}   Stellar acoustic $p$ modes with low $l$ values can propagate deeply inside the stars. For this reason, they may be used to obtain information on the deep stellar structure. However, in the case of strong gradients in the sound velocity, which may be due to the boundary of a convective zone, to the helium ionization region, or to helium gradients, the waves are partly reflected : this creates modulations in the frequency values which may be clearly visible in the second differences. The modulation periods are equal to $ 2 t_s$ where $t_s$ is the time needed for the acoustic waves to travel between the surface and the considered region (acoustic depth), i.e. : \\begin{equation} t_{s}=\\int_{r_s}^{R}\\frac{dr}{c(r)} \\end{equation} where $c(r)$ is the sound velocity at radius $r$, and $r_s$ the radius of the considered region.  As the $l$ dependence of the oscillatory signal is small for high frequency modes of low degrees (Monteiro, Christensen-Dalsgaard {\\&} Thompson \\cite{monteiro94}), computational results for modes of different low $l$ values can be treated together. Here we show that the presence of diffusion-induced helium gradients lead to very clear signatures in the modulation of the second differences.   \\subsection{Models with helium diffusion}  We have chosen to study the evolution of 1.6 M$_{\\odot}$ and 2.0 M$_{\\odot}$  stellar models, with pure helium settling. The stellar evolution code we use is the ``Toulouse-Geneva code'' which has been described many times in the literature (e.g. Richard et al. \\cite{richard96}, Th\\'eado {\\&} Vauclair \\cite{theado03}). The physical input and parameters are those described in Richard, Th\\'eado {\\&} Vauclair (\\cite{richard04}): they include the most recent studies available for the equation of state, opacities and nuclear reaction rates. Helium diffusion is treated as described in these papers (see also Vauclair \\cite{vauclair03}). Metal diffusion is not included here and no mixing process or mass loss is taken into account. In real stars, there must be some mixing which prevents element diffusion from producing extreme abundance anomalies. Computations including mixing as well as metal diffusion will be done in the near future. The aim of the present paper is to study the precise signature on the oscillation frequencies of helium gradients inside stars.   \\begin{table}[t] \\catcode `\\*=\\active \\def*{\\phantom{0}} \\baselineskip12pt \\pagestyle{empty} \\centerline{TABLE 1} \\medskip \\centerline{Model parameters} \\bigskip  \\halign{# & # & # & # & # & # & # & # \\cr \\noalign{\\hrule} \\noalign{\\kern2pt} \\noalign{\\hrule} \\noalign{\\bigskip} & mass (M$_{\\odot}$) & age (Myrs) & *log($L/L_{\\odot}$) & ***$T_{eff}$ & *****$t_{x}$(sec) & ****$t_{s}$(sec) & ****$r_{s}/R$ \\cr \\noalign{\\medskip} \\noalign{\\hrule} \\noalign{\\bigskip} & 1.6* & **95 & ***0.84 & ***7936 & *******5400 & ******1480 & *****0.93 \\cr \\noalign{\\medskip} & 1.6* & *591 & ***0.88 & ***7585 & *******6920 & ******2700 & *****0.86 \\cr \\noalign{\\medskip} & 1.6* & 1600 & ***1.05 & ***7071 & ******11240 & ******5160 & *****0.80 \\cr \\noalign{\\medskip} & 1.6h & 1600 & ***1.06 & ***7352 & ******10227 & ********** & ********* \\cr \\noalign{\\medskip} & 2.0* & **63 & ***1.22 & ***9521 & *******5917 & ******1470 & *****0.94 \\cr \\noalign{\\medskip} & 2.0* & *649 & ***1.31 & ***8293 & ******10150 & ******4100 & *****0.85 \\cr \\noalign{\\bigskip} \\noalign{\\hrule} } \\smallskip  Note.- $t_{x}$ represents the total acoustic radius of the models, $t_{s}$ the acoustic depth at the location of the helium gradients and $r_{s}/R$ the corresponding fractional radii ; all the models presented in this table include helium diffusion, except the model labelled 1.6h which is homogeneous in its outer layers.    \\end{table}  We have computed the oscillation frequencies and their second differences in five models : 1.6 M$_{\\odot}$ with ages 95 Myrs and 591 Myrs, and 1.6 Gyrs and 2.0 M$_{\\odot}$ with ages 63 Myrs and 649 Myrs (Table 1). The results are given in Figs. 1 to 5. Only modes of degrees $l=0,1,2$ and $3$ have been taken into account so that they could be treated together for the study of the second differences. The oscillation frequencies have been limited to the range 0.5 to 2 mHz to take into account the asymptotic approximation validity and the stellar cut-off. In this range 80 to 100 second differences are computed, according to the model.  In each figure six graphs are displayed : a) the helium profile inside the star ; b) the sound velocity, which shows a clear kink at the place of the helium gradient ; c) the first derivative of the sound velocity in which the kink is still clearer ; d) the second differences which show periodical oscillations ; e) the Fourier transform of these oscillations, in which clear peaks are found for precise time values ; f) the time needed for the acoustic waves to travel between the surface and the considered radius, or ``acoustic depth\".  We obtain clear signatures in the second differences of the frequencies of the diffusion-induced helium gradients inside stars. It can be checked that the periods corresponding to the peaks in the Fourier transforms (graphs e) are exactly twice the acoustic depths of the helium gradients (graphs f), as expected. In these figures, the scales are all the same except for graphs d) and f). In particular the Fourier transforms (graphs e) are presented at the same scale in each figure to show how the peak amplitudes decrease for deeper layers (larger times). As shown by Mazumdar {\\&} Antia (\\cite {mazumdar01}), the amplitude of the oscillatory signal in the second differences contains an amplification factor of $4 $sin$^2(\\pi t_s/t_*)$ where $t_*$ is the total acoustic radius of the star and $t_s$ the acoustic depth of the partial reflection region. Here this factor increases from about 2 to about 4 as the models evolve. The amplitude decrease by a factor $\\simeq10$ obtained in the present computations must be due to a competing process. It may be due partly to the fact that the energy in $p$ waves decreases towards the center and partly to the fact that the helium gradient becomes smoother with age, so that the reflected part (oscillations) is smaller for deeper discontinuities.  With these figures we can clearly follow the signatures of the helium gradients as they sink into the stars. The corresponding ages are obtained with the assumption of pure diffusion. In real stars macroscopic effects may change the time scales for helium settling. Comparisons with observed frequencies will give the positions of helium gradients inside the star, but we must keep in mind that the time needed for these gradients to develop depends on the physics.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{f1.eps}} \\vspace*{1cm} \\caption{Helium diffusion and its consequences for the stellar structure and oscillation frequencies for a 1.6 M$_{\\odot}$ at 95 Myrs ; a) helium profile as a function of the fractional radius ; b) sound velocity : a kink is clearly visible at the place of the helium gradient ; c) first derivative of the sound velocity : the kink is still more visible ; d) the second differences of the oscillation frequencies plotted as a function of the frequencies; e) the Fourier transform of graph (d) plotted as a function of time (in thousands of seconds) ; the vertical dashed line corresponds to twice the total acoustic depth of the star (see Table 1) ; f) the ``acoustic depth\", or time needed for the acoustic waves to travel from the surface to the considered radius : it can be checked that the peak of graph (e) corresponds to a time scale which is twice the ``acoustic depth\" of the helium gradient.} \\label{fig1} \\end{figure}   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{f2.eps}} \\vspace*{1cm} \\caption{Same as Fig. 1, for a 1.6 M$_{\\odot}$ model  at 591 Myrs.} \\label{fig2} \\end{figure}   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{f3.eps}} \\vspace*{1cm} \\caption{Same as Fig. 1, for a 1.6 M$_{\\odot}$ model at 1.6 Gyrs.} \\label{fig3} \\end{figure}  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{f4.eps}} \\vspace*{1cm} \\caption{Same as Fig. 1, for a 2.0 M$_{\\odot}$ model at 63 Myrs.} \\label{fig4} \\end{figure}  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{f5.eps}} \\vspace*{1cm} \\caption{Same as Fig. 1, for a 2.0 M$_{\\odot}$ model at 649 Myrs.} \\label{fig5} \\end{figure}   \\subsection{Models without element diffusion}  Similar computations have been done for evolution models in which helium diffusion is suppressed. In this case no helium gradient exists in the stellar outer layers and the oscillatory signal described in the previous section is not present. However, helium ionization induces a typical variation of the $\\Gamma_1$ coefficient which gives rise to another type of oscillatory signal in the frequencies. Computational results are shown in Fig. 6 for the case of the 1.6 M$_{\\odot}$ model  at 1.6 Gyrs. In this case graph a) displays the $\\Gamma_1$ coefficient in the stellar outer layers. As already studied by several authors (e.g. Miglio et al. \\cite {miglio03}), this coefficient is strongly influenced by the helium ionisation gradients. We can see the corresponding kinks in the sound velocity and its first derivative. This oscillatory signal does not appear when diffusion is effective as helium is then depleted in the region considered. Note that in the graphs of Fig. 6 the scales are quite different from those of Figs. 1 to 5. The effect due to the helium ionization gradients is much smaller than the effect due to the overall helium gradient in the case of diffusion. Here the amplification factor  $4 $sin$^2(\\pi t_s/t_*)$ is about 10 times smaller for the homogeneous model than for the model of the same age which includes helium diffusion. The large difference between the signal amplitudes in the two cases may be due to this effect. The right-hand graphs display the second differences, their Fourier transform and the travel time of the waves, as in Figs. 1 to 5. The features are very different from those obtained in case of diffusion, where the signature of the helium gradient is dominant.   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{f6.eps}} \\vspace*{1cm} \\caption{This figure is similar to Fig. 1 when element diffusion is suppressed, in a 1.6 M$_{\\odot}$ at 95 Myrs ; here graph a) represents the $\\Gamma_1$ coefficient which shows a clear feature at the place of the HeII ionisation region ; the other graphs are similar to those of Fig. 1 ; we can clearly see the influence of the helium ionisation regions on the frequencies. Note that the effect is much smaller than that of a helium gradient when diffusion is taken into account.} \\label{fig6} \\end{figure}"
    },
    "0404/astro-ph0404050_arXiv.txt": {
        "abstract": "{XMM--$Newton$ and $Chandra$/HETG spectra of the Compton--thin ($N_{\\rm H}\\sim4\\times10^{23}$ cm$^{-2}$) Seyfert 2 galaxy, NGC~4507, are analyzed and discussed. The main results are: a) the soft X--ray emission is rich in emission lines; an (at least) two--zone photoionization region is required to explain the large range of ionization states. b) The 6.4 keV iron line is likely emitted from Compton--thick matter, implying the presence of two circumnuclear cold regions, one Compton--thick (the emitter), one Compton--thin (the cold absorber). c) Evidence of an Fe {\\sc xxv} absorption line is found in the $Chandra$/HETG spectrum. The column density of the ionized absorber is estimated to be a few$\\times10^{22}$ cm$^{-2}$. ",
        "introduction": "NGC~4507 is a nearby ($z$=0.0118) spiral galaxy and one of the X--ray brightest Compton--thin Seyfert 2s, despite the heavy obscuration (N$_{\\rm H}\\sim4\\times10^{23}$ cm$^{-2}$). {\\sl ASCA} clearly detected, in addition to the absorbed power and the iron K$\\alpha$ line already observed by {\\sl GINGA} (Awaki et al. 1991), a strong soft X--ray excess and an intense emission line at $\\sim$0.9 keV (Comastri et al. 1997; Turner et al. 1997), identified as the Ne {\\sc ix} recombination line. $Beppo$SAX observed the source three times (Risaliti 2002); in the third observation the flux was about half that of the previous two observations. A Compton reflection component was also clearly detected; interestingly, its relative normalization about doubled in the third observation. This suggests a constant reflection flux in spite of the primary flux variations, and therefore an origin in distant matter.  In this paper we present results from XMM--$Newton$ and $Chandra$/HETG observations of NGC~4507. In Secs.~2 and 3 data reduction and analysis are described, while the results are summarized and discussed in Sec.~4.  In the following, we will assume $H_0$=70 km/s/Mpc. ",
        "conclusions": "\\subsection{Soft X--ray continuum and line emission.}  The soft X--ray emission is well fitted  by a steep ($\\Gamma\\sim$3) power law plus several emission lines from oxygen, neon, magnesium and silicon. The ionization state implied by the observed lines is too wide for a single photoionized region, or for collisionally ionized matter. At least two photoionized reflecting regions (or perhaps a region with a range of ionization parameters, as in NGC~1068, see Kinkhabwala et al. 2002) is therefore required. If the soft X--ray emission is indeed due to reflection of the nuclear radiation by ionized matter, this implies a significant steepening of the intrinsic emission at low energies, not unusual in Seyfert galaxies (see e.g. Vaughan et al., 2004, and references therein).  The He--like to H--like K$\\alpha$ lines ratio for both oxygen and neon is very large, i.e. at least a factor of 10 (see Table~\\ref{fit_lines}). The He--like line is actually a triplet, consisting of resonance, intercombination and forbidden lines, while the H--like line is a doublet of resonance lines. The resonance lines can easily be optically thick, with consequent reduction of their equivalent width due to resonant trapping (Matt et al. 1996). This is not true for the intercombination and, especially, forbidden lines which have much lower oscillator strengths: for low densities the forbidden line may indeed dominate the triplet, as shown by Porquet \\& Dubau (2000).   The line energies are consistent with rest frame values: no strong outflow is required by the data. The geometry of the reflecting regions is hard to derive with the present data. These regions must lie outside the cold, Compton--thin absorber in order to be observable. The lack of any extended emission seen by $Chandra$ implies a not very constraining upper limit to the size of the reflecting regions of about 800 pc.  \\subsection{The cold reprocessor}  The iron 6.4  keV emission line is clearly detected, along with its Compton shoulder. The shoulder--to--core ratio clearly points toward Compton--thick matter (Matt 2002), as did the earlier $Beppo$SAX detection of a Compton reflection continuum (Risaliti 2002). The constancy of the reflection component flux despite significant primary flux variations and the narrowness of the iron line as measured by $Chandra$ suggest that the Compton--thick matter is rather distant from the black hole, and likely to be identified with the molecular torus. This in turn implies the presence of two different cold circumnuclear regions, one Compton--thick (the reflector) the other Compton--thin (the absorber), as it is now often found in Seyfert 2s (Matt et al. 2003, and references therein). The Compton--thin region should have a rather low covering factor (as seen by the nucleus) so as not to overproduce the iron line.  \\subsection{The ionized Fe {\\sc xxv} absorption line}  In the $Chandra$/HETG spectrum, evidence for a  Fe {\\sc xxv} absorption line was found. The line is significant at about a 99\\% confidence level. The ionization state of the absorbing matter is much higher than that of the photoionized regions responsible for the soft X--ray lines.  Ionized iron absorption lines have been already observed in several sources. In some cases the line energies are significantly blueshifted, implying large velocity and very massive outflows (e.g. Pounds et al. 2003; Reeves et al. 2003, and references therein). In other cases (Reeves et al. 2004; Vaughan \\& Fabian 2004), iron absorption occurs in matter with no detectable motion, as in our case. It is very likely that the two absorbing regions are completely different, the first probably being much closer to the black hole and possibly due to a disk wind (e.g. King \\& Pounds 2003; Proga et al. 2000), the second likely related to the `hot' scattering regions often observed (via reflected continuum and ionized iron emission lines) in Seyfert 2 galaxies (the best studied case is NGC~1068, Matt et al. 2004 and references therein). Recently, such lines have also been found in unobscured Seyferts, thanks to the improved sensitivity of XMM--$Newton$ (Matt et al. 2001; Bianchi et al. 2003a).  The matter responsible for these lines is possibly located much further away than the outflowing matter, and indeed in one case, NGC~5506 (where highly ionized iron emission lines were observed), it is likely extended over several hundreds of parsecs (Bianchi et al. 2003b); however, in the case of NGC~3783 the variability of the line depth argues in favour of a much smaller (less than 0.1 pc) size (Reeves et al. 2004). The line EW in NGC~4507 corresponds (assuming solar iron abundance, an ionization state peaked on He--like iron and moderate turbulence to ensure the matter is optically thin in the line; see Nicastro, Fiore \\& Matt, 1999, for the relevant formulae and graphs) to an absorbing hydrogen equivalent column density of a few $\\times10^{22}$ cm$^{-2}$, i.e. of the same order as the one estimated by Bianchi \\& Matt (2002) for NGC~5506."
    },
    "0404/astro-ph0404099_arXiv.txt": {
        "abstract": "Instabilities driven by thermal and lepton diffusion (doubly diffusive instabilities) in a Ledoux stable fluid will, if present below the neutrinosphere of the collapsed core of a supernova progenitor (proto-supernova), induce convective-like fluid motions there. These fluid motions may enhance the neutrino emission by advecting neutrinos outward toward the neutrinosphere, and may thus play an important role in the supernova mechanism. ``Neutron fingers,'' in particular, have been suggested as being critical for producing explosions in the sophisticated spherically symmetric supernova simulations by the Livermore group \\citep[e.g.,]{wilsonm93}. These have been argued to arise in an extensive region below the neutrinosphere of a proto-supernova where entropy and lepton gradients are stabilizing and destabilizing, respectively, if, as they assert, the rate of neutrino-mediated thermal equilibration greatly exceeds that of neutrino-mediated lepton equilibration. Application of the Livermore group's criteria to models derived from core collapse simulations using both their equation of state and the Lattimer-Swesty equation of state do indeed show a large region below the neutrinosphere unstable to neutron fingers.  Because of the potential importance of fluid instabilities for the supernova mechanism, and the desire to understand the origin of convective-like fluid motions that may arise in upcoming multi-dimensional radiation-hydrodynamical simulations of core collapse, we develop a methodology introduced by \\citet{bruennd96} for analyzing the stability of a fluid in the presence of neutrinos of all flavors and in the  presence of a gravitational field. Neutrino-mediated thermal and lepton equilibration between a fluid element and its surroundings (background) is modeled as a linear system characterized by four response functions (i.e., thermal and lepton equilibration driven by entropy and lepton fraction differences between a fluid element and the background), the latter evaluated for a given thermodynamic state and fluid element radius by detailed neutrino transport simulations. These transport simulations employ both traditional and improved neutrino physics.  When applied to an extensive two-dimensional grid of core radii and fluid element sizes for each of several time slices of a number of proto-supernovae, we find no evidence for the neutron finger instability as described by the Livermore group. We find, instead, that the rate of lepton equilibration always exceeds that of thermal equilibration. Furthermore, we find that neither of the ``cross'' response functions, that is, entropy equilibration driven by a lepton fraction difference, and lepton equilibration driven by an entropy difference, is zero and that the first of these tends to be the largest of the four response functions in magnitude. These cross response functions play a critical role in the dynamics of the equilibration of a fluid element with the background. An important consequence of this is the presence of a doubly diffusive instability, which we refer to as ``lepto-entropy fingers,'' in an extensive region below the neutrinosphere where the lepton number, $Y_{\\ell}$, is small. This instability is driven by a mechanism very different from that giving rise to neutron fingers, and may play an important role in enhancing the neutrino emission. Deep in the core where the entropy is low and the lepton number higher, our analysis indicates a region unstable to another instability, also involving the cross response functions, which we refer to as ``lepto-entropy semiconvection.'' These instabilities, particularly lepto-entropy fingers, may have already been seen in some multi-dimensional core collapse simulations described in the literature. ",
        "introduction": "\\label{sec:DDInstabilities}  It has long been recognized that fluid instabilities may play an important role in the supernova mechanism. It is generally agreed that entropy driven convection between the neutrinosphere and the shock during the postbounce phase of the collapsed core of a supernova progenitor (hereafter \"proto-supernova\") will contribute turbulent pressure and increase the neutrino energy-deposition efficiency \\citep{bethe90, millerwm93, herantbhfc94, burrowshf95, jankam96, mezzacappacbbgsu98b, fryer99, fryerw02, burasrjk02}, but this may not be enough to produce explosions.  More controversial is the nature and the role of fluid instabilities below the neutrinosphere. Here matter and neutrinos are tightly coupled, and convective-like motions induced by instabilities can advect neutrinos. If these convective-like motions occur between the deeper interior and the neutrinosphere, \\nue's could be advected from the lepton rich deeper interior to the neutrinosphere, thus adding to the rate of \\nue\\ diffusion and enhancing the \\nue\\ luminosity. The analysis of fluid instabilities below the neutrinosphere must take into account the possibilities of thermal and lepton transfer by neutrinos. For example, a displaced fluid element in the presence of an entropy and/or lepton fraction gradient will find that its entropy and/or lepton fraction will differ from that of its surroundings or background. This will induce neutrino mediated thermal and lepton transfer which can modify the hydrodynamic buoyancy forces felt by the fluid element. This in turn can modify the subsequent growth rate of an instability, or the very existence of an instability. In particular, instabilities can arise in a gravitating fluid because of thermal and lepton transport that would otherwise be stable in their absence. These instabilities, involving the transport of two quantities (e.g., energy and leptons), are referred to as doubly diffusive instabilities.  To introduce the phenomenon, let us consider several classic examples of doubly diffusive instabilities. A doubly diffusive instability will arise if there are two differing ``substances,'' say heat and leptons, with gradients that can act oppositely in their stabilizing effects. For the sake of a concrete model, consider the case of heat and salt in water. Let us first imagine the case in which a region of hot salty water overlies a region of cold fresh water \\citep{stern60}. Normally the diffusion of heat occurs more rapidly than the diffusion of salt, and we will assume this to be the case here. Let us also assume that the magnitudes of the gradients in heat and salt are such that the fluid is stably stratified gravitationally (i.e., Ledoux stable). Imagine a parcel of hot salty water to be depressed slightly, as shown in Figure 1. The rapid diffusion of heat will thermally equilibrate this parcel of water with the background, but the slower diffusion of salt will result in the water remaining salty. We thus end up with a pocket of cold salty water in a background cold fresh water, and this pocket, being denser than the background, will continue to sink. The result is that ``fingers'' of salty water will poenetrate the fresh water on a thermal diffusion time scale. Note that what creates the instability in this otherwise stable situation is the presence of diffusion, and thus its description as a ``diffusive instability.'' The above described instability is referred to as salt fingers, and is an example of the kind of doubly diffusive instability that occurs if a slowly diffusing substance with a destabilizing gradient is stabilized in the absence of diffusion by the gradient of a rapidly diffusing substance.  \\begin{figure}[h!] \\setlength{\\unitlength}{1cm} {\\includegraphics[width=\\columnwidth]{Neutron_Fingers.eps}} \\caption{\\label{fig:Neutron_Fingers} Simple model of the salt-finger instability.} \\end{figure}  \\begin{figure}[h!] \\setlength{\\unitlength}{1cm} {\\includegraphics[width=\\columnwidth]{Semi_Convection.eps}} \\caption{\\label{fig:Semi_Convection} Simple model of the semiconvective instability.} \\end{figure}   Now imagine that the spatial configuration described above is reversed, namely, that a region of cold fresh water overlies a region of hot salty water, so that the gradient of salt is now stabilizing, the gradient of heat dstabilizing, and the magnitudes of the gradients such that the fluid is Ledoux stable \\citep{veronis65}. Imagine further that a parcel of cold fresh water is depressed slightly into the hot salty water, as shown in Figure 2. Again the rapid diffusion of heat will thermally equilibrate the parcel with the background, but the slower diffusion of salt will result in the parcel remaining fresh. We thus end up with a parcel of hot fresh water in a background of hot salty water,  and the parcel, being less dense than originally from the background, will accelerate upwards faster than would have been the case with no diffusion. This can result in an overstable situation, driven by the phase lag between the temperature and velocities of the fluid element which causes net work to be done over each cycle \\citep{spiegel72}. In this case the parcel will oscillate with growing amplitude. This instability is referred to as ``semiconvection,'' and is an example of the kind of doubly diffusive instability that occurs if a rapidly diffusing substance with a destabilizing gradient is stabilized in the absence of diffusion by a slowly diffusing substance with a stabilizing gradient.  Doubly diffusive instabilities such as described above can occur in the collapsed core of a supernova progenitor \\citep{smarrwbb81}. Neutron rich (low \\ye, where \\ye\\ is the proton fraction) matter tends to be denser than neutron poor (high \\ye) matter at the same temperature and pressure because the latter contains more light particles (electrons and electron neutrinos) that contribute to the pressure. Thus neutron richness in stellar cores can be thought of as playing a role analogous to that of salt in the discussion above. The ramp-up of the bounce shock on core rebound results in an outwardly positive (stabilizing) entropy gradient throughout much of the inner core. On the other hand, a negative lepton gradient formed at and initially confined to the vicinity of the \\nuesph\\ immediately after shock breakout is advected deeper into the core with time because of the inward matter flow and outward diffusion of \\nue's. A destabilizing negative \\ye\\ gradient stabilized by the outwardly positive entropy gradient can set the stage for a doubly diffusive instability analogous to salt fingers if lepton transport is slower than thermal transport. This instability is referred to as ``neutron fingers.''  The Livermore group has reported successful explosions in their spherically symmetric supernova simulations employing sophisticated multi-energy neutrino transport, and have emphasized the importance of the neutron-finger instability in powering these explosions \\citep{wilsonmww86, wilsonm88, wilsonm93}. The issue of whether or not this instability is operative, and if so whether or not it will power an explosion is important, as no other supernova simulations employing neutrino transport of comparable sophistication but without invoking neutron fingers has produced an explosion \\citep{bruenn93, mezzacappalmhtb01, liebendorfermtmhb01, burasrjk02}.  The Livermore group argue that thermal equilibration between a fluid element and the background is driven by the total neutrino number density (i.e., \\nue's, \\num's, and \\nut's and their antiparticles), whereas equilibration in composition (i.e., lepton fraction, \\yl) is driven only by the difference in the number densities of \\nue's and \\nuebar's. The ratio, $R$, of the two equilibration rates is thus given approximately by  \\begin{equation} R = \\frac{ | n_{\\nue} - n_{\\nuebar}  | }{  n_{\\nue} + n_{\\nuebar} + n_{\\num} + n_{\\numbar} + n_{\\nut} + n_{\\nutbar} } \\label{eq:DDI1} \\end{equation}  \\noindent  where $n_{i}$ is the number density of neutrinos of type $i$. From their simulations they typically find that $R < 0.1$. Based on the magnitude of this inequality the Livermore group tests for the neutron-finger instability by displacing a fluid element at constant composition but equilibrated with the background material in pressure and temperature. If the density of the fluid element relative to its surroundings after displacement is such as to drive the displacement further, the fluid is taken to be neutron-finger unstable. Mathematically, their criterion for neutron-finger instability is given by  \\begin{equation} \\left( \\pderiv{\\rho}{Y_{\\rm e} } \\right)_{T,P} \\frac{d Y_{\\rm e} }{dr} > 0 \\label{eq:DDI2} \\end{equation}  \\noindent where $\\frac{d Y_{\\rm e} }{dr}$ is the $Y_{\\rm e}$  gradient of the background. If a region was established to be neutron-finger unstable, a mixing length algorithm was then employed to simulate the convective fluid motions and mixing. The Livermore group found that the result of the convective-like motions due to the neutron-finger instability in their supernova simulations was the production of a larger \\nue\\ luminosity, as neutrinos and electrons previously trapped in the core were brought to the vicinity of the neutrinosphere. This increase in \\nue\\ luminosity was found to be sufficient to power an explosion.  In several papers, \\citet{bruennmd95} and \\citet{bruennd96} questioned the assumption that the ratio, $R$, of the rates of composition to thermal equilibration is given by equation (\\ref{eq:DDI1}). Situations were envisioned in which a fluid element perturbed with respect to the background could drive counter flows of \\nue's and \\nuebar's so that the numerator of equation (\\ref{eq:DDI1}) would be additive in  $n_{\\nue}$ and $n_{\\nuebar}$ and the denominator subtractive. Detailed simulations, performed by  \\citet{bruennd96}, of the equilibration of fluid elements perturbed with respect to the background showed that in many cases counter flows of \\nue's and \\nuebar's did indeed occur. Furthermore, it was found that lepton transport proceeded quickly by the rapid transport of low energy \\nue's and \\nuebar's, as the opacity for these was relatively small. On the other hand, thermal equilibration, being more efficiently accomplished by high energy neutrinos, was slowed by the relatively large opacities encountered by these neutrinos. Finally, the \\num's, \\nut's, and their antiparticles were less effective at mediating thermal equilibration because of there weak thermal coupling with the matter. The result was that the ratio of  composition to thermal equilibration depended on how the fluid element was perturbed with respect to the background, and in many cases this ratio was greater than unity.  Given its potential role in the core collapse supernova mechanism, there is clearly a need to clarify the issue of the neutron-finger instability in a proto-supernova. Additionally there has been the advent during the past decade of multi-dimensional core collapse supernova modeling \\citep{millerwm93, herantbhfc94, burrowshf95, jankam96, mezzacappacbbgsu98a, mezzacappacbbgsu98b, fryer99, fryerw02, burasrjk02}, with more simulations of greater sophistication to come. The extant simulations show the presence of fluid instabilities both above and below the neutrinosphere. The fluid instabilities above the neutrinosphere are driven by neutrino heating immediately above the gain radius, which is the surface above the neutrinosphere where neutrino cooling and heating balance. This instability is analogous to that of a fluid in a pot heated from below. The fluid instabilities below the neutrinosphere are more complicated to characterize, as neutrino transport drives both thermal and composition changes and the fluid instabilities can be doubly diffusive.  \\citet{bruennd96} developed a procedure for characterizing the stability/instability of a fluid element as a function of the entropy and composition gradient of the background and found that for typical conditions inside the collapsed stellar core of a supernova progenitor the fluid was either unstable to semiconvection or stable. However, they sampled only a small region of the run of thermodynamic conditions inside a supernova progenitor. In this paper we will further develop the methodology of \\citet{bruennd96} and use it to examine the stability of proto-supernovae for the entire run of thermodynamic conditions from the center of the core to the neutrinosphere. Futhermore, for each set of thermodynamic conditions we will consider the stability as a function of the size of the fluid element. The latter is important in that it governs the rate of thermal and composition equilibration. In Section  \\ref{sec:SPA} we begin our discussion of fluid stability in the ambience of neutrino mediated thermal and lepton transport by presenting the equations of motion of a fluid element that we will use to analyze stability. These equations involve four response functions that describe how thermal and lepton diffusion is driven by an entropy or a lepton fraction difference between a fluid element and the background. In Section \\ref{sec:SA} we describe how the solutions of these equations of motion are used to define various types of instability. We also examine the stability of a fluid element as a function of the entropy and lepton gradients of the background for various representative sets of the ``response functions.'' These examples illustrate the different kinds of instabilities that can arise in a proto-supernova. In particular, we find that neutron fingers is unlikely to occur anywhere below the neutrinosphere. However, we discover and describe several new and potentially important modes of doubly diffusive instabilities which we refer to as ``lepto-entropy fingers'' and ``lepto-entropy semiconvection.'' Both of the instabilities involve the cross response functions in essential ways. In Section \\ref{sec:Response} we describe how the response functions are computed for a given thermodynamic state and fluid element size. In Section \\ref{sec:CCModels} we begin the application of our stability analysis to core collapse models by describing these models, and in Section \\ref{sec:Results} we present the results of an extensive set of calculations giving the stability or type of instability and its growth rate as a function both the fluid element size and its location in the core. In Section \\ref{sec:Comp} we compare our analyses and results with prior work, and in in Section \\ref{sec:Cncl} we present our conclusions. ",
        "conclusions": "\\label{sec:Cncl}  The fluid below the neutrinosphere of the collapsed core of a supernova progenitor can be driven unstable by thermal and lepton diffusion as well as by a gravitationally unstable stratification of the fluid. We have developed the methodology introduced by \\citet{bruennd96} for analyzing fluid instabilities in the presence of neutrino mediated thermal and lepton transfer. In this analysis four ``response functions'' are introduced to model the thermal and lepton equilibration of a fluid element perturbed in entropy and lepton fraction from that of the background. These response functions depend on the thermodynamic state of the background and the size of the fluid element, and are obtained by detailed neutrino transport simulations. Given values of the response functions and the ambient gravitational field at the position of the fluid element, a cubic equation can be derived whose roots characterize the response of the fluid element to a perturbation in either its motion or thermodynamic state. From a determination of these roots, the stability or mode of instability of the fluid can be ascertained as a function of the gradients in entropy and lepton fraction. This methodology is applied both to resolve a long standing controversy regarding the presence of the neutron finger instability below the neutrinosphere, and to perform an extensive analysis of the fluid instabilities that do occur below the neutrinosphere for a number of representative post bounce supernova progenitors.  The Livermore group \\citep{wilsonmww86, wilsonm88, wilsonm93} have argued that thermal equilibration proceeds much more rapidly than lepton equilibration in the fluid below the neutrinosphere. As a consequence,  given the gradients in entropy and lepton fraction that arise in their supernova models subsequent to the initial collapse of the core and bounce they found that the neutron finger instability will arise in an extensive region below the neutrinosphere. Modeling by a mixing length algorithm the fluid motions that are expected to result from this instability, they found an enhancement of the \\nue\\ luminosity that helped to produce successful explosions in their sophisticated spherically symmetric supernova simulations. Their work is very important as it has emphasized the potential role of fluid instabilities below the neutrinosphere in the supernova mechanism.  We have performed core collapse simulations using the Livermore \\citep{wilson03} equation of state. Applying their criteria for the neutron finger instability, we do find an extensive region below the neutrinosphere, on either side of a Ledoux unstable region, that is neutron finger unstable. A similar exercise for core collapse simulations computed with the \\citet{lattimers91} equation of state shows a similar region of neutron finger instability, although not quite as extensive. Applying our methodology to the region below the neutrinosphere of these models we always find that lepton equilibration proceeds much more rapidly than thermal equilibration, and therefore that the neutron finger instability never occurs.  On the other hand, we have also found that the cross response functions, viz., the tendency for a perturbation in entropy of the fluid element relative to the background to drive a net lepton flow between them, and the tendency of a lepton fraction perturbation to drive a net energy flow, have substantial positive magnitudes for all thermodynamic states, and that the ratio of the second to the first of these response functions is large. These cross response functions cause a difference in lepton fraction between the fluid element and the background to develop primarily in response to an entropy difference. This is in contradistinction to the case without diffusion (the Ledoux case) in which a lepton fraction difference develops from a displacement of the fluid element through a gradient in the background lepton fraction.  As a result of this, several new doubly diffusive instabilities appear which we have referred to as lepton-entropy fingers and lepton-entropy semiconvection. These instabilities are expected to present in significant regions below the neutrinospheres of post collapsed stellar cores, may play an important role in the supernova mechanism. We have analyzed these instabilities in detail.  For each of a number of post bounce core collapse supernova models we have performed an extensive survey of fluid  instabilities on a two-dimensional grid of fluid element size and radial location in the core. Our results show a common pattern of unstable regions in these core collapse models. Outside of a Ledoux unstable region is a region where the lepton fraction is small, due to prior shock dissociation and deleptonization, and the entropy gradient is positive, due to the powering up of the shock. The fluid is leptom-entropy finger unstable in this region, with maximum growth rates for fluid element scales approximately 1/20 the distance of the fluid element from the core center. This unstable region can extend up almost to the neutrinosphere, and may therefore be important for the same reason adduced by the Livermore group for their neutron fingers. Below the Ledoux unstable region is region of unshocked material of relatively high lepton fraction. This region is unstable to lepto-entropy semiconvection with growth rates favoring small scales. Reviewing the multidimensional core collapse simulations that have been reported in the literature we conclude that the lepto-entropy finger instability may have already been seen."
    },
    "0404/astro-ph0404500_arXiv.txt": {
        "abstract": "The propagation of cellularly stabilized thermonuclear flames is investigated by means of numerical simulations. In Type Ia supernova explosions the corresponding burning regime establishes at scales below the Gibson length. The cellular flame stabilization---which is a result of an interplay between the Landau-Darrieus instability and a nonlinear stabilization mechanism---is studied for the case of propagation into quiescent fuel as well as interaction with vortical fuel flows. Our simulations indicate that in thermonuclear supernova explosions stable cellular flames develop around the Gibson scale and that a deflagration-to-detonation transition is unlikely to be triggered from flame evolution effects here. ",
        "introduction": "\\label{sec:1}  The standard model of Type Ia supernovae (SNe Ia) describes these astrophysical events as thermonuclear explosions of white dwarf (WD) stars \\cite{hoyle1960a}. In our study, we refer to the specific scenario (for a review on SN Ia explosion models see \\cite{hillebrandt2000a}), where the WD consists of carbon and oxygen and the thermonuclear reaction propagates in form of a flame that starts out in the so-called \\emph{deflagration mode}. That is, the flame is mediated by microphysical transport processes and its burning velocity is subsonic. One key ingredient in such a SN Ia model is the determination of the effective propagation velocity of the deflagration flame. The so-called laminar burning speed $s_l$, i.e.~the propagation velocity of a \\emph{planar} deflagration flame, is much too low (a few percent of the sound speed in the unburnt material) to explain powerful SN Ia explosions. The solution to this problem is provided by the concept of turbulent combustion. Instabilities on large scales result in the formation of a \\emph{turbulent cascade}, where large-scale eddies decay into smaller ones thereby transporting kinetic energy from large to small scales. Interaction of the flame with those eddies wrinkles the flame front and enlarges its surface. This is equivalent to an increase in the net fuel consumption rate and hence causes an acceleration of the flame.  Recent SN Ia explosion models on scales of the WD could show that in this way enough energy can be released to gravitationally unbind the star \\cite{reinecke2002d,gamezo2003a}. However, these models cannot resolve all relevant length scales down to the flame width and therefore have to rely on assumptions on the physics on unresolved small scales. The goal of our small-scale simulations of the propagation of the thermonuclear flame is to test those assumptions and eventually to reveal new physics that additionally needs to be included in the SN Ia models. The most significant assumptions of the large scale models by Reinecke et al.~ (e.g.~\\cite{reinecke2002d}) are that the energy input into the turbulent cascade originate solely from large scales and the flame be stable on unresolved scales.  A fundamental feature of the turbulent cascade is that in a certain intermediate scale range, the turbulent velocity fluctuations decrease monotonically with smaller length scales of the eddies. This is captured by the corresponding scaling law. From that effect it is obvious, that there must exist a cutoff scale, below which flame propagation is not affected by the turbulent cascade anymore. Below this \\emph{Gibson scale}, the velocity fluctuations are so small, that the flame burns faster through the eddies than they can deform it. Thus the flame here propagates through ``frozen turbulence''. Our simulations aim on the study of effects around and below the Gibson scale. Except for the very late stages of the SN Ia explosion, the Gibson scale is well-separated from the width of the flame and the flame may safely be regarded as a discontinuity between burnt and unburnt material. ",
        "conclusions": "The presented simulations of flame propagation into quiescent fuel are in good agreement with theoretical expectations. The linear stage of flame evolution is consistent with Landau's dispersion relation \\cite{roepke2003a, landau1944a} and in the nonlinear regime the flame stabilizes in a cellular pattern. Thus, our hydrodynamical model of flame evolution shows that the cellular burning regime exists for SN Ia explosions.  Interaction of the flame with vortical flow fields may lead to a break-down of cellular stabilization if the velocity fluctuations are sufficiently large. In this case, however, the flame shows the tendency to adapt to the imprinted flow field. Thus, we observe a moderate increase in the effective flame propagation velocity but no drastic effects. No indication of active turbulent combustion could be found. Therefore, it seems unlikely that effects around the Gibson scale account for a DDT at late stages of the SN Ia explosion as has been anticipated by \\cite{niemeyer1997b}. More detailed discussions of flame propagation into quiescent fuel at different fuel densities and interaction with vortical flow fields of varying strengths are currently in preparation \\cite{roepke2003b,roepke2003c}."
    },
    "0404/astro-ph0404446_arXiv.txt": {
        "abstract": "The interplay between stellar populations and gas in local starburst galaxies is analyzed using images from the Hubble Space Telescope to map the ages of the young stellar components and to isolate the contribution of shocks on spatial scales ranging from a few tens of pc to $\\sim$1~kpc. The shocked gas represents a small fraction of the total ionized gas in these objects, yet it can have profound effects on the long--term evolution of the starburst, which may include the triggering of new star formation. ",
        "introduction": "The evolution of galaxies and of the intergalactic/intracluster medium (IGM/ICM) are closely connected, via the continuous interchange between the ISM and the IGM/ICM. The heating and metal enrichment of the IGM/ICM, and the enrichment of galaxies' ISM with pristine gas via infall from the IGM are manifestations of this close connection. Evidence for gas outflows from galaxies has been found at redshift as high as $\\sim$3 (\\cite{pettini00}), and may have contributed to the early pollution of the IGM (\\cite{ellison00}).  The driving mechanism for the energy and metal pollution of the IGM/ICM is the feedback from star formation. In the basic scenario, the combined effects of ejecta and energy release from supernovae and massive stars in regions of intense star formation form superbubbles that sweep up the surrounding medium and produce cavities containing hot, shocked gas (\\cite[Weaver et al. 1977]{weaver77}). If the bubble has enough energy to expand to a size of a few times the scale height of the ambient medium, the swept-up shell may break out and the wind internal to the bubble may expand into the IGM (\\cite[Chevalier \\& Clegg 1985]{chevalier85}).  Many theoretical and observational studies have investigated the mechanical feedback from star formation and its efficiency in transfering energy and material to regions far removed from the site of star formation (\\cite[e.g., De Young \\& Heckman 1994]{deyoung94}; \\cite[Suchkov et al. 1994]{suchkov94}; \\cite[Mac Low \\& Ferrara 1999]{maclow99}; \\cite[Strickland \\& Stevens 2000]{strickland00}; \\cite[Heckman, Armus \\& Miley 1990]{heckman90}; \\cite[Dahlem 1997]{dahlem97}; \\cite[Ferrara \\& Tolstoy 2000]{ferrara00}; \\cite[Cecil, Bland-Hawthorn \\& Veilleux 2002]{cecil02}), as well as in triggering renewed star formation (\\cite[e.g., Wada \\& Norman 2001]{wada01}). Yet many questions remain to be fully answered. Among these, progressing from small to large scales: the impact of the spatial distribution, history, and duration of the star--forming event (\\cite[Clarke \\& Oey 2002]{clarke02}); the coupling of the feedback from star formation with the surrounding ISM; the dependence of the feedback efficiency on the conditions of the ISM (e.g. porosity; \\cite[Clarke \\& Oey 2002]{clarke02}); and the role of the global galaxy parameters (mass, metallicity, environment; \\cite[Dahlem 1997]{dahlem97}; \\cite[Strickland et al. 2004]{strickland04}).  A review of some of these issues is given by D. Strickland (2004, these proceedings).  One of the complicating factors in tackling the various facets of feedback is the vast range of physical scales that needs to be probed. While the relevant scales for charting star formation are those of the stars and star clusters (a few pc), a study of the structure and conditions of the ISM needs to sample scales of tens to hundreds of pc. Moving on to larger scales, superwinds are best investigated on kpc scales, while probing the influence of global parameters involves scales of tens of kpc (galactic sizes) to Mpc (interactions).  Even for the closest galaxies, the small--to--intermediate scales relevant for the interaction feedback--ISM are accessible only via high angular resolution observations (e.g., with the Hubble Space Telescope), as 10~pc=0.4$^{\\prime\\prime}$ for a galaxy at 5~Mpc distance. The study reviewed here involves four local starburst galaxies observed with the Wide Field Planetary Camera~2 on the HST. Because of their intense-to-extreme star formation rates, starburst galaxies are the sites where stellar feedback can have its most dramatic influence on the structure and evolution of the ISM (and surrounding IGM; \\cite{heckman90}), and are, therefore, optimal laboratories to investigate mechanical feedback. The four galaxies in the present study cover a range in luminosity, metallicity, star formation rate, and environment, but they are all closer than 5~Mpc (Table~\\ref{tab1}). This investigation attempts at addressing the issue of the coupling feedback--ISM, and its relation to the duration of the star forming event. The problem is tackled from two fronts: stellar populations and ISM conditions. The ages of the young ($\\lesssim$300~Myr) stellar clusters are used to set a minimum value to the duration of the current starburst event, while measurements of the shocked gas provide a constraint on the fraction of starburst mechanical output that is recovered in the ISM within a small distance ($\\lesssim$1~kpc) of the site of star formation.  \\begin{table} \\begin{center} \\begin{tabular}{lrrrrrrr} Galaxy  & $D$   & $d$ & $M_B$ &   $(O/H)$ & $SFR(H\\alpha)$ & $    L(H\\alpha)_{nph}/L(H\\alpha)$ &    $    L(H\\alpha)^{pred}_{nph}/L(H\\alpha)_{nph}$\\\\ Name    & (Mpc)   & (kpc) & &           & (M$_{\\odot}$~yr$^{-1}$) & & \\\\[3pt] NGC3077&3.85&1.4&$-17.5$&8.9&0.076&0.043&0.51--1.40\\\\ NGC4214&2.94&2.0&$-17.2$&8.2&0.089&0.041&0.29--1.24\\\\ NGC5236&4.5&0.75&$-20.3$&9.2&0.308&0.030&1.27--2.90\\\\ NGC5253&4.0&1.5&$-17.5$&8.2&0.270&0.032&0.38--1.19\\\\ \\end{tabular} \\caption{Properties of the observed galaxies. Columns are as follows; $D$: distance; $d$: physical size of the region observed; $M_B$: absolute B magnitude; $(O/H)$: oxygen abundance, 12$+$log(O/H); $SFR(H\\alpha)$: star formation rate derived from the extinction corrected H$\\alpha$ luminosity; $L(H\\alpha)_{nph}/L(H\\alpha)$: fraction of the H$\\alpha$ luminosity associated with shocks; $L(H\\alpha)^{pred}_{nph}/L(H\\alpha)_{nph}$: ratio of the predicted-to-observed H$\\alpha$ luminosity associated to shocks. NGC3077 is in the M81 group, in close interaction with M81 itself; NGC4214 is an isolated galaxy, while NGC5236 and NGC5253 form a loose pair.} \\label{tab1} \\end{center} \\end{table} ",
        "conclusions": "\\label{sec:concl}  The characteristics of the stellar populations and ISM in the starbursts of four nearby galaxies suggest a strong coupling between the starbursts and the surrounding ISM. This is particularly true for the three dwarf galaxies in the sample. Here the amount of observed H$\\alpha$ luminosity associated with shocks corresponds to about 70\\%--100\\% of the total mechanical output from the starbursts, and places a lower limit of about 30~Myr to the duration of the starbursts (i.e., these cannot be `bursts'). Such large fractions of mechanical energy available in the immediate surroundings of starbursts may suggest triggered star formation as a mechanism to explain the long duration timescales inferred from stellar population studies ($\\approx$100--300~Myr), in the absence of direct, recent triggers."
    },
    "0404/hep-ph0404128_arXiv.txt": {
        "abstract": "We construct a phenomenological model of electroweak-scale inflation that is in accordance with recent cosmic microwave background observations by WMAP, while setting the stage for a zero-temperature electroweak transition as assumed in recent models of baryogenesis. We find that the scalar spectral index especially poses tight constraints for low-scale inflation models. The inflaton--Higgs coupling leads to substantial mixing of the scalar degrees of freedom. Two types of scalar particles emerge with decay widths similar to that of the Standard Model Higgs particle. ",
        "introduction": "With the results from the WMAP mission \\cite{WMAPparam,WMAPweb} it has become relevant to critically review models of inflation, especially with regard to the scalar spectral index. While clearly still susceptible to improvement, the cosmic microwave background (CMB) observations are accurate enough to rule out certain (classes of) models, see e.g.\\ \\cite{WMAPinfl,Kinneyetal,LeachLiddle}. In this paper we consider electroweak-scale inflation, which turns out to be indeed tightly constrained by the spectral index.  The motivation for looking at electroweak-scale (i.e.\\ of order 100~GeV) inflation is twofold. Firstly, it is interesting to see if one can construct a working model of inflation with just minimal extensions of the Standard Model (SM) of particle physics, and to derive what kind of additional constraints such a coupling to the SM puts on an inflation model. The second (main) motivation has to do with baryogenesis, the production of the observed baryon asymmetry in the universe. As reviewed in \\cite{RubakovShaposhnikov}, all necessary ingredients for baryogenesis (baryon number violation, C and CP violation, and non-equilibrium) are present in the SM. This provides a strong motivation for trying to construct a working model of electroweak baryogenesis. However, the current lore is that in a standard finite-temperature electroweak transition both the CP-violating and the non-equilibrium effects are too small to be able to account for the observed baryon asymmetry. These problems may be resolved in the context of tachyonic preheating at the electroweak scale after low-scale inflation \\cite{KraussTrodden,GarciaBellidoetal,Felderetal, Copelandetal,GarciaBellido,TranbergSmit}. A tachyonic electroweak transition is strongly out of equilibrium, and the fact that the process takes place at zero temperature at the end of inflation may maximize the effectiveness of CP violation \\cite{TalkSmit}. In addition, the low reheating temperature prevents sphaleron wash-out of the baryon number produced \\cite{RubakovShaposhnikov}.  In this context it becomes important to check that the models that combine low-scale inflation with tachyonic electroweak preheating satisfy all the new observational constraints from WMAP. Low-scale inflation has been considered in many papers (see for instance \\cite{Knox,Kinney,Copelandetal,Germanetal,Lythrecent}). In this paper we build in particular on \\cite{Copelandetal}, which was also motivated by the problem of electroweak baryogenesis. The main idea is that we have a kind of hybrid inflation model \\cite{ShafiVilenkin,Lindehybr}, in which inflation is driven by a nearly constant potential energy of order (100~GeV)$^4$ \\cite{KraussTrodden,GarciaBellidoetal}, while one field (the inflaton) slowly rolls down its potential and the other field (the Higgs) is in a local minimum at zero. Once the inflaton passes a critical value, the local minimum for the Higgs field develops into a local maximum and both fields roll down rapidly to the absolute minimum at a non-zero value of the Higgs field, thus breaking the electroweak symmetry. As was shown in \\cite{Copelandetal}, ordinary hybrid inflation models in which the inflaton rolls from large field values towards zero are not viable at low scales because of large quantum loop corrections. This problem can be avoided in inverted hybrid inflation models, in which the inflaton rolls away from zero and inflation takes place at very small field values. Note that unlike standard hybrid inflation, slow-roll inflation ends in this case before the critical value is reached, instead of the end being caused by the phase transition, so that the slow-roll inflation stage and the phase transition can be considered as two separate processes.  The paper \\cite{Copelandetal} was written before WMAP, and the authors did not study the spectral index. As we will show in this paper, their model gives a spectral index that is too low according to WMAP. In \\cite{Germanetal} somewhat more general low-scale inflation models were considered, although not from the point of view of electroweak baryogenesis, but these models still appear to be incompatible with WMAP. We shall show that one can improve these models to obtain a spectral index that lies comfortably within the WMAP confidence levels.  At first sight it may seem that one can always fine-tune a model with sufficient parameters to satisfy the constraints, but this is not necessarily the case within a set of reasonable rules. To formulate these rules we start from the point of view that we are constructing a purely phenomenological model, a minimal extension of the SM that introduces only one extra inflaton field in order to describe the history of the universe during and after inflation. We stress that there is nothing wrong with fine-tuned parameters in a phenomenological model, as the phenomenologically very successful SM shows. In incorporating a slow-roll inflationary regime compatible with the CMB measurements, one is led to an inflaton potential with non-renormalizable couplings. To constrain this potential we assume a polynomial approximation, such that there is only a limited number of terms to be parametrized.  This leads to a tight fit when we also incorporate the scenario of tachyonic electroweak baryogenesis, which requires that the inflaton field ends up far from the slow-roll regime with a vacuum expectation value similar to that of the Higgs field. The inflaton is a gauge singlet and it couples only to the radial (gauge-invariant) mode of the Higgs field. This coupling should induce a sufficiently fast tachyonic electroweak transition to make baryogenesis possible, without being unrealistically large. It implies a considerable mixing between the inflaton and Higgs modes, and the model predicts the existence of (only) {\\em two} scalar particle species with electroweak-scale masses. Up to mixing-angle factors, their decay widths are similar to that of a SM Higgs with the same mass. The model should therefore be falsifiable by accelerator experiments, in particular with the LHC.  An important issue with any slow-roll inflation model is the question whether the assumed flatness of the effective potential is consistent with basic properties of quantum fields, with `quantum corrections'. We investigate this by calculating one-loop corrections to the effective potential. This exercise also led us to a rough estimate of the scale at which the model may be expected to break down because of its non-renormalizable and strong couplings.  The outline of the paper is as follows. In section~\\ref{Nk} we first address the number of e-folds of inflation between horizon crossing of a WMAP-observable scale and the end of inflation, which number is crucial for the computation of CMB observables. Contrary to the generic situation, there is little uncertainty here because the (p)reheating of the universe and the onset of the radiation-dominated era are reasonably well understood in this model. Next, in section~\\ref{model1}, we review the model of \\cite{Copelandetal} and show that its spectral index disagrees with WMAP. The implied inflaton--Higgs mixing is studied in section~\\ref{mixing}. We then show in the following section (plus appendix~A) that, by adding two additional terms to the potential and tuning the coupling parameters to a certain extent, values for the scalar spectral index in agreement with WMAP can be obtained. In section~\\ref{loop} (plus appendix~B) we calculate one-loop quantum corrections to the effective potential and study the implications. Finally, section~\\ref{concl} summarizes our conclusions. ",
        "conclusions": "\\label{concl}  In this paper we investigated the implications of the WMAP results for low-scale inflation and found that there are severe constraints.  This is essentially because the proximity of the observed scalar spectral index $\\tn\\equiv n-1$ to zero is hard to reconcile with the  small number of e-folds between horizon crossing of the observable scales and  the end of inflation in these models.  Working in the context of a phenomenological electroweak-scale inflation model that allows for tachyonic electroweak baryogenesis and consists of one additional scalar field $\\gs$ coupled to the Standard Model Higgs, we were led to further constraints on the inflaton--Higgs potential. However, we found that there is a range of parameters compatible with a spectral index close to (and even larger than) zero in this model.  The polynomial approximation together with all the constraints led to the conclusion that we need a $\\gs^5$ term in the potential during inflation (as in \\cite{Copelandetal}). In addition $\\gs^2$ and $\\gs^4$ terms are needed, and (or instead of the $\\gs^4$) there might be a $\\gs^3$ term, but no powers higher than 5 can be present during inflation. The appearance of an odd power ($\\gs^5$) implies a local minimum at small negative values of the inflaton field ($\\gs=0$ being the value where the potential has a local maximum). A universe with initial conditions in this negative region would classically never stop inflating, which is why we assumed a small positive initial condition for the inflaton field. However, quantum tunnelling might very well make the case of a negative initial condition viable as well. The use of a polynomial approximation is quite natural for the inflationary region of the potential where the inflaton field is small, but it seems somewhat artificial in the large field region where the Higgs field comes into play. The odd and non-renormalizable power $\\gs^5$ in the potential appears to be the price we have to pay for keeping the number of parameters limited. It is not excluded of course that the model with its non-symmetric potential and non-renormalizable couplings can be embedded satisfactorily in a model with more symmetry that is also renormalizable, e.g.\\ a supersymmetric extension of the Standard Model.  In the one-loop calculation we found that the very different renormalization conditions in the inflationary and electroweak regimes did not lead to unresolvable conflicts. In our results quantum corrections do not disrupt the required flatness of the potential in the inflationary region. The non-renormalizable couplings and also the relatively strong inflaton self-coupling suggest a breakdown of the model already occurring at a fairly low scale, perhaps below 1 TeV, depending on the choice of parameters.  We conclude with the important remark that the phenomenological model arrived at here can be falsified experimentally through its conspicuous generic feature: the existence of two (and only two) particle species with zero spin and masses around the electroweak scale, and couplings to the rest of the SM equal to that of the Higgs up to factors related to the mixing angle.   \\ack This work is supported by FOM/NWO and by PPARC Astronomy Rolling Grant PPA/G/O/2001/00476.  \\appendix"
    },
    "0404/astro-ph0404137_arXiv.txt": {
        "abstract": "The High Resolution Fly's Eye (HiRes) experiment is an air fluorescence detector which, operating in stereo mode, has a typical angular resolution of $0.6^{\\circ}$ and is sensitive to cosmic rays with energies above $10^{18}$\\,eV. HiRes is thus an excellent instrument for the study of the arrival directions of ultrahigh energy cosmic rays.  We present the results of a search for anisotropies in the distribution of arrival directions on small scales ($<5^{\\circ}$) and at the highest energies ($>10^{19}$~eV).  The search is based on data recorded between 1999 December and 2004 January, with a total of 271 events above $10^{19}$\\,eV. No small-scale anisotropy is found, and the strongest clustering found in the HiRes stereo data is consistent at the 52\\,\\% level with the null hypothesis of isotropically distributed arrival directions. ",
        "introduction": "Identifying the sources of ultrahigh energy cosmic rays remains one of the central challenges in astrophysics.  After three decades of systematic searches for the origin of these particles, source identification still remains elusive. Sky maps of cosmic ray arrival directions at all energies are generally isotropic, with no obvious source or source region standing out.  A direct way to search for sources of ultrahigh energy cosmic rays is to analyze the distribution of their arrival directions for small-scale clustering. Any significant clustering in arrival directions could be evidence of nearby, compact sources, whereas the lack of clustering is consistent with models in which ultrahigh energy cosmic ray sources are distributed at large distances from our Galaxy.  Arrival directions do not necessarily point back to sources, as charged cosmic ray primaries suffer deflections traveling through Galactic and intergalactic magnetic fields.  The strength and orientation of these fields is not well established, so the size and direction of the deflection is difficult to ascertain. However, since the Larmor radius increases with energy, the possibility of observing small-scale anisotropy associated with cosmic rays pointing back to their origins is expected to grow.  Indeed, small-scale clustering of cosmic ray arrival directions at the highest energies has been previously claimed.  The AGASA (Akeno Giant Air Shower Array) experiment reported possible clustering in their sample of events with energies above $4\\cdot10^{19}$\\,eV \\citep{agasa1996}. The analysis has been updated several times \\citep{agasa1999, agasa2001, agasa2003}, most recently reporting six clusters (five doublets and one triplet) in a sample of 59 events, where a cluster is defined as a set of events with angular separation less than $2.5^{\\circ}$. The chance probability of this signal was reported to be less than $10^{-4}$ \\citep{agasa2003}.  Given the potential importance of this result for our understanding of the origin of cosmic rays, it is crucial to test the claim that clustering is a feature of cosmic ray arrival directions with independent experimental data. Since 1999, the High Resolution Fly's Eye (HiRes) air fluorescence experiment has been operating in stereo mode, collecting data of unprecedented quality on the arrival direction, energy, and composition of ultrahigh energy cosmic rays. In this Letter, we report results of a search for small-scale anisotropy in the arrival directions of ultrahigh energy cosmic rays observed by the HiRes stereo detector between 1999 December and 2004 January. ",
        "conclusions": "We perform the scan on the HiRes stereo sample of 271 events above $10^{19}$\\,eV.  Because we start well below the $4\\cdot 10^{19}$\\,eV energy associated with the AGASA clustering signal, our search should safely encompass the energy region of interest even in the presence of a systematic energy shift of 30\\% between the two experiments, as suggested by \\citet{demarco2003}. Starting at this energy does not appreciably dilute the significance of a clustering signal if one is found at higher energy, since the scan involves repeated searching with successively higher energy thresholds. An additional motivation for starting at $10^{19}$\\,eV is the fact that the HiRes angular resolution ($0.6^{\\circ}$) is much sharper at this energy than AGASA's ($2.8^{\\circ}$) \\citep{agasa1999}.  The results of the scan are shown in Figure\\,\\ref{scan}. The strongest clustering signal ($P_{min}=1.9\\%$) is observed using the energy threshold $E_c = 1.69\\cdot 10^{19}$\\,eV where we observe $n_p=10$ pairs separated by less than $\\theta_c=2.2^{\\circ}$ within a set of $N_c=120$ events. The statistical significance of this result corresponds to $P_{ch}=52\\%$.  \\begin{figure} \\epsscale{.95} \\plotone{f3.eps} \\caption{Autocorrelation scan of the HiRes data set above $10^{19}$\\,eV.  $P(N,\\theta)$ is the probability of obtaining the same or greater number of pairs as is actually observed in the data using a maximum separation angle $\\theta$ and searching among the $N$ highest-energy events.  These probabilities do not include the statistical penalty due to scanning. \\label{scan}} \\end{figure}  The HiRes stereo data above $10^{19}$\\,eV is therefore consistent with the null hypothesis of isotropic arrival directions.  Comparison with the AGASA clustering result is not straightforward. The HiRes stereo event sample above $4\\cdot 10^{19}$\\,eV is still smaller than AGASA's, though how much smaller depends critically on the level of agreement in absolute energy scale for the two experiments. The possibility of a systematic energy shift of 30\\% would imply that above the rescaled energy threshold, $(0.7)\\cdot 4\\cdot 10^{19}\\,{\\rm eV} = 2.8 \\cdot 10^{19}\\,{\\rm eV}$, HiRes has seen 47, rather than 27, events. More importantly, there is the question of how many pairs an independent data set might be expected to contain, given the lack of an obvious source model and the widely varying estimates of the strength of the AGASA clustering.  Without assuming a model and source strength, there is no natural way to translate the AGASA observation of five doublets and one triplet separated by less than $2.5^{\\circ}$ into a meaningful prediction for HiRes.  However, what can be tested using a statistically independent data set is the claim that significant small-scale clustering is a general feature of ultrahigh energy cosmic ray arrival directions.  The HiRes stereo data set does not support such a claim. We observe no statistically significant evidence for clustering on any angular scale up to $5^{\\circ}$ at any energy threshold above $10^{19}$\\,eV.  Comparing the observed value of $P_{ch}$ with the values obtained from simulations in Section 4 (shown in Table\\,\\ref{simtable}), we note that if the current HiRes data above $4\\cdot 10^{19}$\\,eV contained two or more pairs of events contributed by compact sources at the angular resolution limit of the detector, then the typical value of $P_{ch}$ would be 0.018 or less, and more than 90\\% of the time the value of $P_{ch}$ would be much smaller than the observed value of 0.52.  Results of searches for correlations with known astrophysical source classes will be published in a separate paper."
    },
    "0404/astro-ph0404301_arXiv.txt": {
        "abstract": "{ We investigate the role of dust in star formation activity of extremely metal-poor blue compact dwarf galaxies (BCDs). Observations suggest that star formation in BCDs occurs in two different regimes: ``active'' and ``passive''. The ``active'' BCDs host super star clusters (SSCs), and are characterised by compact size, rich \\H2 content, large dust optical depth, and high dust temperature; the ``passive'' BCDs are more diffuse with cooler dust, and lack SSCs and large amounts of \\H2. By treating physical processes concerning formation of stars and dust, we are able to simultaneously reproduce all the above properties of both modes of star formation (active and passive). We find that the difference between the two regimes can be understood through the variation of the ``compactness'' of the star-forming region: an ``active'' mode emerges if the region is compact (with radius $\\la 50$ pc) and dense (with gas number density $\\ga 500$ cm$^{-3}$). The dust, supplied from Type II supernovae in a compact star-forming region, effectively reprocesses the heating photons into the infrared and induces a rapid \\H2 formation over a period of several Myr. This explains the high infrared luminosity, high dust temperature, and large \\H2 content of active BCDs. Moreover, the gas in ``active'' galaxies cools ($\\la 300$ K) on a few dynamical timescales, producing a ``run-away'' star formation episode  because of the favourable (cool) conditions. The mild extinction and relatively low molecular content of passive BCDs can also be explained by the same model if we assume a diffuse region (with radius $\\ga 100$ pc and gas number density $\\la 100$ cm$^{-3}$). We finally discuss primordial star formation in high-redshift galaxies in the context of the ``active'' and ``passive'' star formation scenario. ",
        "introduction": "\\label{sec:intro}  The surfaces of interstellar dust grains are known to be sites where efficient formation of hydrogen molecules (\\H2) takes place. Without dust grains, \\H2 forms only in the gas phase with a production rate much lower than the dust surface reaction (e.g., Peebles \\& Dicke \\cite{peebles68}; Matsuda et al.\\ \\cite{matsuda69}). The shielding of \\H2 dissociating photons by dust grains also enhances molecule formation. In general, dust also absorbs a part of the radiation from stars and reemits it in the infrared (IR)\\footnote{In this paper, IR indicates the wavelength range where the emission from dust dominates the radiative energy (roughly 8--1000 $\\mu$m).}, thereby modifying the spectral energy distribution (SED) of galaxies (e.g., Silva et al.\\ \\cite{silva98}). Therefore, dust plays an important role in both chemical and radiative properties of galaxies.  Recently Hirashita \\& Ferrara (\\cite{hirashita02}, hereafter HF02) have proposed that the dust enrichment in extremely metal-poor primeval objects is essential for the enhancement of star formation activity. Their scenario suggests the following cycle between dust production and star formation: {\\it (i)}~massive stars end their lives as Type II supernovae (SNe II), which supply dust grains into the interstellar medium (ISM) (Kozasa et al.\\ \\cite{kozasa89}; Todini \\& Ferrara \\cite{todini01}; Nozawa et al.\\ \\cite{nozawa03}; Schneider et al.\\ \\cite{raffa04}); {\\it (ii)}~those grains enhance the formation of molecular clouds in which stars form; {\\it (iii)}~some of those stars are massive and the dust supply described in {\\it (i)} occurs again. This cycle {\\it (i)}--{\\it (iii)} effectively enhances the star formation rate as shown by HF02.  In fact, a large amount of dust has been suggested to exist at high redshift (high $z$) (e.g., Smail et al.\\ \\cite{smail97}). However, it is not easy to explore the first dust enrichment in primeval galaxies at high $z$ ($z\\ga 5$) with the present observational facilities. Therefore, nearby templates for primeval galaxies are useful to test galaxy formation scenarios. The best candidates for such a template are metal-poor blue compact dwarf galaxies (BCDs), since they are at the initial stage of metal enrichment and their current star formation activity is generally young (Searle \\& Sargent \\cite{searle72}; Kunth \\& \\\"{O}stlin \\cite{kunth00}). In other words, BCDs can be used as laboratories in which to study high-$z$ primeval galaxies.  Two classes of star formation activity have recently emerged observationally, as proposed for a BCD sample by Hunt et al.\\ (\\cite{hunt-cozumel}, hereafter HHTIV; \\cite{hunt-active}). They argue that the star-formation activity in the two most metal-poor galaxies, \\sbs\\ and \\izw, shows very different properties, in spite of their similar metallicities (1/41 and 1/50 $Z_\\odot$, respectively; Skillman \\& Kennicutt \\cite{skillman93}; Izotov et al.\\ \\cite{izotovetal99}). The major star-forming region of \\sbs\\ is compact and dense (radius $r_{\\rm SF}\\la 40$ pc, number density $n\\ga 600$ cm$^{-3}$; Dale et al.\\ \\cite{dale01}; Izotov \\& Thuan \\cite{izotov99}). Moreover, \\sbs\\ hosts several super star clusters (SSCs), detectable \\H2 emission lines in the near-infrared (Vanzi et al.\\ \\cite{vanzi00}), a large dust extinction ($A_V\\sim 16$ mag; Thuan et al.\\ \\cite{thuan99}; Hunt et al.\\ \\cite{hunt01}; Plante \\& Sauvage \\cite{plante02}), and high dust temperature (Hunt et al.\\ \\cite{hunt01}; Dale et al.\\ \\cite{dale01}; Takeuchi et al.\\ \\cite{takeuchi03}). On the contrary, the star-forming regions in \\izw\\ are diffuse ($r_{\\rm SF}\\ga 100$ pc, $n\\la 100$ cm$^{-3}$), and contain no SSCs. Near-infrared \\H2 emission has not been detected (Hunt et al., private communication), and the dust extinction is moderate ($A_V\\sim 0.2$ mag; Cannon et al.\\ \\cite{cannon02}). We call a region with such properties ``passive'' following HHTIV. The similar metallicities of \\sbs\\ (active) and \\izw\\ (passive) imply that the chemical abundance is not a primary factor in determining the star-forming properties. We argue that the compactness of star-forming regions, which affect gas density, gas dynamics, and so on, is important in the dichotomy of active and passive modes.  Hirashita et al.\\ \\cite{hhf02} (hereafter HHF02) show that the IR luminosity, the dust mass, and the rich \\H2 content of \\sbs\\ can be explained through dust accumulation by successive SNe II in a compact ($r_{\\rm SF}\\la 100$ pc) region. Moreover, \\sbs\\ is not unique among BCDs; the BCDs with dense compact star-forming regions similar to \\sbs\\ were dubbed ``active'', and tend to be characterized by high surface brightness (see Figure 1 of HHTIV). The physical state of ISM is also similar among ``active'' BCDs. They always have large dust extinction, and a significant fraction of stellar radiation is reprocessed into IR; the dust temperature is also high and there is a ``hot'' component with 600--1000 K (Hunt et al.\\ \\cite{hunt02}). Dust properties can be further constrained by future IR observations for a large sample of BCDs (e.g., Takeuchi et al.\\ \\cite{takeuchi03}). HHTIV also demonstrate that there are BCDs with converse properties, namely ``passive'' ones, which are more diffuse, less dense, and of lower surface brightness. A representative of this category is \\izw. Contrary to ``active'' BCDs, IR dust emission has not been detected so far in ``passive'' BCDs. Although their star formation rate is lower than ``active'' BCDs, ``passive'' BCDs are actually forming stars, and are completely different from passively evolving galaxies such as ellipticals.  The above clear difference in dust properties implies that in addition to the compactness, dust should be considered as a key to understand the ``active'' and ``passive'' modes, and hence to understanding how star formation occurs in extremely low-metallicity environments. We consider the role of dust in various compactness of regions by using our theoretical framework. This paper is organised as follows. First, in Section \\ref{sec:model} we explain the model that describes the evolution of dust content and gas state in a star-forming region. Then, in Section \\ref{sec:result}, we present our results concerning the differences between ``active'' and ``passive'' star-forming regions. In Section \\ref{sec:bcds}, we discuss the interpretation of observational properties of active and passive BCDs in the context of our model, and consider in particular the two prototypes, \\izw\\ and \\sbs. Implications for high-redshift star formation are described in Section \\ref{sec:highz}, and our conclusions are presented in Section \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  In this paper, we have theoretically modeled the ``active'' and ``passive'' star formation modes observed in metal-poor BCDs. The ``active'' mode is characterised by dense and compact ($n\\ga 500~{\\rm cm}^{-3}$ and $r_{\\rm SF}\\la 50$ pc) star-forming regions. On the contrary, the ``passive'' mode takes place in diffuse ($n\\la 100~{\\rm cm}^{-3}$ and $r_{\\rm SF}\\ga 100$ pc) star-forming regions.  In the dense regions observed in ``active'' BCDs, the gas free-fall time is typically shorter than $\\sim 5$ Myr. Such a short free-fall time can enhance star formation activity and generate an efficient supply of dust from SNe II. An accumulation of dust in such a compact region leads to a large dust optical depth, and the region consequently becomes luminous in the IR. Even in the environment of active star formation, the gas retains the physical conditions favourable for the gas collapse: cool ($\\la 1000$ K) and highly molecular ($\\fH2\\ga 10^{-2}$). The above characteristics explain the properties of star-forming regions in the BCDs categorised as ``active'' in HHTIV, especially SBS 0335$-$052 (i.e., IR luminous, containing hot dust, rich in molecules, compact, and dense).  In ``passive'' BCDs, the dynamical time of diffuse regions is longer than $10^7$ yr, and the star formation rate can be as low as $\\la 0.1~M_\\odot$. The increase of dust optical depth is milder and therefore ``passive'' star-forming regions become IR luminous much later in their evolution, $> \\ 10^7$ yr. Such a passive region has a low $\\la 10^{-4}$ molecular content, and it would be difficult to detect \\H2 in passive BCDs.  However, the physical state of gas is strongly affected by the size and spatial distribution of grains. With a fixed total dust mass, if the grain radius is large, the optical depth of dust against the UV light is small. This means for large grains ($\\sim$0.01~$\\mu$m), the gas temperature and ionisation degree tend to be high, and the molecular fraction tends to be low. As for the spatial distribution of grains, efficient shielding of UV light takes place in a screen geometry, which leads to efficient cooling and molecule formation. On the contrary, in the mixed geometry, because of inefficient shielding, the gas is heated to nearly $10^4$ K, and the molecules are efficiently dissociated.  We also discussed the evolution of dust-to-gas ratio and metallicity. In particular, our model is consistent with the observations of \\izw\\ and \\sbs, but future observations of dust in metal-poor (${\\rm [O/H]}\\la -1$) BCDs are needed for further constraints. The consistency also implies that around 10--20 \\% of metals supplied from SNe are in the dust phase.  We have shown how differences in two physical parameters of a star-forming region, its size and density, can lead to substantially different evolution over time. The distinction has been made between ``active'' and ``passive'' modes, but such a dichotomy is perhaps misleading, since they are rather extremes on a continuum. Pressure must also drive evolution of a star-forming region, and may determine a region's initial size and density. However, after the onset of star formation, dust shielding of UV photons determines the fate of star-forming regions, which finally bifurcates into active and passive regimes."
    },
    "0404/astro-ph0404071_arXiv.txt": {
        "abstract": "{ In this paper we show normalized differential source counts n(S) at 408~MHz and 1.4~GHz of radio sources separately for FRI and FRII classes with extended and compact  morphologies. The maps from the FIRST, NVSS, and WENSS surveys are used to define the source morphology and flux density. The counts provide a basis for a direct test as well as constraining the cosmological evolution of powerful extragalactic radio sources in terms of the dual-population model (Jackson \\& Wall 1999), where radio sources of Fanaroff-Riley (1974) types I and II are regarded as two physically separate types of active galactic nuclei (AGN). The predicted count values are compared with the observational data to find the best fits for the evolution and beaming parameters, and to further refine the model. ",
        "introduction": "Galaxies (and/or quasars) hosting active galactic nuclei (AGN) are powerful extragalactic radio sources which produce jets and extended radio emitting regions (lobes) of plasma (Scheuer 1974). Most of these sources can be easily separated into two distinct classes: (i) the edge-brightened FRII sources with well-collimated jets, radio lobes and prominent hot spots, and (ii) edge-dimmed FRI sources with bending jets and peaked radio emission near the nucleus. There is a clear division in the luminosity between FRI and FRII, which lie, respectively, below and above $\\rm L_{178}\\approx 2\\times10^{25}$ $\\rm W Hz^{-1} sr^{-1}$ (Fanaroff \\& Riley 1974). Owen \\& Ledlow (1994) demonstrated that this separation is reproduced in the two-dimensional radio-optical luminosity plane in an unambiguous way. Furthermore, it is now quite clear that FRII and FRI sources differ radically their cosmological evolution (Wall 1980), i.e. the density and/or luminosity evolution rate of FRII is high, while that of FRI seems to be rather low.  Radio source counts n(S) involving distributions of the other observational parameters (i.e. luminosity, spectral index, redshift) were used to constrain a number of numerical models of cosmological evolution of extragalactic radio sources (e.g. Wall et al. 1980; Peacock \\& Gull 1981; Condon 1984; Peacock 1985; Dunlop \\& Peacock 1990). All these models assumed two distinct populations of sources: steep-spectrum and flat-spectrum ones, which in principle were selected based on their spectral characteristics only. Since the late 70s, however,  there have been efforts to find a common scheme for some types of extragalactic radio sources (Rowan-Robinson 1977; Scheuer \\& Readhead 1979). These unification attempts aimed to explain the diversity of observational properties of extragalactic radio sources in terms of one simple model. The starting point for such models is the intrinsic structure of a radio source. Because they are not spherically symmetric, their observed form depends on the angle of view. The realization that orientation angle in AGN is of great importance for unification came with the discovery of observational superluminal motion in some compact radio sources (Cohen et al. 1971; Moffet et al. 1972) and its interpretation in terms of bulk relativistic movements of the jet's plasma (de Young 1972). It is obvious that when some emitting body moves in a relativistic manner, the radiation received by an observer is a function of the angle between the line of sight and the direction of motion. The small fraction of objects with compact morphologies are those viewed with their axes inclined at small angles to the line of sight. Furthermore, the beamed sources should be linked to an unbeamed parent population. The unified scheme of AGN has been quite successful in explaining differences between various classes of radio galaxies on the basis of the jet luminosity, its speed, relativistic beaming effect, orientation towards the observer, and environmental effects (for a review, see Urry \\& Padovani 1995).  Recently,  Wall and Jackson (Wall \\& Jackson 1997; Jackson \\& Wall 1999) [hereafter WJ] introduced the {\\it dual-population unified scheme}. Their model consists of two parent populations of powerful extragalactic radio sources, FRI and FRII, which appear (dependent on the viewing angle) as lobe-dominated (LD) steep-spectrum sources, or core-dominated (CD) flat-spectrum sources.  Our main goal is to test the dual-population model of FRI and FRII radio source cosmological evolution and to refine its parameters with new well-defined samples of sources at two radio frequencies and within a wide range of radio flux density. The recent radio surveys with sufficient sensitivity, spatial resolution and flux-density limit, i.e. FIRST (Becker et al. 1995), NVSS (Condon et al. 1998)  and WENSS (Rengelink et al. 1997), allowed us to determine morphologies, flux densities, and spectral indices for a large enough number of objects and thereby proved to be suitable for constructing such samples.  In the following section the complete source samples used in this analysis and the counts of extragalactic radio sources n(S) based on data at 1.4~GHz and 408~MHz are described (Sect. 2). In Sect. 3 the application of the dual-population numerical model of radio source evolution is shown. In Sect. 4 the results of modelling of the counts n(S) compared to the observational data and discussion are given. Finally, Sect. 5 contains the summary. Throughout the paper we assume $\\rm H_{0}=50$ $\\rm km s^{-1} Mpc^{-1}$ and the Einstein-de Sitter cosmological model ($\\rm \\Omega=\\Omega_{m}=1$). ",
        "conclusions": "Following the procedure given in the preceding section we derive the normalized differential counts and the source fractions. The set of parameter values originally proposed by WJ is called `model A' (see Table 3). It is necessary to modify these values to fit the observational data. The best-modeled fit to the source counts at 1.4~GHz and 0.408~GHz is called `model B' and its parameter values are given also in Table 3. The values are determined using statistical methods, by minimizing $\\chi^{2}$ evaluated between the observed and modeled source distributions, both at 1.4 and 0.408~GHz. The value of $\\chi^{2}$ is evaluated as:  \\[\\rm \\chi^{2}= \\sum^{N}_{i=1}(\\frac{n_{data_{i}} - n_{model_{i}}}{err_{data_{i}}})^{2},\\]  \\noindent where $\\rm n_{data_{i}}$ is the value of the data in the ith bin, similarly $\\rm n_{model_{i}}$ and $\\rm err_{data_{i}}$ are the model value and data error in the ith bin, respectively. These all are determined in the logarithm space and the sum is taken over the N bins.\\\\  The observed and calculated n(S) distributions for 1.4 and 0.408~GHz are presented in Fig.~1 and Fig.~2, respectively. Fig. 3 shows the expected and observed source fraction.  The best-fitted  parameters of cosmological evolution and relativistic beaming that we estimate are slightly different from those obtained by WJ on the basis of the 5~GHz data. However, it is important to note that our calculations do not strictly follow all the steps proposed by Jackson \\& Wall (1999). For example, we do not demand the evolution function  E(L, z) to peak at $\\rm z_{c}$/2. Furthermore, our approach does not involve any dependence of the intrinsic core-to-extended flux ratio $\\rm R_{c}$ on radio power. We estimate that the transition between evolving and non-evolving FRII sources appears for luminosities $\\rm L_{151} = 10^{24.5}$ $\\rm (W Hz^{-1} sr^{-1})$, whereas objects with $\\rm L_{151} > 10^{27.95}$ $\\rm (W Hz^{-1} sr^{-1})$ show maximal evolution with the rate of $\\rm M_{max}$=12.2.\\\\ To fit the FRI observational data, it is necessary to introduce some luminosity evolution for these type of sources. However,  Urry \\&  Padovani (1995) report that the behavior of FRI radio galaxies is consistent with no evolution. Also, in the WJ scheme, the cosmological evolution of FRI does not explicitly appear, however, there are no fundamental aspects in their model forbidding the evolution of FRI sources in general. The idea that FRI should evolve is confirmed by the recent work of Snellen \\& Best (2001), who on inspecting Hubble Deep Field images came across two sources of FRI morphology. This could mean that this kind of source is significantly more abundant at high redshifts than at the present epoch. They suggest also that cosmological evolution of FRI sources could be similar to that of FRII (i.e. is radio luminosity dependent). We propose that the evolution of FRI sources is similar to that of FRII, (Sect. 3.1, eq. 4). The evolution rate M, which depends on luminosity, has the form:  \\vspace{3mm} \\noindent $\\rm M(L)=0$ \\hspace{16.8mm} for $\\rm L < L_{1}$,\\\\ $\\rm M(L)=M_{max}$ \\hspace{10mm} for $\\rm L \\geq L_{1}.$ \\vspace{3mm}   \\noindent We estimate that the small fraction of FRI with luminosities $\\rm L_{151}\\geq 10^{26.5}$ $\\rm (W Hz^{-1} sr^{-1})$ evolve with cosmic time. This is supported by the fact that at $\\rm L_{151}\\approx 10^{26.5}$ $\\rm (W Hz^{-1} sr^{-1})$ in the low-frequency samples the fraction of objects  with observed broad-line nuclei changes radically from $\\sim 0.4$ at higher radio luminosities to $\\sim 0.1$ at lower luminosities (Willott et al. 2000).\\\\ The best-fit relativistic beaming parameters ($\\rm R_{c}(\\nu)$ and $\\rm \\gamma$) obtained are similar to those estimated by Wall \\& Jackson (1997). They are also in agreement with those derived from observations. The observed core-to-extended flux ratio of individual sources spans a wide range of values from $\\sim10^{-5}$ for the lobe-dominated sources (Morganti et al. 1993) to $\\sim10^{3}$ for the core-dominated sources (Murphy et al. 1993). The observations show also that there is a large spread of Lorentz factor $\\gamma$ values from 2 to 20 (Vermulen \\& Cohen 1994), while some numerical simulations give $4\\leq\\gamma\\leq40$ (Urry \\& Padovani 1995). However, the Lorentz factor value of FRII sources should be greater from those of FRI (Urry \\& Padovani 1995).\\\\ Some recent papers (e.g. Caccianiga et al. 2002) suggest the existence of another class of sources with compact morphologies that do not originate from FRI or FRII populations. Our investigation shows that there is no explicit need to introduce any additional population of compact powerful radio sources, apart from the beaming version of FRI and FRII parent populations. However, if the values of beaming parameters given here are overestimated, then the model predictions for compact sources will be too small when compared to the data.  \\begin{table}[] \\begin{center} \\caption{Values of beaming and evolution parameters  for models A and B.} \\vspace{1mm} \\begin{tabular}{lrrrr}  {\\bf Parameter}       &\\multicolumn{2}{c}{\\bf FRI} &\\multicolumn{2}{c}{\\bf FRII}\\\\ &\\multicolumn{1}{c}{\\bf  A}&\\multicolumn{1}{c}{\\bf  B} &\\multicolumn{1}{c}{\\bf  A} &\\multicolumn{1}{c}{\\bf  B}                \\\\ {\\bf $\\rm \\gamma$}          & 10.3   ($\\dagger$)  & 10.3  & 20.0    ($\\dagger$)  & 20.0         \\\\ {\\bf $\\rm R_{c}(\\nu)(*)$}   & 0.010  ($\\dagger$)  & 0.010 & 0.004   ($\\dagger$)  & 0.006        \\\\ {\\bf $\\rm \\log(L_{1} (**))$}& --                  & 26.50 & 25.44   ($\\ddagger$) & 24.50        \\\\ {\\bf $\\rm \\log(L_{2} (**))$}& --                  & --    & 27.34   ($\\ddagger$) & 27.95        \\\\ {\\bf $\\rm M_{max}$}         & --                  & 5.0   & 10.93   ($\\ddagger$) & 12.20        \\\\ {\\bf $\\rm z_{max}$}         & 5.0    ($\\ddagger$) & 5.0   &  5.62   ($\\ddagger$) &  5.62        \\\\  \\end{tabular} \\end{center} $\\dagger$  the value is taken from Wall \\& Jackson (1997), and\\\\ $\\ddagger$ the value is taken from Jackson \\& Wall (1999),\\\\ (*) $\\rm \\nu=5~GHz$, (**) $\\rm W Hz^{-1} sr^{-1}$ \\end{table}   \\begin{figure*}[] \\resizebox{\\hsize}{130mm}{\\includegraphics{mjf1.ps}} \\caption{Normalized differential counts of sources of different morphological types at {\\bf 1.4~GHz}. Abscissa: log flux-density (Jy). Ordinate: log differential number of sources multiplied by $\\rm S^{2.5}[sr^{-1}Jy^{1.5}]$. Observational data are provided in the form of stepped curve with vertical error bars. The solid  curve shows the counts predicted from model B and the dotted solid one -- counts from model A. The respective panels give the counts for {\\bf a)} all sources in the sample, {\\bf b)} compact FRI and compact FRII sources, {\\bf c)} FRI lobe-dominated sources, and {\\bf d)}  FRII lobe-dominated and CSS sources, (for details, see the text).} \\label{counts14} \\end{figure*}  \\begin{figure*}[] \\resizebox{\\hsize}{130mm}{\\includegraphics{mjf2.ps}} \\caption{Normalized differential counts of sources of different morphological types at {\\bf 408~MHz}. Abscissa: log flux-density (Jy). Ordinate: log differential number of sources multiplied by $\\rm S^{2.5}[sr^{-1}Jy^{1.5}]$. Observational data are provided in the form of stepped curve with vertical error bars. The solid  curve shows the counts predicted from model B and the dotted solid one -- counts from model A. The respective panels give the counts for {\\bf a)} all sources in the sample, {\\bf b)} compact FRI and compact FRII sources, {\\bf c)} FRI lobe-dominated sources, and {\\bf d)} FRII lobe-dominated and CSS sources, (for details, see the text).} \\label{counts400} \\end{figure*}  \\begin{figure*}[] \\vspace{50mm} \\resizebox{\\hsize}{80mm}{\\includegraphics{mjf3a.ps} \\includegraphics{mjf3b.ps}} \\caption{Fraction of {\\bf a)} compact flat-spectrum sources, {\\bf b)} FRI lobe-dominated sources, and {\\bf c)} FRII  lobe-dominated together with CSS sources in the sample at 1.4~GHz -- {\\bf left panels}, and at 408~MHz -- {\\bf right panels}. The solid curve shows the source ratio predicted from model B.} \\end{figure*}"
    },
    "0404/astro-ph0404592_arXiv.txt": {
        "abstract": "A major uncertainty in the spectroscopic dating of extragalactic globular clusters concerns the degenerate effect that age and horizontal branch morphology have on the strength of Balmer lines. In this {\\it Letter} we show that the ratio between the equivalent widths of \\hdf\\ and \\hb\\ is far more sensitive to horizontal branch morphology than to age, thus making it possible to break the degeneracy.  We show that it is possible to distinguish intermediate-age globular clusters from those whose Balmer lines are strengthened by the presence of blue horizontal branch stars, {\\it purely on the basis of the clusters' integrated spectra}. The degeneracy between age and horizontal branch morphology can be lifted with \\hb\\ and \\hdf\\ line strengths from spectra with S/N $\\simgreat$ 30/${\\rm\\AA}$, which is typical of current studies of integrated spectroscopy of extragalactic globular clusters. ",
        "introduction": "In the last decade, the use of equivalent widths (EWs) of Balmer absorption lines in the integrated spectra of old stellar populations (SPs) has become a standard procedure for estimating their light-weighted mean ages (Worthey 1994). The rationale of the method is to utilize the sensitivity of Balmer lines to the temperature and luminosity of turn-off stars, which depend strongly on the age of a SP.  However, a well known caveat is the existence of old hot stars that may be sufficiently bright and numerous to boost the EWs of Balmer lines, mimicking the signature of young stars in the integrated spectra of old SPs. That is the case of blue horizontal branch (BHB) stars, which are present in such old systems as Galactic globular clusters (GCs) with a range of metallicities (see Moehler 2001 for a review), the halo and bulge fields (e.g. Kinman et al.\\ 2000), the cores of Local Group galaxies (Brown et al.\\ 1998; Brown et al.\\ 2000; Brown 2003), and old Galactic clusters like NGC~6791 (e.g. Peterson \\& Green 1998). So far the only way of unequivocally distinguishing between an old SP that hosts BHB stars from a truly young SP is by construction of color-magnitude diagrams (CMDs), which of course can only be achieved for nearby systems, within the Local Group.  Theoretical studies (Freitas Pacheco \\& Barbuy 1995; Lee et al.\\ 2000; Maraston \\& Thomas 2000) have shown that the impact of BHB stars on the integrated spectra of old SPs can be significant, and in particular it may affect the ages inferred from the EWs of Balmer lines. However, in view of the lingering theoretical uncertainties regarding the morphology of the horizontal branch (HB) and its connection with mass loss on the red giant branch (Catelan 2000), an accurate assessment of the contribution of BHB stars to the integrated optical light of old SPs still depends on the availability of CMDs.  In this {\\it Letter} we present the first results of an ongoing project to calibrate our SP synthesis models with high quality integrated spectra of star clusters. While this has already been attempted by a number of groups (Beasley et al.\\ 2002; Schiavon et al. 2002a,b; Puzia et al.\\ 2002; Maraston et al.\\ 2003; Leonardi \\& Rose 2003), our work benefits from the combination of a homogeneous set of high S/N integrated GC spectra newly collected by our group, with the HST CMDs of several tens of Galactic GCs, most of them reaching the GC turn-off, by Piotto et al.\\ (2002).  We demonstrate the ability of well-calibrated models, combined with high S/N spectra, to single out the contribution of BHB stars to the integrated light of GCs. We use a combination of an Fe-sensitive index and an index comprised of the ratio of two Balmer line EWs to uniquely constrain the BHB contribution to the integrated light. This enables us to distinguish, {\\it purely on the basis of their integrated spectra}, truly intermediate-age or young clusters from those which are old, but whose Balmer lines are strengthened by the contaminating light of BHB stars. We envisage a direct application of our method to studies of extragalactic GC systems, where the determination of GC ages and metal abundances can yield insight into the star formation and merger histories of the host galaxies (e.g. Cohen, Blakeslee \\& C\\^ot\\'e 2003; Larsen et al. 2003; Hempel et al.\\ 2003). ",
        "conclusions": ""
    },
    "0404/astro-ph0404247_arXiv.txt": {
        "abstract": "We present the results of photometric measurements from images of the LMC globular clusters NGC 1928, 1939 and Reticulum taken with the Advanced Camera for Surveys on the {\\em Hubble Space Telescope}. Exposures through the F555W and F814W filters result in high accuracy colour-magnitude diagrams (CMDs) for these three clusters. This is the first time that CMDs for NGC 1928 and 1939 have been published. All three clusters possess CMDs with features indicating them to be $> 10$ Gyr old, including  main sequence turn-offs at $V\\sim 23$ and well populated horizontal branches (HBs). We use the CMDs to obtain metallicity and reddening estimates for each cluster. NGC 1939 is a metal-poor cluster, with $[{\\rm Fe}/{\\rm H}] = -2.10 \\pm 0.19$, while NGC 1928 is significantly more metal-rich, with $[{\\rm Fe}/{\\rm H}] = -1.27 \\pm 0.14$. The abundance of Reticulum is intermediate between the two, with $[{\\rm Fe}/{\\rm H}] = -1.66 \\pm 0.12$ -- a measurement which matches well with previous estimates. All three clusters are moderately reddened, with values ranging from $E(V-I) = 0.07 \\pm 0.02$ for Reticulum and $E(V-I) = 0.08 \\pm 0.02$ for NGC 1928, to $E(V-I) = 0.16 \\pm 0.03$ for NGC 1939. After correcting the CMDs for extinction we estimate the HB morphology of each cluster. NGC 1928 and 1939 possess HBs consisting almost exclusively of stars to the blue of the instability strip, with NGC 1928 in addition showing evidence for an extended blue HB. In contrast, Reticulum has an intermediate HB morphology, with stars across the instability strip. Using a variety of dating techniques we show that these three clusters are coeval with each other and the oldest Galactic and LMC globular clusters, to within $\\sim 2$ Gyr. The census of known old  LMC globular clusters therefore now numbers $15$ plus the unique, somewhat younger cluster ESO121-SC03. The NGC 1939 field contains another cluster in the line-of-sight, NGC 1938. A CMD for this object shows it to be less than $\\sim 400$ Myr old, and it is therefore unlikely to be physically associated with NGC 1939. ",
        "introduction": "\\label{s:intro} The Large and Small Magellanic Clouds (LMC/SMC) possess extensive systems of rich stellar clusters. These objects exhibit a much wider variety in age, structure, environment and mass than do Galactic clusters, and this, combined with their relatively close proximity, has rendered them central to a surprising number of fields of modern astrophysics -- from star and cluster formation, and stellar evolution, to gravitational dynamics, galactic evolution, and distance scale measurements. They are also vital probes and tracers of the chemical and dynamical evolution of the LMC and SMC themselves. It is therefore important to understand how their properties, and in particular their ages and abundances, are distributed.  It has long been known that the LMC, which contains the more numerous cluster system of the two Clouds, houses a small sub-population of extremely ancient objects. Although early studies revealed half a dozen LMC clusters to possess colour-magnitude diagrams (CMDs) with features similar to those of the Galactic globular clusters, it has only been since the advent of the {\\em Hubble Space Telescope} ({\\em HST}) that imaging resolution and sensitivity has been sufficiently high as to allow accurate colour-magnitude diagrams (CMDs) suitable for relative age dating. A number of relatively recent studies have demonstrated that the number of LMC clusters coeval with each other and the oldest Galactic globular clusters is somewhat more than a dozen -- NGC 1466, NGC 2257, and Hodge 11 (e.g., Johnson et al. \\shortcite{johnson:99}); NGC 1754, 1835, 1898, 1916, 2005, and 2019 (e.g., Olsen et al. \\shortcite{olsen}); and NGC 1786, 1841 and 2210 (e.g., Brocato et al. \\shortcite{brocato}). Several published CMDs show the remote outer cluster Reticulum to be very old \\cite{johnson:99,marconi,monelli}; however a full age analysis is yet to be published for this cluster. In addition there are two more clusters located in the LMC bar region -- NGC 1928 and 1939 -- which integrated spectroscopy suggests could be old \\cite{dutra}, but which are so compact and lie against such highly crowded LMC fields that it has not previously been possible to obtain accurate CMDs for them (see for example, the search for old LMC clusters conducted by Geisler et al. \\shortcite{geisler}). The census of old LMC clusters is therefore still incomplete.  We have used the Advanced Camera for Surveys (ACS) on {\\em HST} to obtain images of NGC 1928, 1939, and Reticulum as part of program 9891 -- a snapshot survey of $40$ LMC and $40$ SMC clusters. This program is primarily designed to extend the detailed investigation of Mackey \\& Gilmore \\shortcite{sbp1,sbp2} concerning the structural evolution of Magellanic Cloud clusters. However, the ACS observations are of sufficient quality and resolution as to allow colour-magnitude diagrams (CMDs) to be constructed for NGC 1928 and NGC 1939 for the first time, as well as a high quality CMD for Reticulum. In this paper, we present the results of these observations (Section \\ref{s:obsred}) and the measured colour-magnitude diagrams (Section \\ref{s:colourmag}). Abundances and reddenings are derived for each cluster, and it is demonstrated that NGC 1928 and 1939, along with Reticulum, are coeval with both well studied Galactic globular clusters (such as M92, M3 and M5) and other old LMC objects (such as NGC 2257 and Hodge 11) (Section \\ref{s:clusterprop}). ",
        "conclusions": "\\label{s:summary} We have used ACS/WFC snapshot observations to obtain colour-magnitude diagrams for the LMC clusters NGC 1928, 1939 and Reticulum. This is the first time that CMDs for NGC 1928 and 1939 have been published. These two CMDs suffer from very dense field star contamination requiring a thorough subtraction algorithm. Using the final CMDs we obtained photometric reddening and metallicity measurements for all three clusters. NGC 1939 is one of the most metal-poor LMC bar clusters, with $[{\\rm Fe}/{\\rm H}] = -2.10 \\pm 0.19$, while NGC 1928 is significantly more metal-rich, with $[{\\rm Fe}/{\\rm H}] = -1.27 \\pm 0.14$. Reticulum is of a more intermediate abundance, with $[{\\rm Fe}/{\\rm H}] = -1.66 \\pm 0.12$. This measurement matches well the previous estimates for this cluster.  All three clusters possess CMDs with features characteristic of the oldest Galactic globular clusters -- main sequence turn-offs at $V\\sim 23$, and well populated horizontal branches. Both NGC 1928 and 1939 possess very blue HB morphologies, with little or no population stretching to the red of the blue edge of the instability strip. In contrast, Reticulum has a HB populated right across the instability region. To quantify the ages of the three clusters we employed a variety of differential dating techniques, comparing their CMDs to a set of three of the oldest Galactic globular clusters, and three of the oldest LMC globular clusters. We conclude that the entire set of clusters is coeval to within approximately $2$ Gyr. This work firmly establishes NGC 1928 and 1939 as members of the LMC globular cluster population, confirming the conclusion obtained by Dutra et al. \\shortcite{dutra} from integrated spectroscopy. The LMC globular cluster census therefore now numbers $15$, in two distinct groups -- the outer clusters NGC 1466, 1841, 2210, 2257, Hodge 11, and Reticulum; and the inner (bar) clusters NGC 1754, 1786, 1835, 1898, 1916, 1928, 1939, 2005, and 2019. The only other LMC cluster older than the lower end of the age gap is ESO121-SC03, at $\\sim 9$ Gyr."
    },
    "0404/astro-ph0404467_arXiv.txt": {
        "abstract": "The SC and CS stars are thermal-pulsing AGB stars with C/O ratio close to unity. Within this small group, the Mira variable BH Cru recently evolved from spectral type SC (showing ZrO bands) to CS (showing weak C$_2$). Wavelet analysis shows that the spectral evolution was accompanied by a dramatic period increase, from 420 to 540 days, indicating an expanding radius.  The pulsation amplitude also increased.  Old photographic plates are used to establish that the period before 1940 was around 490 days. Chemical models indicate that the spectral changes were caused by a decrease in stellar temperature, related to the increasing radius. There is no evidence for a change in C/O ratio. The evolution in BH Cru is unlikely to be related to an on-going thermal pulse.  Periods of the other SC and CS stars, including nine new periods, are determined. A second SC star, LX Cyg, also shows evidence for a large increase in period, and one further star shows a period inconsistent with a previous determination.  Mira periods may be intrinsically unstable for C/O$\\,\\approx 1$; possibly because of a feedback between the molecular opacities, pulsation amplitude, and period. LRS spectra of 6 SC stars suggest a feature at $\\lambda>15\\mu$m, which resembles one recently attributed to the iron-sulfide troilite. Chemical models predict a large abundance of FeS in SC stars, in agreement with the proposed association. ",
        "introduction": "The observed properties of stars on the Asymptotic Giant Branch (AGB) are largely determined by the ratio of carbon to oxygen in their atmospheres. Carbon abundances in AGB stars increase due to dredge-up following their regular thermal pulses.  After a series of these events, which happen every $10^4$--$10^5$\\,yr, the enhanced carbon may cause a transition from an oxygen-rich star to a carbon-rich star, via the intermediate S-stars \\citep[e.g.,][]{LE84}.  The SC stars \\citep{KB80} form a continuous spectral sequence intermediate between the S and C stars \\citep{CF71}.  They show very strong sodium D-lines, and strong CN bands \\citep{S73}, and either weak ZrO bands (SC stars) or C$_2$ bands (CS stars). The molecular abundances indicate a C/O number ratio very close to unity, so that CO formation leaves little C or O for the other molecules.  The SC star BH~Crucis has simultaneously been found to show a unique combination of a lengthening period \\citep*{BMV88, WM91, WIW95}, and evolution in spectral type from SC to CS \\citep{LE85}, with ZrO bands \\citep{CF71} disappearing and C$_2$ band appearing.  The spectral evolution has been interpreted as due to an increase in C/O ratio \\citep{Whi99}.  Both the period and abundance changes have been suggested to be caused by a recent thermal pulse \\citep{Whi99,WZ81}. BH Cru could therefore present a unique case of real-time AGB evolution.  \\begin{figure*} \\includegraphics[width=0.8\\textwidth, clip=true]{bhcru-visual-bin20.0125.wwz.ps} \\caption{\\label{bhcru.ps} The wavelet analysis for BH~Cru, 1970--2002, based on visual observations. Shown are: the light curve, the frequency, the semi-amplitude of the main frequency component and its period.} \\end{figure*}  In this paper we present an analysis of the light curve of BH~Cru, to quantify the period and amplitude variation.  Chemical modelling is used to investigate whether the change in relative molecular abundances can be explained as due to a decrease in effective temperature, at constant C/O ratio.  We also present new periods for other SC stars, including two stars which may show period changes similar to BH Cru. ",
        "conclusions": "Our investigation of the evolving variable BH~Cru shows an increase in its period of 25\\%\\ within 25\\,yr. The period has stabilized at about 540 days. The visual semi-amplitude has increased simultaneous with the period increase, up to a value of 1.25\\,mag.  The changing period shows that the radius of the star has increased, and the temperature decreased, the latter confirmed by a slow reddening.  The spectral type of BH Cru changed simultaneously from SC to CS.  Chemical equilibrium modelling, both for the potosphere and a hydrostatic atmosphere, explains this as due to the decrease in the effective temperature. The lower temperature favours formation of C$_2$ and causes the fractional abundances of ZrO and YO to decrease.  Our calculations explore a range of C/O$>1$, for which oxides are still found although at reduced abundances.  The distinction between SC and CS stars does not require an evolution from C/O$<1 \\longrightarrow >1$, as sometimes is suggested, although this can also play a role.  Infrared spectra suggest the possible presence of the iron-sulfide troilite in SC stars. Chemical calculations show that FeS is abundant in the photosphere, and in the absence of silicates or carbon grains, may be an important dust component.  We determined new periods for a number of SC/CS stars, including three for stars with no previously known period. One star, LX~Cyg, shows a much longer period than previously determined, and is confirmed as undergoing period evolution similar to BH~Cru.  VX Aql also shows an inconsistency between the present period and one previously reported. Among the few Miras known with evolving periods, BH~Cru was unique in showing an {\\it increasing} period. Its time scale for the evolution is ten times faster than that of the well-studied case of R Hya.  With LX~Cyg, a second case of rapid period increase is now known.  Period changes in Miras are commonly attributed to thermal pulses.  This appears unlikely in BH Cru, because other SC stars show the same type of changes.  As an alternative model we suggest the possiblity that a feedback between molecular opacities, pulsation amplitude and periods cause unstable periods among the long-period SC stars."
    },
    "0404/astro-ph0404184_arXiv.txt": {
        "abstract": "It appears that the dynamical status of clusters and groups of galaxies is related to the large-scale structure of the Universe. A few interesting trends have been established:  \\noindent (1) {\\em The Cluster Substructure - Alignment Connection} by which clusters show a strong correlation between their tendency to be aligned with their neighbors and their dynamical state (as indicated by the existence of significant substructres).  \\noindent (2) {\\em The Cluster Dynamics -Cluster Clustering Connection} by which dynamically young clusters are more clustered than the overall cluster population.  \\noindent (3) {\\em The Cluster- Supercluster Alignment Connection} by which clusters of galaxies show a statistical significant tendency to be aligned with the projected major axis orientation of their parent supercluster.  \\noindent (4) {\\em The Galaxy Alignment - Cluster Dynamics Connection} by which red-sequence cluster bright galaxies show a significant trend to be aligned with their parent cluster major axis, especially in dynamically young clusters.  \\noindent (5) {\\em The Group Richness - Shape Connection} by which groups of galaxies are flatter the poorer they are.  These are strong indications that clusters develop in a hierarchical fashion by anisotropic merging of smaller units along the large-scale filamentary structures within which they are embedded. ",
        "introduction": "In the framework of the hierarchical model for the formation of cosmic structures, galaxy clusters are supposed to form by accretion of smaller units (galaxies, groups etc). After the epoch of mass aggregation (which depend on the cosmological model), violent relaxation processes will tend to alter the cluster phase-space configuration producing `regular', quasi-spherical, having smooth density profile clusters.  In the last decade due to the increased spatial resolution in X-ray imaging ({\\sc Rosat}, {\\sc Xmm}-{\\em Newton}, {\\sc Chandra}) and to the availability of wide-field cameras, many of the previously thought ``regular'' clusters have shown to be clumpy to some level, a fact that could have important consequences for structure formation theories, since the present fraction of dynamically young clusters, as well as the rate of cluster evolution (as measured for example by their luminosity and temperature functions and their morphology), are cosmology dependent (eg., Richstone, Loeb \\& Turner 1992, Evrard et al. 1993, Lacey \\& Cole 1996). Detailed optical and X-ray studies have yielded evidence for significant amounts of subclustering within local and distant clusters (eg. Buote \\& Tsai 1996; Plionis \\& Basilakos 2002; Jeltema et al. 2004) indicating that possibly a large fraction is still forming at the present time. Hints do exist for a very recent (within the last Gyr) dynamical evolution of the cluster population (Melott et al. 2001; Plionis 2002, see however Floor et al. 2003). Note however that the existence of substructure does not necessarily mean that the corresponding cluster is dynamically young due to the ambiguity of cluster post-merging relaxation times (see discussion in Plionis 2001).  Furthermore, the dynamical evolution of member galaxies and of the ICM gas is also an open issue. Star formation seems to be active in clusters showing substantial substructure and velocity gradients, as expected if a recent merger has taken place. For example, the fraction of blue galaxies is strongly correlated with cluster ellipticity (Wang \\& Ulmer 1997), while ellipticity is strongly correlated with the dynamical state of the cluster (eg. Kolokotronis et al. 2001). It appears that the violent merging events trigger star-formation, possibly through a multitude of different mechanisms; for example, the excess number of galaxy-galaxy interactions, the rapid variation of the cluster gravitational field (Bekki 1999), etc. An interesting observable that may be related to the dynamics of clusters is their tendency to be aligned with their nearest neighbor as well as with other clusters that reside in the same supercluster (eg. Bingelli 1982; Plionis 1994). Analytical (Bond 1986) and numerical work (eg. West et al. 1991, van Haarlem \\& van de Weygaert 1993, Splinter et al. 1999, Onuora \\& Thomas 2000; Faltenbacher et al. 2002; Knebe et al. 2004) have shown that such alignments occur in many hierarchical clustering models and are probably the result of an interesting property of Gaussian random fields which is the \"cross-talk\" between density fluctuations on different scales (eg. West 1994). ",
        "conclusions": "We have presented evidence, based on the {\\sc Apm} and Abell cluster samples as well as the {\\sc Uzc-Sssrs2} and {\\sc 2dfgrs} group samples, that there is a strong link between the dynamical state of clusters/groups of galaxies and their large-scale environment. Dynamically young clusters are significantly more aligned with their nearest neighbors and they are also much more spatially clustered. Clusters belonging in dense superclusters are preferentially aligned with their parent supercluster projected orientation and bright galaxies, in dynamically young clusters, are also orientated along their parent cluster major axis. These coherent orientation effects, from the scale of galaxies to that of superclusters support the hierarchical formation scenario in which clusters form by anisotropic merging along the large-scale filamentary structures within which they are embedded."
    },
    "0404/astro-ph0404190_arXiv.txt": {
        "abstract": "In a composite system of gravitationally coupled stellar and gaseous discs, we perform linear stability analysis for axisymmetric coplanar perturbations using the two-fluid formalism. The background stellar and gaseous discs are taken to be scale-free with all physical variables varying as powers of cylindrical radius $r$ with compatible exponents. The unstable modes set in as neutral modes or stationary perturbation configurations with angular frequency $\\omega=0$. The axisymmetrically stable range is bounded by two marginal stability curves derived for stationary perturbation configurations. By the gravitational coupling between the stellar and the gaseous disc components, one only needs to consider the parameter regime of the stellar disc. There exist two unstable regimes in general: the collapse regime corresponding to large-scale perturbations and the ring-fragmentation regime corresponding to short-wavelength perturbations. The composite system will collapse if it rotates too slowly and will succumb to ring-fragmentation instabilities if it rotates sufficiently fast. The overall stable range for axisymmetric perturbations is determined by a necessary $D-$criterion involving of the effective Mach number squared $D_s^2$ (i.e., the square ratio of the stellar disc rotation speed to the stellar velocity dispersion but scaled by a numerical factor). Different mass ratio $\\delta$ and sound speed ratio $\\eta$ of the gaseous and stellar disc components will alter the overall stability. For spiral galaxies or circumnuclear discs, we further include the dynamical effect of a massive dark matter halo. As examples, astrophysical applications to disc galaxies, proto-stellar discs and circumnuclear discs are discussed. ",
        "introduction": "% \\label{sect:intro} Axisymmetric instabilities in models of disc galaxies have been investigated extensively in the last century (e.g., Safronov 1960; Toomre 1964; Binney \\& Tremaine 1987; Bertin \\& Lin 1996). For a single disc of either gaseous or stellar content, Safronov (1960) and Toomre (1964) originally introduced a dimensionless $Q$ parameter to determine the local stability condition (i.e. $Q>1$) against axisymmetric ring-like disturbances in the usual Wentzel-Kramers-Brillouin-Jeffreys (WKBJ) or tight-winding approximation. A more realistic model of a disc galaxy would involve both gas and stars as well as an unseen massive dark matter halo, all interacting among themselves through the mutual gravitation.\\footnote{Magnetic field and cosmic-ray gas component are dynamically important on large scales (Fan \\& Lou 1996; Lou \\& Fan 1998a, 2003; Lou \\& Zou 2004 and references therein) in the galactic gas disc of interstellar medium (ISM) but are not considered here for simplicity.} Many theoretical investigations have been conducted along this track for a composite system of two coupled discs (Lin \\& Shu 1966; Kato 1972; Jog \\& Solomon 1984a, b; Bertin \\& Romeo 1988; Romeo 1992; Elmegreen 1995; Jog 1996; Lou \\& Fan 1998b). In these earlier treatments, either an approach of combined distribution function and fluid or the formalism of two fluids have been adopted in a WKBJ modal analysis. While these model results were initially derived in various galactic contexts, they can be applied or adapted also to, with proper qualifications, relevant self-gravitating disc systems including accretion discs in general, circumnuclear discs, protostellar discs, planetary discs and so forth.  The local WKBJ or tight-winding approximation has been proven to be a powerful technique in analysing disc wave dynamics. Meanwhile, theorists have long been keenly interested in a class of relatively simple disc models referred to as scale-free discs (Mestel 1963; Zang 1976; Lemos et al. 1991; Lynden-Bell \\& Lemos 1993; Syer \\& Tremaine 1996; Evans \\& Read 1998; Goodman \\& Evans 1999; Shu et al. 2000; Lou 2002; Lou \\& Fan 2002; Lou \\& Shen 2003; Shen \\& Lou 2003; Lou \\& Zou 2004; Lou \\& Wu 2004). Scale-free discs, where all pertinent physical variables (e.g., disc rotation speed, surface mass density, angular speed etc.) scale as powers of cylindrical radius $r$, have become one effective and simple vehicle to explore disc dynamics. Perhaps, the most familiar case is the so-called singular isothermal discs (SIDs) or Mestel discs with an isothermal equation of state and flat rotation curves (Mestel 1963; Zang 1976; Goodman \\& Evans 1999; Shu et al. 2000; Lou 2002; Lou \\& Shen 2003; Lou \\& Zou 2004). In contrast to the usual WKBJ approximation for perturbations, perturbations in axisymmetric scale-free discs can be treated, in some cases, globally and exactly without the local restriction (i.e., valid only in the short-wavelength regime). It is therefore possible to derive global properties of perturbations. Using scale-free disc models, Lemos et al. (1991) and Syer \\& Tremaine (1996) both studied the axisymmetric stability problem for a single disc and found that instabilities first set in as neutral modes or stationary configurations with angular frequency $\\omega=0$.  The main motivation of this paper is to examine the global axisymmetric stability problem in a composite system of two gravitationally coupled scale-free discs. As a more general extension to the previous two-SID analysis (Lou \\& Shen 2003; Shen \\& Lou 2003), we further consider a much broader class of rotation curves as well as the equation of state. This contribution gives an explicit proof that stationary configurations ($\\omega=0$) do mark the marginal stability in the two-fluid system, a cogent supplement to our recent investigation on stationary perturbation configurations (Shen \\& Lou 2004). ",
        "conclusions": "\\label{sect:summary}  The main thrust of this investigation is to model linear coplanar perturbations of axisymmetry ($m=0$) in a composite system of two-fluid scale-free discs with one intended for a stellar disc and the other intended for a gaseous disc. The two discs are dynamically coupled through the mutual gravitational interaction. In order to include the dynamical effect of a massive dark matter halo with axisymmetry, we further describe a composite system of two coupled partial discs (e.g. Syer \\& Tremaine 1996; Shu et al. 2000; Lou 2002; Lou \\& Shen 2003; Lou \\& Zou 2004; Lou \\& Wu 2004). In a global perturbation analysis, we show that axisymmetric instabilities set in as stationary perturbation configurations with $\\omega=0$. The marginal $D_s^2$ stability curves (characterized by the stationary configurations) delineate two different unstable regimes, namely, the collapse regime for large-scale perturbations and the ring-fragmentation regime for short-wavelength perturbations. Apparently, the composite disc system becomes less stable than a single-disc system and can be unstable with the two discs being stable separately (Lou \\& Fan 1998b). In our analysis, stationary perturbation configurations turn out to be more than just an alternative equilibrium state, especially in view of the stability properties.  The basic results of this paper are generally applicable to self-gravitating disc systems with or without axisymmetric dark matter halos. The two-fluid treatment contains more realistic elements than a single-disc formulation in the context of disc galaxies. In addition to astrophysical applications to disc galaxies, the studies presented here can be valuable for exploring the dynamical evolution of protostellar discs and circumnuclear discs.  In the context of a proto-stellar disc, it is the usual case to ignore the self-gravity effect. By considering the self-gravity of a composite disc system, our analysis indicates several qualitative yet interesting results. For example, if the initial disc system rotates sufficiently fast, the ring fragmentation (see the upper-right part of Fig. 1) can occur at relatively small radial scales. By further non-axisymmetric fragmentations, these condensed rings of materials may eventually become birthplaces of planets. On the other hand, if the initial disc system rotates sufficiently slow, then gravitational collapse can be induced by perturbations of relatively large radial scales (see the lower left corner of Fig. 1). Once such a perturbation develops in the background equilibrium disc, it grows rapidly and destabilizes the disc. Subsequently, the system undergoes global Jeans collapse to form a central young stellar object. Finally, if the initial disc system rotates in a regime stable against all axisymmetric perturbations (see Fig. 1), there might be two possibilities: (1) the composite disc system might become unstable caused by non-axisymmetric perturbations (not analyzed here) and (2) the disc rotation may be gradually slowed down by some braking mechanisms (e.g. magnetic field not included here and outflows or winds) and the disc eventually succumbs to a central collapse induced by large-scale perturbations.  Likewise, in the context of a circumnuclear disc around the center of a galaxy, we can readily conceive similar physical processes in parallel. One important distinction is that a dark-matter halo should play an important dynamical role so that a formulation of a partial composite disc system would be more appropriate. Here, the ring fragmentation can be induced by relatively small-scale perturbations in a disc system of sufficiently fast rotation. Such a ring of relatively dense materials around the galactic center would be a natural birthplace for circumnuclear starburst activities (e.g. Lou et al. 2001). Depending on the evolution history of a circumnuclear disc system, it may be stable initially and gradually lose angular momentum by generating and damping spiral magnetohydrodynamic (MHD) density waves (Lou et al. 2001). When the disc rotation becomes sufficiently slow, Jeans collapse induced by large-scale perturbations can set in to form a bulge or a supermassive black hole.  In summary, our global analysis shows the possible presence of an evolution stage for a composite disc system against all axisymmetric coplanar perturbations. More importantly, we reveal the parameter regime of ring fragmentation and the parameter regime of large-scale collapse. Astrophysical applications are discussed in the contexts of disc galaxies, proto-stellar discs and circumnuclear disks."
    },
    "0404/astro-ph0404473_arXiv.txt": {
        "abstract": "We present a new SCUBA image of the cluster \\ms, which exhibits strong, extended sub-mm flux at 850\\mum.  The most striking feature in the map is an elongated region of bright sub-mm emission, with a flux density of ${\\sim}\\,10\\,$mJy over several beam-sizes.  This region is apparently coincident with a previously known optical arc (which turns out to be a strongly lensed Lyman Break Galaxy at $z=2.911$), as well as with a newly identified multiply imaged ERO (Extremely Red Object) pair predicted to be at a similar, if not identical redshift. By combing a detailed lensing model with deep images from \\hst, \\chandra, CFHT, JCMT, and spectra from the VLT, we conclude that both the strongly lensed optical arc and ERO systems have properties consistent with known sub-mm emitters.  Using a simple model for the two sources, we estimate that the multiply lensed EROs contribute the majority of the flux in the SCUBA lensed arc. Correcting for the lensing amplification, we estimate that the inherent 850\\mum\\ fluxes for both objects are \\lsim0.4\\,mJy. If the LBG and ERO pair are truly at the same redshift, then they are separated by only $\\sim10\\,$kpc in the source plane, and hence constitute an interacting system at $z\\sim2.9$. Higher angular resolution observations in sub-mm/mm will permit us to more accurately separate the contribution from each candidate, and better understand the nature of this system. ",
        "introduction": "\\label{sec:intro} Since the installation of the Submillimetre Common User Bolometer Array (SCUBA; Holland et al.~1999) at the JCMT telescope, hundreds of luminous dusty galaxies have been detected (see Blain et al.~2002 for a review). Such observations are still far from routine, however, with the detection rate in random fields being typically one or two sources of modest signal-to-noise ratio (SNR) per shift of telescope time. Bright sources, which are easier to obtain detailed sub-mm/mm follow-up observations for using heterodyne receivers or bolometers tuned to other wavelengths, are correspondingly rarer.  This, coupled with the fact that the dust responsible for the sub-mm luminosity absorbs radiation at other wavelengths, means that SCUBA detected galaxies are often extremely faint in the optical.  Hence, redshifts, morphologies and spectral energy distributions (SEDs) have proven elusive for most sources.  One technique which can help, pioneered by Smail, Ivison \\& Blain (1997), is to use the lensing amplification of rich galaxy clusters to boost the detectability of background sub-mm galaxies. The rate at which sources are detected is increased by a factor ${\\sim}\\,3$, with an improved prospect of discovering the occasional bright object at relatively high SNR.  Here we present detailed observations of just such a source in the field of the massive cluster \\ms, which we discovered in an earlier survey of 9 rich cluster fields (Chapman et al.~2002a, hereafter C02).  This paper uses new SCUBA data, as well as rigorous chop deconvolution, in order to produce a map free from the chopping artifacts common in sub-mm observations.  The sub-mm emission in this image is extended along a giant optical arc discovered in this cluster by Luppino et al.~(1999, hereafter L99), and thus there is little doubt that the SCUBA emission is lensed as well.  In this paper, we combine deep optical, near-IR (NIR), and X-ray data with a detailed lensing model in order to interpret the sub-mm emission in this remarkable SCUBA field. ",
        "conclusions": "We have presented broad-band sub-mm and near-IR images of the cluster \\ms\\ and attempted to understand the nature of the highly extended sub-mm emission.  The only other case of a multiply lensed sub-mm detection in a cluster field  has recently been reported in Abell 2218 (Kneib et al.~2004).  They identified the source of the triple sub-mm image as a LBG galaxy (showing a complex structure, indicative of a close merger) in a $z\\sim2.5$ group.  In the case of \\ms, the evidence suggests that the sub-mm emission is coming from at least two components: a blue arc at $z=2.911$ and a multiply imaged ERO pair predicted from the lens model to be at $z\\sim2.85$. Using a simple model for the 2 sources, we estimate that the multiply lensed EROs contribute $\\sim$2/3 of the observed flux in the sub-mm arc. Correcting for the lensing amplification, we estimate that the unlensed fluxes at 850\\mum\\ are 0.1\\,mJy for the LBG and 0.4\\,mJy for the ERO pair.  The uncertainty in the redshift for the ERO keeps the possibility open that the ERO pair and LBG are at identical redshifts.  This would place them only $\\sim$10 kpc apart at $z\\sim2.9$, and implies that the objects are interacting and their interaction is likely to be the origin of a violent star-burst revealed by the strong sub-mm emission.  To confirm that the LBG is effectively spatially connected with the ERO pair can only be determined by measuring the redshift of the EROs using NIR spectroscopy or via millimeter CO line search; Given the constraints on the redshifts of either candidate, sub-mm CO observations can be used to detect the CO(3-2) transition line that falls within the observation window of several ground-based heterodyne receivers.  To verify these conclusions, and to accurately separate the contributions from each component will require sub-mm/mm interferometry.  The SCUBA flux is bright enough that it should be detectable using the IRAM Plateau de Bure interferometer.  The SMA and ultimately ALMA will also have the power to resolve the components in great detail. With such further study, and using the lensing amplification of nature's telescope, we should be able to provide a full accounting of the source of sub-mm emission in this intriguing system."
    },
    "0404/astro-ph0404159_arXiv.txt": {
        "abstract": "We propose synchrotron absorption in a magnetosheath forming a cocoon around the magnetosphere of pulsar \\B to be the origin of the eclipse phenomena seen in the recently discovered double pulsar system PSRJ07370-3039 \\A \\& \\Bn. The magnetosheath enfolds the magnetosphere of pulsar \\Bn, where the relativistic wind from \\A collides with \\Bn's magentic field. If this model is correct, it predicts the eclipses will clear at frequencies higher than those of the observations reported to date (nominally, above $\\nu \\sim 5$ GHz.) The model also predicts synchrotron emission at the level of a few to 10 $\\mu$Jy, peaking at $\\nu \\sim 2-5$ GHz with possible orbital modulation. We use simplified semi-analytic models to elucidate the structure of the \\B magnetosphere, showing that the \\A wind's dynamic pressure confines \\Bn's magnetic field to within a radius less than 50,000 km from \\Bn, smaller than \\Bn's light cylinder radius, on the ``daytime'' side (the side facing \\An).  Downstream of \\B (``nightime''), \\B forms a magnetotail.  We use particle-in-cell simulations to include the effects of magnetospheric rotation, showing that the magnetosheath has an asymmetric density distribution which may be responsible for the observed eclipse asymmetries. We use simple estimates based upon the magnetic reconnection observed in the simulations to derive a ``propellor'' spindown torque on \\Bn, which is the dominant mode of angular mementum extraction from this star.  Application of this torque to \\Bn's observed spindown yields a polar dipole field  $\\sim 7 \\times 10^{11}$ Gauss (magnetic moment $\\mu_B \\sim 3.5 \\times 10^{29}$ cgs). This torque has a braking index of unity. We show that the model can explain the known eclipses only if the \\A wind's density is at least 4 orders of magnitude greater than is expected from existing popular models of pair creation in pulsars.  We discuss the implications of this result for our general understanding of pulsar physics.  Our proposal was qualitatively outlined in Kaspi {\\it et al.} (2004) and Demorest {\\it et al.} (2004).  Since those papers' appearance, a similar proposal has been made by Lyutikov (2004). ",
        "introduction": "}  Rotation Powered Pulsars (RPPs) lose their rotational energy becuase of electromagnetic torques. While this fact has been known since the earliest days of pulsar research (Gold 1968, Goldreich and Julian 1969, Ostriker and Gunn 1969), and indeed was predicted before pulsars' discovery (Pacini 1967), 1) the physics of the processes through which the extraction works, 2) the physics of how the rotational energy is transmitted to the surounding world, and 3) the physics of how that energy transforms into the observed synchrotron radiation from the nebulae around pulsars have all remained open questions. Answers to all three questions are of significance not only to the understanding of RPPs themselves, but also to the physics of Active Galactic Nuclei and to the workings of Gamma Ray Burst sources, especially if these outflows are driven by large scale Poynting fluxes from systematically magnetized disks (or perhaps magnetars, in the GRB case.)  Modern pulsar theory suggests that a RPP throws off its rotational energy in the form of a relatively dense, magnetized relativistic plasma wind, largely composed of electron-positron pairs with an embedded wound up magnetic field. Particle acceleration in electrostatic ``gaps'' (polar cap gaps, outer gaps or slot gaps) is thought to be the origin of the $e^\\pm$ plasma, through emission and conversion of gamma rays from accelerated particles within a RPP's magnetosphere ({\\it e.g.} Hibschman and Arons 2001, Harding and Muslimov 1998, Muslimov and Harding 2003, Hirotani {\\it et al.} 2003.) The outflow densities suggested by these models justify the use of relativistic MHD in modeling the winds ({\\it e.g.} Beskin, Kuznetsova and Rafikov 1998, Bogovalov 1999, Contopoulos, Kazanas and Fendt 1999, Vlahakis 2004.) Theoretical models of MHD winds exhibit negligible radiative emission (by construction), and indeed, there has been no positive observational identification of the winds themselves - observational study of the winds' properties has depended on detection of the winds' consequences.  The winds are like a river flowing on dark nights - invisible until the water strikes a dam, or rocks in the stream, when the glimmer of starlight from the spray thrown by the obstacles allows one to infer the river's presence and properties.  To date, the main useful probe of RPPs' energy flow has been the winds' collisions with the ``dams'' created by interstellar and circumstellar media surrounding RPPs. These collisions create prominent Pulsar Wind Nebulae (PWNe) around the young pulsars with large rates of rotational energy loss $\\dot{E}_R = c \\Phi^2$ ($\\Phi$ is the electric potential drop across the magnetopsheric open field lines.) The radiative emissions from these nebulae allow inferences of the plasma content and magnetization at the winds' termination working surfaces (shock waves, in most interpretations.) See Arons (1998, 2002, 2004), Slane (2002), Chevalier (2002, 2004), Reynolds (2003), Komissarov and Lyubarsky (2003), Spitkovsky and Arons (2004), Del Zanna {\\it et al.} (2004) for recent reviews and results on this class of interactions.  Rocks in the relativistic stream provide another window into relativistic wind behavior. Examples of such interactions are the collision between the wind and the ``excretion'' disk around the Be star in PSR 1263-59 (e.g., Kaspi {\\it et al.} 1995, Johnston {\\it et al.} 1996, 2001, Tavani and Arons 1997), and the collision of the wind from the millisecond pulsar PSR 1957+20 with the non-relativistic wind from its companion star (Fruchter {\\it et al.} 1988, Ruderman {\\it et al.} 1989, Arons and Tavani 1993.) As with the PWNe, most of what has been gleaned about the wind properties has come from interpretations of the X-ray detections.  The recent discovery of the double pulsar system PSR J07370-3039 \\A \\& \\B (Burgay {\\it et al.} 2003, Lyne {\\it et al.} 2004) offers a new window into studying a relativistic wind, in this case through the tools of radio astronomy. The binary has an orbital period  $P_b$ of 2.4 hours, an orbital eccentricity $e = 0.08$ and an inclination angle $87 \\pm 3$ degrees. Pulsar  \\A has a spin period $P_{A}$ of 22.7 ms, a rotational energy loss rate $\\dot{E}_{A}$ of  $0.6 \\times 10^{34} \\; {\\rm erg~s}^{-1}$, and a light cylinder radius $R_{\\rm LA}  =cP_A /2\\pi = $1,084 km, which is small compared to the orbit semi-major axis $a = (4.25 \\mp 0.05) \\times 10^{5}$ km. Pulsar \\B has a pulse period $P_B$ of 2.77 s, spin down rate $\\dot{P}_B$ of  $0.8 \\times 10^{-15}$ s s$^{-1}$ which leads to a rotational energy loss rate $\\dot{E}_B$ of  $2 \\times 10^{30}$ ergs s$^{-1}$, and a light cylinder radius $R_{\\rm LB}= 1.32 \\times 10^{5}$ km. Pulsar \\A shows a brief eclipse that lasts approximately 30 seconds when \\A passes behind \\B along the line of sight, at {\\it inferior conjunction}\\footnote{This is standard terminology assuming that \\A is the {\\it primary} star and \\B the {\\it secondary}. Likely binary evolution scenarios suggest that this is the appropriate nomenclature for the two evolved remnants of the original main sequence system. {\\it Superior conjunction} occurs a half orbit later when \\B passes behind \\A along the sight line.}. This eclipse shows substantially slower ingress (7 s) than egress (4 s), and the full eclipse profile is nearly achromatic (Kaspi {\\it et al.} 2004). The flux density and emission profile of pulsar \\B vary around the orbit in nearly achromatic manner over the range 430 MHz to 3.2 GHz (Lyne {\\it et al.} 2004, Demorest {\\it et al.} 2004, Ramachandran {\\it et al.} 2004). The strongest emission episodes of \\B are during  two orbital longitude ranges about $70^\\circ$ apart and asymmetrically spaced by $\\sim30^\\circ$ with respect to inferior conjunction. Two weak \\B emission episodes are  located $\\sim 115^\\circ$ before inferior conjunction  (lasting $\\sim40^\\circ$) and $90^\\circ$ after inferior conjunction  (lasting $\\sim60^\\circ$).  Pulsar \\B is not detected, or perhaps is seen with pulsed flux at the level of $ \\sim 0.4 $\\% of its maximum flux, during a range of orbital longitude that starts $\\sim 60^\\circ$ before superior conjunction, and ends $\\sim 30^\\circ$ after this epoch - effectively, this episode is an eclipse of \\Bn.  These eclipses and emission episodes of \\B offer an opportunity to probe the wind around \\A much closer to the energizing pulsar than has been possible using PWNe in higher voltage systems.  Furthermore, the radio observations are sensitive to low energy relativistic electrons and positrons, providing a look into the instantaneous state of this component of a relativistic wind's plasma - PWNe observations only constrain their winds' low energy particle content averaged over the lives of the nebulae.  We propose that \\Bn's magnetosphere has a structure more similar to that of the Earth's magnetosphere than to the magnetospheres of pulsars not interacting with a companion. In contrast, \\An's magnetospheric properties are decoupled from the binary. Thus, the collision of \\An's wind with \\Bn's magnetosphere causes the formation of a bow shock. The pressure of the post-shock particles and fields confines the \\B magnetopshere on the side that instantaneously faces \\An, with a magnetotail extending behind \\Bn.  Magnetic reconnection allows the shocked wind to create a tangential stress on \\Bn's magnetopshere, which creates the dominant spin-down torque on \\B (a variant of the propellor effect). There is also a less significant relativistic wind component to the torque on \\Bn, created by wind from \\B flowing out the magnetotail, whose transverse size is comparable to  $R_{\\rm LB}$. We speculate that the propellor torque also includes components that align \\Bn's rotation axis with the orbital angular momentum. If so, \\B must be an orthogonal rotator, with its magnetic axis perpendicular to its spin axis.  Polarization observations obtained by Demorest {\\it et al.} (2004) show that \\A is almost an aligned rotator (angle bewteen \\An's magnetic and rotation axes $\\sim 5^\\circ$), with its spin axis substantially misaligned with the orbital angular momentum, $\\angle ({\\boldsymbol \\Omega}_A , {\\boldsymbol \\Omega}_{\\rm orbit}) \\sim 50^\\circ $. Then the equatorial current sheet in \\An's wind is likely to have thickness $ \\sim 10^\\circ $ around the wind's equator, with the result that for $340^\\circ$ of orbital phase, \\B is immersed in the high latitude, possibly slow and dense $e^\\pm$ wind. We suggest that latitudinal variation in the confining pressure exerted by the wind causes variation of \\Bn's beaming morphology, which may be in part responsible for the orbit dependent variations of \\Bn's pulse morphology.  The bow shock creates a magnetosheath of relativistically hot, magnetized plasma  which enfolds the confined \\B magnetosphere.  We show that synchrotron absorption in the magnetosheath can explain the eclipse of \\A at inferior conjunction and of \\B at superior conjunction. The model, which requires surprisingly high density in the \\A wind, predicts eclipse clearing at frequencies higher than observed to date - nominally, above 5 GHz - and also predicts observable, orbitally modulated synchrotron emission, at the level of 10 $\\mu$Jy at $\\nu \\sim 5$ GHz.  In the context of this model the eclipses of  \\A and \\B and the other emission phenomenology of  \\B provide the first significant constraint on the properties of a relativistic wind near its source  - \\B forms a magnetospheric rock in the relativistic stream from \\A only 785 light cylinder radii outside the fast pulsar's magnetosphere. The model suggests that $\\sigma$, the ratio of Poynting flux to kinetic energy flux, in \\An's wind  just upstream of the bow shock is certainly less than 2.5 and probably $\\sim 0.2$, much less than what is expected from existing ideal MHD theories of relativistic wind outflow. Thus, the interaction of \\An's wind with \\B suggests magnetic dissipation in the wind begins quite close to the source. ",
        "conclusions": "Our results show that if synchrotron absorption in the magnetosheath is the cause of eclipse phenomena in this fascinating system, the wind from \\A at latitudes outside the equatorial current sheet ($|\\lambda | > 5^\\circ $) is dense ($\\kappa \\sim 10^6$), slow ($\\gamma_{wind} \\sim 75$) and weakly magnetized at $r \\approx 850,000 $ km from \\An. The magnetosheath synchrotron absorption model predicts the eclipses will clear at higher frequencies (nominally, $\\nu > 5$ GHz), and that synchrotron emission (probably with some orbital modulation) will be detectable at the level of 5-10 $\\mu$Jy at $\\nu \\sim 5$ GHz.  If this is the correct interpretation of the eclipse phenomena, the eclipses are the first (semi-)direct detection of a RPP's wind outside of the equatorial current sheet. Essentially all the phenomena in the young PWNe can be ascribed to the equatorial winds (Kennel and Coroniti 1984a,b, Coroniti 1990, Gallant and Arons 1994, Bogovalov and Khangoulian 2002, Lyubarsky 2002, Komissarov and Lyubarsky 2003, 2004, Spitkovsky and Arons 2004, Del Zanna {\\it et al.} 2004) as they interact with the plasma of the surrounding PWNe. The density paramater $\\kappa$ is very high compared to expectations derived from current models of magnetospheric pair creation. These have been reasonably successful in accounting for the plasma fluxes inferred to be in the equatorial winds, where we see the results of a very high $\\gamma_{wind} $ outflow.  They have not been successful in accounting for the larger populations of lower energy particles which produce the radio synhrotron emission from the young PWNe.  The low energy particle injection rates averaged over the history of these systems are factors of 50,  and more, larger than derived from standard pair creation models (polar cap, outer gap,  slot gap,...).  Our results, and analogous results found by Lyutikov (2004), are even more radical - standard pair creation models applied to \\A yield pair injection rates at least 4 orders of magnitude smaller ($\\kappa < 100$) than are required by the magnetosheath synchrotron absorption model.  One solution to the excess low energy plasma problem in young PWNe has been unusual evolutionary spindown history. If the energizing pulsars had much larger spin rates (or magnetic fields) early in their lives than one derives from a constant braking index, constant magnetic moment model, the pair production rates at earlier times might have been much larger, and the equatorial winds much slower, than they are at present. This is a possible (if perhaps unlikely) resolution of the problem, since radio emitting electrons and positrons in the PWNe live ``forever'', with synchrotron lifetimes much in excess of the PWNe ages.  Such an evolutionary solution cannot explain the plasma overdensity inferred here for \\An's wind - plasma striking \\Bn's magnetosphere emerged from \\A only 3 seconds before it enters the magnetosheath, and flows out of the magnetosheath even more quickly, after absorbing pulsed radiation both from \\A and from \\Bn. The most efficient explanation of the discrepancy is that the standard pair creation theories are inadequate (i.e., {\\it wrong}), as applied to field lines which feed the non-equatorial wind. We point out that where pair creation models have worked reasonably well, they apply to the feeding of the equatorial wind, and to the origin of pulsed gamma rays - both phenomena occur on field lines connected to the stars near the boundary between the closed and open field lines of their magnetospheres.  Following Ruderman and Sutherland's (1975) vacuum gap model of polar cap pair creation, all subsequent theories have assumed an electrostatic gap structure (strictly steady in the co-rotating frame), with pair plasma taking on the role of poisoning the gap accelerator as fast as the density builds up. The spatial rate of such build up varies, depending on gap geometry and dominant gamma ray emission and pair creation processes, but in all cases, the production rates required by radio observations are not achieved.  In the case of pulsar \\A, one can readily show that if one gives up the concept of gap poisoning and simply asssumes that particle acceleration along the magnetopsheric magnetic field continues uninihibited by pair creation, as in Tademaru's (1973) early cascade model, pair outflows from \\A as large as we have infered from the magnetosheath absorption model are possible.  The physics behind such behavior remains to be elucidated.  The PSRJ07370-3039 \\A \\& \\B system clearly has  promise for helping us to unravel the mysteries of relativistic outflows from compact objects, as well as providing fascinating phenomenological food for thought and further study.  The details of the model described here, along with a number of aspects not touched on in this brief report, will be reported elsewhere."
    },
    "0404/gr-qc0404058_arXiv.txt": {
        "abstract": "We calculate the response functions of a freely falling Unruh detector in de Sitter space coupled to scalar fields of different coupling to the curvature, including the minimally coupled massless case. Although the responses differ strongly in the infrared as a consequence of the amplification of superhorizon modes, the energy levels of the detector are thermally populated. ",
        "introduction": "It is expected that an observer, corresponding to an Unruh detector~\\cite{Unruh:1976,BirrellDavies:1984} coupled to a scalar field, when freely falling in de Sitter space, will perceive radiation with a thermal spectrum of the de Sitter temperature $T_H = H/(2\\pi)$~\\cite{GibbonsHawking:1977,BirrellDavies:1984}, where $H$ denotes the Hubble parameter. It is the purpose of this article, to clarify in what sense this result is universal to scalar fields of different couplings to the de Sitter background and how the detector apprehends the differences.  Let us therefore refine what is meant by the observation of thermal radiation: At first order in perturbation theory, the detector response function is proportional to the Fourier transform of the scalar propagator {\\it w.r.t.} the proper time of the detector, and it describes how many particles are absorbed and emitted per unit time. When being in equilibrium with the de Sitter background, the energy levels of the detector are thermally populated according to the temperature $T_H$.  The easiest way to derive this is to note that the propagator for nonminimally coupled scalars has in the imaginary direction of its proper time $t$ the periodicity $t\\rightarrow t+2\\pi i/ T_H$~\\cite{GibbonsHawking:1977,BoussoMaloneyStrominger:2001,SpradlinStromingerVolovich:2001}. As we will point out, this however does not completely characterize the response function of the detector, which describes the number of particles detected per unit time. The rate turns out to depend on the scalar mass and on its coupling to the curvature, as was first shown in Ref.~\\cite{Higuchi:1986}, circumventing the use of the scalar propagator.  The fact that in de Sitter space the invariance of the quantum vacuum becomes manifest when the scalar propagator only depends on the proper time seperation along a geodesic has led to the practice of defining the de Sitter vacua through this quantity~\\cite{ChernikovTagirov:1968,BunchDavies:1978,Mottola:1984,Allen:1985,AllenFolacci:1987}. However, in the case of a massless scalar, which is minimally coupled to the curvature, this leads to a problem since the propagator is infrared divergent~\\cite{Allen:1985}. We argue that, when regulated by a cutoff, this divergence gives rise to a contribution which is irrelevant to the total detector response. In addition, we compare with the response functions of a detector immersed in a thermal bath in Minkowski space and discuss the situation in de Sitter space-times with dimension other than four. ",
        "conclusions": "We found the response functions for different scalar fields to differ strongly in the infrared, where $\\Delta E< H$. Moreover, they do not in general coincide with the response to an equilibrum state in flat space. A disagreement with the thermal response does  not yet imply that the detector does not equilibrate with the de Sitter background. In fact, the energy levels of the detector are thermally populated. Similar deviations from a Minkowski-space thermal response are also known for accelerated detectors~\\cite{BroutMassarParentaniSpindel:1995,GabrielSpindelMassarParentani:1997}. The disagreement of the response functions should be attributed to the fact, that for the conformally and the minimally coupled scalar field the density of modes per frequency is different.  However, for fields which are massive or nonconformally coupled to the metric, the infrared enhancement is a consequence of the amplification of superhorizon modes leading to cosmological density perturbations, an effect which is absent in the conformally coupled massless case. This makes the different fields clearly distinguishable by observables. The fact that physically distinct situations such as thermally populated flat space and the different types of de Sitter-invariant vacua result in the same thermal distribution function for the energy levels of the detector is due to the insensitivity of this quantity to whether there is a mixed (thermal) state or a pure (de Sitter-invariant) state and to how often the detector responds per unit time \\footnote{ An example for the interpretation of the de Sitter invariant states as thermal can be found in Ref.~\\cite{Higuchi:1986}, where scalar field quantization is performed in static coordinates. The mode functions are chosen to vanish beyond the horizon, where the static coordinates exhibit a coordinate singularity. Since the horizon distance is singled out, the mode functions in static coordinates violate spatial homogeneity. An Unruh detector placed at the origin of the static vacuum measures no particles. On the other hand, when the static vacuum is thermally populated, the response function is the same as for the de Sitter-invariant vacuum, and it is given by~(\\ref{Higuchi:4dim}). Note first that a thermally populated state in static coordinate does not correspond to a usual thermal equilibrium state, since spatial homogeneity is broken. In addition, the mode functions in the static and the de Sitter-invariant vacuum have different support. Indeed, the static mode functions vanish beyond the de Sitter horizon, while the de Sitter-invariant mode functions exhibit superhorizon correlations. }.  Finally, we want to point out that de Sitter invariance of the stress-energy tensor of the conformal vacuum implies that the stress-energy tensor must be proportional to the metric, such that $p_{\\rm conf} = -\\rho_{\\rm conf}$~\\cite{Brandenberger:1983}. This is inconsistent with a thermal equation of state and leads us to the investigation, whether the energy density is captured at higher orders in perturbation theory~\\cite{GarbrechtProkopec:2004}, allowing for the pair creation operators of the scalar field Hamiltonian, which are not captured by the first order perturbation expansion used here. While we argued in this paper that the thermal state of the detector does not contain the full information about the quantum field, it is yet remarkable, that at first order in perturbation theory the detector is insensitive to the (regulated) stress-energy tensor, which is clearly nonthermal."
    },
    "0404/gr-qc0404091_arXiv.txt": {
        "abstract": "In case of spacetimes with single horizon, there exist several well-established procedures for relating the surface gravity of the horizon to a thermodynamic temperature. Such procedures, however, cannot be extended in a straightforward manner when a spacetime has multiple horizons. In particular, it is not clear whether there exists a notion of global temperature characterizing the multi-horizon spacetimes. We examine the conditions under which    a global temperature can exist for a spacetime with two horizons using the example of Schwarzschild-De Sitter (SDS) spacetime. We systematically extend different procedures (like the expectation value of stress tensor, response of particle detectors, periodicity in the Euclidean time etc.) for identifying a temperature in the case of spacetimes with single horizon to the SDS spacetime. This analysis is facilitated by using a global coordinate chart which covers the entire SDS manifold. We find that all the procedures lead to a consistent picture characterized by the following features: (a) In general, SDS spacetime behaves like a non-equilibrium system characterized by two temperatures. (b) It is not possible to associate a global temperature with SDS spacetime except when the ratio of the two surface gravities is rational (c) Even when the ratio of the two surface gravities is rational, the thermal nature depends on the coordinate chart used. There exists a global coordinate chart in which there is global equilibrium temperature while there exist other charts in which SDS behaves as though it has two different temperatures. The coordinate dependence of the thermal nature is reminiscent of the flat spacetime in Minkowski and Rindler coordinate charts. The implications are discussed. ",
        "introduction": "It is possible to associate a notion of a thermodynamic temperature with metrics having a single horizon. For example, the general class of spacetimes described by spherically symmetric metrics of the form \\be ds^2 = f(r) dt^2 - [f(r)]^{-1} dr^2 - r^2 [d\\theta^2 + \\sin^2\\theta d\\phi^2] \\label{eq:genmet} \\e with $f(r)$ having a, single, simple zero at $r=r_0$ [i.e., $f(r) \\simeq f'(r_0) (r - r_0)$ near $r=r_0$], have a fairly straightforward thermodynamic interpretation. In fact, it can be shown that the Einstein's equations can be expressed in the form of a thermodynamic relation $T dS=dE - P dV$ for such spacetimes \\cite{paddy-horizons,padmanabhan03c}, with the temperature being determined by the surface gravity of the horizon: \\be \\kappa = \\f{1}{2} |f'(r_0)| \\label{eq:kappa} \\e The most familiar metric amongst the above class is, of course, the Schwarzschild metric with a black hole event horizon, with the temperature being directly related to the mass of the black hole. The De Sitter metric can also be analysed in a similar manner and one  can again identify  a unique  temperature for  the  metric \\cite{gh77}. However,  since (i) De Sitter spacetime is not asymptotically flat and (ii) the De Sitter horizon is observer dependent, certain new difficulties arise in this case. In particular, concept  like ``evaporation'' of the cosmological horizon is  not very obvious unlike  in the case of ordinary black holes (Some of these issues are discussed in e.g., \\cite{paddy-horizons,padmanabhan03c}). But  these  issues have nothing to do with the existence of {\\it multiple} horizons in spacetimes and hence are beyond the scope of this paper.  A completely different class of conceptual and mathematical difficulties arise while dealing with spacetimes having multiple horizons which we shall analyse in this paper. The simplest spacetime with multiple horizons is that of a black hole in a spacetime with a cosmological constant, described by the Schwarzschild-De Sitter (SDS) metric \\cite{complex-path,tadaki,mu91,sds,mkni98}. The metric is characterized by the presence of a black hole event horizon and a cosmological horizon. In recent times, studying such a spacetime has acquired further significance because of the cosmological observations suggesting the existence of a non-zero positive cosmological constant (\\cite{pag++99}; for reviews, see \\cite{lambdareviews}). While the observations can be explained by a wide class of models (see, e.g \\cite{pc-sn}), including those in which the cosmic equation of state can depend on spatial scale (see e.g., \\cite{pc-tachyon}), virtually all these models approach the De Sitter (DS) spacetime at late times and at large scales. Thus any black hole which forms in the real universe with a cosmological constant provides an idealization of a SDS spacetime.   The SDS metric has the same form (in one coordinate chart) as the metric in (\\ref{eq:genmet}) but with a $f(r)$ that has two simple zeros at $r=r_\\pm$ and two surface gravities $\\kappa_\\pm = (1/2) |f'(r_\\pm)| $. Naively, one could associate two different temperatures to the black hole and  cosmological horizon using the two different   surface gravities. Since the surface gravities are (generically) different, the spacetime behaves like a system with two temperatures --- somewhat like a solid with its edges kept at two thermal baths of two temperatures. In that case, there will be no well-defined notion of global temperature associated with the spacetime \\cite{mkni98,medved02}. While this sounds plausible, one must bear in mind that the notion of a temperature in spacetimes with horizon is neither local nor coordinate independent. Hence it is not clear whether one can associate two separate temperatures with the two horizons. On the other hand, there are also some indications that  one can associate  a single, effective temperature  with the SDS spacetime  at least in some special cases and possibly in specific coordinate charts \\cite{tadaki,mkni98,complex-path}. This suggests that both the viewpoints could be correct but in different coordinate charts. It should be stressed that the temperature is not a property of the spacetime geometry in general but depends on the coordinate chart used by a class of observers \\cite{paddy-horizons,padmanabhan03c}. [The best known example is the flat spacetime itself which acquires an observer dependent  horizon and temperature in the Rindler coordinate chart.] In view of these complications, we feel it is worth analysing the temperature of SDS spacetime in some detail, which we attempt to do in this paper.  Our approach will be as follows.  We will simplify mathematical complexities by working with 1+1 spacetimes since it is well known that the issues we are attempting to address exist even in two dimensions. We shall then use the standard procedures which lead to the concept of a temperature in the case of single horizon metrics to the SDS metric. This analysis shows that the naive notion of associating two different temperatures with the two horizons is indeed justifiable at least in some approximate sense. One of the approaches, based on periodicity of Euclidean time indicates that there could exist a notion of global temperature in SDS when the ratio of surface gravities is a rational number. To investigate this issue carefully, we use a global coordinate chart which covers the entire SDS manifold and show that \\emph{in this coordinate chart} there does exist a global notion of temperature for the SDS metric when the ratio of surface gravities is rational. However, even in this case, the spacetime behaves as though it has two different temperatures in certain coordinate charts while it has one equilibrium temperature in another global coordinate chart. This is reminiscent of flat spacetime which exhibits different thermal characteristics in different coordinate chart. The implications of this result are discussed in the last section. ",
        "conclusions": "Our analysis shows that SDS spacetime possesses a dual thermal interpretation in a manner similar to flat spacetime in Minkowski and Rindler coordinate charts, if the surface gravities satisfy the condition $\\kappa_+/\\kappa_- = n_+/n_-$. In this case, there is a global coordinate chart in which the metric can be described as having a temperature $\\beta^{-1}=\\kappa_-/(2 \\pi n_-) = \\kappa_+/(2 \\pi n_+)$. More conventional coordinate charts lead to the interpretation of SDS having two different temperatures. We stress that while this result is quite interesting (and new, as far as the authors are aware of), such coordinate dependence is well known in quantum field theory in curved spacetime. The only difference between the SDS situation and the Minkowski/Rindler situation is that in the latter the global temperature is zero while in SDS the global temperature is non-zero.  The rational numbers are infinitely dense in the space of real numbers and in practical sense the condition on the surface gravities may {\\it not} impose any condition on $M$ and $H$ at all -- however the very fact that demanding equilibrium in a spacetime imposes a restriction on the surface gravities is an interesting point just as a matter of principle. It would be interesting to examine the possibility of whether any semiclassical calculations can lead to such quantization of surface gravities. It has been suggested from two different classes of semiclassical gravity calculations that the areas of the horizons might be quantized [in units of (Planck length)$^2$] (for reviews, see \\cite{rovelli98,padmanabhan03c}). In that case we have $r_+^2/r_-^2 = N_+/N_-$, where $N_{\\pm}$ are (relatively prime) positive integers. The condition on the ratio of surface gravities is then \\be \\f{\\kappa_+}{\\kappa_-} = \\f{N_- + 2 \\sqrt{N_+ N_-}}{N_+ + 2 \\sqrt{N_+ N_-}} \\e The ratio $\\kappa_+/\\kappa_-$ can still be an irrational number, hence the quantization of areas does {\\it not} necessarily imply the existence of thermal equilibrium. It is clear from the above expression that $\\kappa_+/\\kappa_-$ will be rational only when $\\sqrt{N_+/N_-}$ is rational, i.e., $N_+/N_- = N_1^2/N_2^2$, where $N_1$ and $N_2$ form another set of relatively prime integers. Hence, for the existence of a global temperature one requires that $r_+/r_- = N_1/N_2$. This means that the quantization of areas is not sufficient, there must be further restrictions on the horizon radii. {\\it In particular, it is adequate if the radii of horizons are quantized in the units of Planck length which, of course, is consistent with the notion of area quantization.} In such case, the condition on $MH$ will be \\be M H = \\f{N_2}{2} \\f{N_1 N_2 (N_1 + N_2)} {(N_1^2 + N_1 N_2 + N_2^2)^{3/2}} \\label{cond} \\e The implications of this result are under investigation.  The situation is more unclear when the ratio of surface gravities is not a rational number. The global coordinate system, defined in Appendix B, still exists covering the whole manifold but the metric in this coordinate system does not lead to any thermal interpretation. The other, singular coordinate charts, of course, lead to the conventional view of two different temperatures for the SDS. The somewhat disturbing feature in this case is that, the Euclidean metric, obtained by the analytic continuation in $t$ will necessarily have a conical singularity. Hence, a non-singular Euclidean quantum field theory does not exist in this case. It is, however, unclear whether one should {\\it demand} the existence of the non-singular Euclidean field theory while working on a given curved spacetime. After all, an arbitrary, time dependent background spacetime may not even have a Euclidean continuation, let alone a non-singular one.  But if we make such a demand then we obtain certain bizarre conclusions. For example, if the universe has a cosmological constant, then any black hole that forms in it  must have a mass which satisfies the quantization condition in equation (\\ref{cond}).  Finally, we mention that the thermal behaviour of horizons is closely related to the quasi-normal modes (QNM) as pointed out in recent analyses of QNM's using Born approximation \\cite{pc-qnm}. This investigation shows that the QNM's of the SDS spacetime arises essentially from those of the Schwarzschild metric. In the study performed in this paper, however, both the horizons contribute in equal footing. It will be, therefore, interesting to analyse the semiclassical wave modes in the global coordinate system and compare them with the results in the singular coordinate charts. This, and related issues, are under study.   \\appendix"
    },
    "0404/astro-ph0404535_arXiv.txt": {
        "abstract": "We discuss the contribution of proton photoproduction interactions on the  isotropic infrared/optical background to the cosmic neutrino fluxes. This contribution has a strong dependence on the proton injection energy spectrum, and is essential at high redshifts. It is thus closely correlated with the cosmological evolution of the ultra high energy proton sources and of the inrared background itself. These interactions may also contribute to the source fluxes of neutrinos if the proton sources are located in regious of high infrared emission and magnetic fields.  \\vspace{1pc} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404253_arXiv.txt": {
        "abstract": "We study the angular power spectrum estimate in order to search for large scale ani\\-so\\-tro\\-pies in the arrival directions distribution of the highest-energy cosmic rays. We show that this estimate can be performed even in the case of partial sky coverage and validated over the full sky under the assumption that the observed fluctuations are statistically spatial stationary. If this hypothesis - which can be tested directly on the data - is not satisfied, it would prove, of course, that the cosmic ray sky is non isotropic but also that the power spectrum is not an appropriate tool to represent its anisotropies, whatever the sky coverage available. We apply the method to simulations of the Pierre Auger Observatory, reconstructing an input power spectrum with the Southern site only and with both Northern and Southern ones. Finally, we show the improvement that a full-sky observatory brings to test an isotropic distribution, and we discuss the sensitivity of the Pierre Auger Observatory to large scale anisotropies. ",
        "introduction": "The origin of the highest energy cosmic rays is a theoretical challenge of modern astrophysics, and is subject of much experimental efforts.  Above $10^{20}$ eV, the current data are too scarce for one to make any definitive statement about the existence or the lack of the GZK cutoff, as well as a statistically meaningful information about the arrival direction distribution. Whereas the AGASA experiment ~\\cite{agasa} is over since January 2004, a new generation of experiments especially dedicated to the highest energies is emerging, and the first of them is the Pierre Auger Observatory currently under construction~\\cite{auger}. For a recent review on the state of the art of the highest-energy cosmic rays, we refer the reader to \\cite{cronin} for instance.  The distribution of the arrival directions is certainly one of the most crucial observable in order to yield some evidences about the sources of the highest-energy cosmic rays~\\cite{isola,sigl2,sigl}. When trying to point out large scale anisotropies, one is naturally led to work with the angular power spectrum of the arrival direction distribution. The detection of large scale anisotropy could probe certain classes of sources and/or test certain propagation models in presence of magnetic fields to be associated with such large scale celestial patterns. In addition, the evidence for large scale anisotropy around 1~EeV\\footnote{1 EeV$\\equiv 10^{18}$ eV} claimed by the AGASA collaboration~\\cite{anisagasa2,anisagasa} motivates even more this kind of studies.  From the PeV energy range, the flux of cosmic rays is so low that we need ground based experiments with large collecting areas measuring the secondary products of the interaction of the cosmic ray in the upper atmosphere. As any ground based experiment has only at one's disposal a limited field of view in declination distribution, anisotropy analysis are generally done owing to the nearly uniform exposure in right ascension by using a 1-dimensional coordinate system instead of the natural 2-dimensional one over the sphere. This is the case for example of the Rayleigh formalism~\\cite{linsley} which necessarily corrupts the sensitivity to tiny anisotropies. In order to exploit the angular power spectrum analysis methods, it is assumed within the cosmic rays community that a full exposure of the sky is required \\cite{sommers,anchordoqui}. The aim of this paper is to show that this conclusion arises only because of the choice of the spherical harmonic coefficients estimate, and to show that with another choice of estimate, standard anisotropy analysis methods can be used even with a partial and non-uniform coverage of the celestial sphere.  By denoting $\\vec{n}_i$ each cosmic ray arrival direction, the standard estimate of the spherical harmonic coefficients is computed through \\[ a_{\\ell m} = \\frac{1}{\\mathcal{A}}\\sum_{i=1}^N \\frac{Y_{\\ell m}(\\vec{n}_i)}{\\omega(\\vec{n}_i)} \\] where $N$ is the total number of events, $\\omega(\\vec{n})$ is the relative exposure function of the considered experiment, and $\\mathcal{A}$ a normalization constant taken as $\\sum_{i=1}^N 1/\\omega(\\vec{n}_i)$. As well known, the use of $1/\\omega$ allows for decoupling the modes when working with a variable exposure over the whole celestial sphere, but breaks down in case of partial exposure of the sky, because it is no longer possible to perform the full sky integrations that are required to measure the multi-poles of the celestial cosmic ray intensity \\cite{sommers}.  In this paper, we choose to introduce and adapt the quadratic estimator method that is widely used in the Cosmic Microwave Background  analysis (see e.g.~\\cite{hivon}) where effects of a partial exposure can be deconvoluted from the observations in order to recover the true underlying power spectrum. We show that the application of this method allows for the standard anisotropy analysis with an exposure possibly going to zero in some parts of the sky. This point is of major interest for most cosmic rays experiments, as the Southern site of the Auger Observatory for instance.  Another approach to power spectrum estimation is through maximum likelihood (see ~\\cite{bjk,borrill} and~\\cite{cras_jch} for a review) which has the advantage of solving exactly the problem. It however requires an explicit representation of the sky covariance and is computationally very time consuming and numerically hard to achieve on large datasets. The quadratic estimate proposed here avoids this difficulty using a Monte-Carlo simulation (this is why it is often refered to as a ``frequentist approach'').  This paper is organized as follows~: in the next section, we describe our method and compute all the statistical properties of our choice of angular power spectrum estimates. From section 3 and on we apply it to the forthcoming Auger observatory. In section 3, we present the relevant informations about this experiment that we need in the context of angular power spectrum estimation. In section 4, we discuss the constraints that the Auger Observatory can put on isotropic distribution of cosmic rays at ultra-high energy. At last, in section 5, we extend the analysis to the case of a large scale anisotropic distribution. ",
        "conclusions": "We showed that in the general case of a varying and incomplete exposure on the sky, the true power spectrum of the cosmic ray sources distribution can be recovered. This result is not new in itself as it was introduced a few years ago in the framework of CMB data analysis~\\cite{hivon}. Its application to cosmic ray data is however new and might open new possibilities as the general feeling up to now was that no $C_\\ell$ power spectrum can be reconstructed without a complete sky coverage. The power spectrum that our procedure allows to recover is equivalent to the full sky one if the anisotropies in the arrival directions of the cosmic rays are well modeled by a spatially stationary random field on the sphere. If this is not the case, the recovered power spectrum is still valid, but only for the region that was used to determine it. Anyway in the non spatially stationary case, different power spectra are required in different regions of the sky and our approach is still relevant.  Additionally, we have analytically solved the calculation of the bias and of the variance introduced by the finite sampling of the sky in the general case of a varying and eventually incomplete exposure. Using the deconvolution proposed here, any cosmic ray dataset will be usable for anisotropy determination purpose, provided the fact that the arrival directions and coverage map are known within reasonable precision."
    },
    "0404/astro-ph0404586_arXiv.txt": {
        "abstract": "{We present the results of the ongoing {\\em XMM-Newton} Survey of nearby spiral galaxy M31. 17 X-ray sources detected in the survey have bright radio counterparts, and 15 X-ray sources coincide with supernova remnant (SNR) candidates from optical and radio surveys.  15 out of 17 sources with radio counterparts, not SNR candidates, have spectral properties similar to that observed for background radio galaxies/quasars or Crab-like supernova remnants located in M31. The remaining two sources, XMMU J004046.8+405525 and XMMU J004249.1+412407, have soft X-ray spectra, and are associated with spatially resolved H-alpha emission regions, which makes them two new SNR candidates in M31.  The observed absorbed X-ray luminosities of SNR candidates in our sample range from $\\sim 10^{35}$ to $\\sim 5 \\times 10^{36}$ ergs s$^{-1}$, assuming the distance of 760 kpc. Most of the SNR candidates detected in our survey have soft X-ray spectra. The spectra of the brightest sources show presence of emission lines and can be fit by thermal plasma models with $kT \\sim 0.1 - 0.4$ keV. The results of spectral fitting of SNR candidates suggest that most of them should be located in a relatively low density regions.  We show that X-ray color-color diagrams can be useful tool for distinguishing between intrinsically hard background radio sources and Crab-like SNR and thermal SNR in M31 with soft spectra. } ",
        "introduction": "The Andromeda Galaxy (M31), the closest giant spiral galaxy to our own, is a unique object for the study of optical and X-ray astronomy. Its proximity and favorable orientation allow to observe stellar populations over the full extent of the galaxy at a nearly uniform distance, and with less severe effects of line-of-sight contamination from interstellar gas and dust. Due to similarities between the two galaxies, the results from the study of M31 provide an important benchmark for comparison with the results from the study of our own Milky Way Galaxy. M31 was observed extensively in X-rays with {\\em Einstein}, {\\em ROSAT}, {\\em Chandra} and {\\em XMM} missions, detected hundreds of sources, with some bright X-ray sources coincident with radio-emitting sources and SNR candidates from optical surveys. In the extensive {\\em ROSAT}/PSPC survey of M31 (Supper et al. 2001), 16 X-ray emitting SNR were identified. Recently, Kong et al. (2003) identified two new SNRs using the data of {\\em Chandra} and {\\em VLA} observations. Trudolyubov et al. (2004) used {\\em XMM-Newton} observations to study X-ray spectral properties of the 3 bright thermal SNR in the northern disk of M31. ",
        "conclusions": "Nearly 600 X-ray sources were detected in the ongoing {\\em XMM-Newton} survey of M31. 17 of them have bright radio counterparts not identified with SNR candidates, and 15 sources coincide with SNR candidates from radio and optical surveys.  15 out of 17 sources with radio counterparts, not known to be SNR candidates from optical surveys, have spectral properties similar to that observed for background radio galaxies/quasars or Crab-like supernova remnants located in M31. The energy spectra of the brightest objects are well represented by absorbed power law with photon indices between $\\sim 1.6$ and $\\sim 2.2$, and low-energy absorption in excess of the Galactic foreground value in the direction of M31. The remaining two sources, XMMU J004046.8+405525 and XMMU J004249.1+412407, have soft X-ray spectra, and are associated with spatially resolved H-alpha emission regions, which makes them two new SNR candidates in M31.  Most of the SNR candidates detected in our survey have soft X-ray spectra. The spectra of the brightest sources show presence of emission lines and can be fit by thermal plasma models with $kT \\sim 0.1 - 0.4$ keV. The observed absorbed X-ray luminosities of our SNR candidates range from $\\sim 10^{35}$ to $\\sim 5 \\times 10^{36}$ ergs s$^{-1}$, assuming the distance of 760 kpc. The results of spectral fitting suggest that most of the bright SNR candidates in our sample may be located in the low density regions.  We show that X-ray color-color diagrams can be useful tool for distinguishing between intrinsically hard background radio sources and Crab-like SNR and thermal SNR in M31."
    },
    "0404/astro-ph0404065_arXiv.txt": {
        "abstract": "We present an analysis of optical HST/STIS and HST/FOS spectroscopy of 6 blue stragglers found in the globular clusters M~3, NGC~6752 and NGC~6397. These stars are a subsample of a set of $\\sim$50 blue stragglers and stars above the main sequence turn-off in four globular clusters which will be presented in an forthcoming paper. All but the 6 stars presented here can be well fitted with non-LTE model atmospheres. The 6 misfits, on the other hand, possess Balmer jumps which are too large for the effective temperatures implied by their Paschen continua. We find that our data for these stars are consistent with models {\\it only} if we account for {\\it extra} absorption of stellar Balmer photons by an ionized circumstellar disk. Column densities of H~{\\sc i} and Ca~{\\sc ii} are derived as are the the disks' thicknesses. This is the first time that a circumstellar disk is detected around blue stragglers. The presence of magnetically-locked disks attached to the stars has been suggested as a mechanism to lose the large angular momentum imparted by the collision event at the birth of these stars. The disks implied by our study might not be massive enough to constitute such an angular momentum sink, but they could be the leftovers of once larger disks. ",
        "introduction": "\\label{sec:introduction}  In the color-magnitude diagrams (CMD) of open and globular clusters (GCs), some stars appear brighter and bluer than the main sequence turn-off, yet not as bright as stars on the horizontal branch (HB; Sandage 1953; Paresce et al. 1991). Their brightnesses and colors are consistent with main-sequence evolutionary models for stars more massive than the cluster turnoff (e.g., Saffer et~al. 2002). Blue stragglers are likely to form through a mass-acquisition episode such as a binary merger or a stellar collision (for an overview of stellar collisions, see Shara 2002). In the dense environments of GCs such collision mergers should in fact be common (Hills~\\& Day 1976).  If the stragglers originate in stellar collisions, the rotation rate of the product star should be high (Sills et~al. 2002). After an initial phase when the collision product is bloated, a contracting phase should follow. In the absence of angular momentum sinks, the star should spin up to faster than break-up speed and tear itself apart. A magnetically-locked circumstellar disk has been proposed as an efficient angular momentum reducer (Leonard \\& Livio 1995). Sills et~al. (2002) do not find evidence for the formation of a {\\it substantial} circumstellar disk from Smoothed Particle Hydrodynamics (SPH) simulations, including up to $10^6$ SPH particles, with the lightest particles having masses of $3 \\times 10^{-10}$~\\msun.  In this Letter we present the first observational evidence consistent with thin disks around six GC stragglers. Our conclusion follows a careful analysis that excludes other systematic effects as potential explanations of the data. Whether these disks are the remnants of once more massive, dynamically important disks, or have always been thin disks unable to serve as angular momentum sinks, this finding is likely to provide an important window into the formation and evolution of blue stragglers. ",
        "conclusions": "\\label{sec:discussion}  From the Boltzmann equation, the disk electron temperatures must typically be larger than 9000~K in order to populate the first excited level of hydrogen. The derived column densities depend moderatly on the disk temperatures, which were assumed to be the same as the stellar effective temperatures. From the Saha equation, the hydrogen ionization fraction is 10$^{-3}$ - 10$^{-4}$, implying a collisionally-ionized disk. The ratio of the \\ion{Ca}{2} to \\ion{H}{1} column densities is typically 1$\\times$10$^{-4}$, which is consistent with the stellar Ca/H abundance ratio when we account for the hydrogen ionization fraction and assume no dust depletion of calcium.  With a {\\it total} H column density of 10$^{25}$\\,cm$^{-2}$ (Table~1), assuming a disk radius of 0.1~AU (larger disk radii should be rare in the crowded environments of GC centers), and a thickness of 0.03~AU derived from a covering factor of 20\\% and a stellar radius of 2.0~\\rsun , from simple geomentry we obtain a total H density of 1$\\times$10$^{-11}$~g~cm$^{-3}$ and a total mass of 2$\\times$10$^{-8}$~\\msun . An increase of the disk radius increases the disk mass proportionally.  We suggest that {\\it most} stragglers might have circumstellar disks. The disk pushes data points in the color-color diagram down (Fig.~1). Only those data points that shift outside the locus covered by the models will be recognized as stars with disks. For those stars whose disks do not change their colors enough to displace them outside the locus covered by the models, an alternative model (cooler and with lower gravity) will be found that fits the data. A method of detecting these disks might be by UV spectroscopy of low ionization metal line.  The disks implied by our analysis are irrevocably non-massive. Although typical SPH simulations could not be used to predict disks as light as these, the one-million-particle simulation of Sills et al. has particle masses as small as $\\sim$10$^{-10}$~\\msun\\ and might have forecast the presence of disks as light as $\\sim$10$^{-8}$~\\msun. On the other hand, magneto-hydrodynamical transport of angular momentum, not included in the SPH code of Sills et al., might be important. More work is needed both observationally and theoretically to understand the presence and role of disks in blue straggler formation."
    },
    "0404/astro-ph0404315_arXiv.txt": {
        "abstract": "We present the Final Analysis of the European Large Area ISO Survey (ELAIS) 15 $\\mu$m observations, carried out with the ISOCAM instrument on board the Infrared Space Observatory (ISO).  The data reduction method, known as LARI method, is based on a mathematical model of the detector's behaviour and was specifically designed for the detection of faint sources in ISO-CAM/PHOT data. The method is fully interactive and leads to very reliable and complete source lists.  The resulting catalogue includes 1923 sources detected with $S/N > 5$ in the \\mbox{0.5 -- 100 mJy} flux range and over an area of 10.85 \\mbox{deg$^2$} split into four fields, making it the largest non-serendipitous extragalactic source catalogue obtained to date from ISO data.  This paper presents the concepts underlying the data reduction method together with its latest enhancements. The data reduction process, the production and basic properties of the resulting catalogue are then discussed. ",
        "introduction": "The Infrared Astronomical Satellite \\citep[IRAS,][]{Neugebauer_et_al_1984,Soifer_et_al_1987} was extremely successful in characterizing for the first time the global properties of the mid- and far-infrared sky, carrying out an all-sky survey at wavelengths of 12, 25, 60 and 100 $\\mu$m and leading to discoveries such as those of Luminous, Ultraluminous and Hyperluminous Infrared Galaxies (LIRGs, ULIRG and HLIRGs, respectively), a substantial population of evolving starbursts and the detection of large-scale structure in the galaxy distribution \\citep{Saunders_et_al_1991}.  Unfortunately, the IRAS view was typically limited to the very local Universe ($z \\stackrel{<}{_\\sim} 0.2$), thus hampering statistical studies of infrared-luminous galaxies at cosmological redshifts. Only few sources were detected by IRAS at higher redshifts, typically ULIRGs magnified by gravitational lenses, like F10214+4724 \\citep[$z=2.28$,][]{Rowan-Robinson_et_al_1991}. In particular only about 1\\,000 galaxies were detected all over the sky in IRAS 12 \\mbox{$\\mu$m} band. Infrared source counts based on IRAS data \\citep{Rowan-Robinson_et_al_1984,Soifer_et_al_1984} showed some marginally significant excess of faint sources with respect to no evolution models \\citep{Hacking_et_al_1987,Franceschini_et_al_1988,Lonsdale_et_al_1990,Gregorich_et_al_1995,Bertin_et_al_1997}, but not enough statistics and dynamic range in flux to discriminate between evolutionary scenarios were available.  Although conceived as an observatory-type mission, the Infrared Space Observatory \\citep[ISO,][]{Kessler_et_al_1996} was in many ways the natural successor to IRAS, bringing a gain of a factor $\\sim 1000$ in sensitivity and $\\sim 10$ in angular resolution in the mid-infrared. A substantial amount of ISO observing time was therefore devoted to field surveys aimed at detecting faint infrared galaxies down to cosmological distances. Such surveys were conceived as complementary in flux depth and areal coverage, allowing a systematic investigation of the extragalactic sky down to so far unattainable flux densities at both mid and far infrared wavelengths, whose results are summarized by \\citet{Genzel_and_Cesarsky_2000}. In particular, extragalactic 15 \\mbox{$\\mu$m} source counts determined with ISOCAM \\citep{Elbaz_et_al_1999,Gruppioni_et_al_2002} have revealed a significant departure from Euclidean slope within the 1 - 5 mJy flux range, which has been interpreted as evidence for a strongly evolving population of starburst galaxies.  The European Large Area ISO Survey \\citep[ELAIS,][]{Oliver_et_al_2000,Rowan-Robinson_et_al_2003} was the most ambitious non-serendipitous survey and the largest Open Time project carried out with ISO, aimed at bridging the flux gap between IRAS all-sky survey and ISO deeper surveys. ELAIS observations mapped areas of about 12 \\mbox{deg$^2$} at 15 and 90 $\\mu$m and smaller areas at 7 and 175 $\\mu$m with the ISOCAM \\citep[7 and 15 $\\mu$m]{Cesarsky_et_al_1996} and ISOPHOT \\citep[90 and 175 $\\mu$m]{Lemke_et_al_1996} cameras. Most importantly, ELAIS 15 $\\mu$m observations are the only ones allowing to sample the 1 - 5 mJy flux range, where most of the source evolution appears to take place.  Since the project approval, the ELAIS consortium, grown in time to a total of 76 collaborators from 30 European institutes, has undertaken an extensive program of ground-based optical and near-infrared imaging and spectroscopy. Thanks to such an extensive multi-wavelength coverage, the ELAIS fields have now become among the best studied sky areas of their size, and natural targets of on-going or planned large-area surveys with the most powerful ground- and space-based facilities. Further details on ELAIS multi-wavelength observations and catalogues are presented in \\citet{Rowan-Robinson_et_al_2003}. After the loss of the WIRE satellite, notwithstanding the observations at several infrared wavelengths soon to come from Spitzer and later from SOFIA and Herschel, ISO observations will remain a valuable database for many years to come. In particular, until the advent of JWST, ELAIS 15 \\mbox{$\\mu$m} observations will provide a complementary view on three areas (S1, N1 and N2) which will be covered at different wavelengths as part of the Spitzer Wide-Area Extragalactic Survey \\citep[SWIRE,][]{Lonsdale_et_al_2003}. Thus the need of reducing such data with the uttermost care and provide the community with an agreed-upon legacy from the ELAIS project.  This paper presents the Final Analysis of ELAIS 15 \\mbox{$\\mu$m} observations, and is structured as follows. In Section~\\ref{elais15.sec} a brief description of the most relevant aspects of ELAIS 15 $\\mu$m dataset is given. Section~\\ref{datared.sec} describes the data reduction method and its improvements. In Section~\\ref{autosim.sec} the technique employed for flux determination and its results are presented. Section~\\ref{simulations.sec} details the results of the simulations that were carried out in order to assess the performance of the data reduction method and thus the quality of the resulting catalogue. In Sections~\\ref{astroacc.sec} and \\ref{photoacc.sec}, respectively, estimates of the achieved astrometric and photometric accuracy are given. Section~\\ref{optids.sec} summarizes the identification of 15 $\\mu$m sources in optical and near-infrared images, while Section~\\ref{photocal.sec} describes the procedure adopted to establish the catalogue photometric calibration. Finally, Section~\\ref{catalogue.sec} describes gives a basic description of the catalogue contents. ",
        "conclusions": "A technique for ISO-CAM/PHOT data reduction, the LARI method, was variously refined and applied to ELAIS 15 $\\mu$m observations. The mathematical model for the detector's behaviour is the same as originally presented in \\citet{Lari_et_al_2001}, but thanks to various improvements, and particularly to a new Graphical User Interface, the method is now more robust and, most importantly, quicker and easier to apply to large datasets. Its application, in the new form, to the four fields composing the dataset (including a re-reduction of S1 observations already presented in \\citet{Lari_et_al_2001}) has produced a catalogue of 1923 sources spanning the 0.5 -- 100 mJy range, detected with a $S/N$ greater than 5 over a total area of $10.85~\\mathrm{deg}^2$. Optical identification of 15 $\\mu$m sources has been carried out on heterogeneous optical and near-infrared imaging material, allowing to determine a robust association for about 85\\,\\% of the sources and identify 22\\,\\% of them as bona fide stars, furtherly demonstrating the reliability of our data reduction process.  The evaluation of the catalogue's quality has been carried out through both accurate simulations and multi-wavelength identification. The astrometric accuracy is of order 1 arcsec in both RA and Dec for $S/N > 10$, while it increases up to about 2 arcsec in both RA and Dec for $S/N \\sim 5$, and somewhat better for sources detected in higher-redundancy sky regions. The photometric accuracy is estimated to be below 25\\,\\% over the whole range of fluxes and redundancy levels probed by our catalogue, and better than 15\\,\\% for $S/N > 10$ sources.  The comparison of measured stellar fluxes with fluxes estimated on the basis of stellar atmospehere models calibrated on IRAS data and on near-infrared photometry allowed to achieve an IRAS/ISO relative photometric calibration. An IRAS/ISO relative calibration factor of $1.0974 \\pm 0.0121$ was determined, shedding doubts on the goodness of the two independently determined calibrations at the 10\\,\\% level. For lack of a simple way to identify error sources in IRAS and/or ISO calibratioon process, it was decided to put our catalogue on the more commonly used IRAS flux scale.  In a forthcoming paper \\citep{Lari_et_al_2004} completeness estimates and extragalactic source counts from this catalogue will be presented, covering the crucial flux range 0.5 -- 100 mJy between ISOCAM 15 $\\mu$m Deep Surveys and IRAS All Sky Survey."
    },
    "0404/astro-ph0404409_arXiv.txt": {
        "abstract": "For the past 10 years there has been an active debate over whether fast shocks play an important role in ionizing emission line regions in Seyfert galaxies.  To investigate this claim, we have studied the Seyfert 2 galaxy Mkn 78, using HST UV/optical images and spectroscopy. Since Mkn 78 provides the archetypal jet-driven bipolar velocity field, if shocks are important anywhere they should be important in this object.  Having mapped the emission line fluxes and velocity field, we first compare the ionization conditions to standard photoionization and shock models.  We find coherent variations of ionization consistent with photoionization model sequences which combine optically thick and thin gas, but are inconsistent with either autoionizing shock models or photoionization models of just optically thick gas.  Furthermore, we find absolutely no link between the ionization of the gas and its kinematic state, while we do find a simple decline of ionization degree with radius. We feel this object provides the strongest case to date against the importance of shock related ionization in Seyferts. ",
        "introduction": "Ionization studies of Seyferts have a long history. Early work led to the establishment of nuclear photoionization as the favored NLR ionizing mechanism. But in the past decade or so, standard models have been called into question because, among other reasons, they strongly underestimate the strengths of many of the weaker high-ionization and high-excitation lines (see \\cite{binette96,robinson00} for a more complete discussion). This led to the development of alternative models, as well as refinements to standard nuclear photoionization. In particular, photoionizing shocks, driven by AGN jets and outflows, have emerged as a viable ionizing source, following work by \\cite{viegas89} and \\cite{dopita96}.  We try to resolve this debate by taking the following approach: we choose a Seyfert with strong, NLR-wide jet-gas interactions. If shocks are important in providing the ionizing power in Seyferts, we should expect to see unambiguous signs of their presence in this object's spectrum. If not, current refinements to nuclear photoionization can be tested.  \\section {Mkn 78 : A Jet-Gas Interaction Archetype}  The Seyfert 2 Mkn 78 was selected as a target because it lies well off the virial correlation for Seyferts (\\cite[Whittle 1992]{whittle92}), indicating the presence of widespread non-gravitational motions in the ionized gas. This makes it one of the best candidates for a strong radio jet/ISM interaction among the sample of nearby Seyferts. \\cite{paper1} discuss the structural aspects of the interaction in detail.  We use a extensive dataset consisting of HST-STIS longslit spectra from four slits sampling all the major emission line features in the NLR at high spatial resolution ($\\sim 0.05$ arcsec). Our spectra give us almost complete FUV and optical wavelength coverage, allowing the measurement of many lines of different ionization state and excitation level. In addition, medium resolution ($\\sim 30$ km/s) spectra allow us to accurately estimate the kinematics of the line emitting gas. ",
        "conclusions": ""
    },
    "0404/astro-ph0404123_arXiv.txt": {
        "abstract": "We present echelle spectrophotometry of the Galactic \\ion{H}{2} region NGC 3576. The data have been taken with the VLT UVES echelle spectrograph in the 3100 to 10400 \\AA\\ range. We have measured the intensities of 458 emission lines, 344 are permitted lines of H$^0$, He$^0$, C$^{+}$, N$^{0}$, N$^{+}$, N$^{++}$, O$^{0}$, O$^{+}$, Ne$^{+}$, S$^{++}$, Si$^{0}$, Si$^{+}$, Ar$^{0}$ and Ar$^{+}$; some of them are produced by recombination and others mainly by fluorescence. Electron temperatures and densities have been determined using different continuum and line intensity ratios. We have derived He$^{+}$, C$^{++}$, O$^{+}$, O$^{++}$ and Ne$^{++}$ ionic abundances from pure recombination lines. We have also derived abundances from collisionally excited lines for a large number of ions of different elements. Remarkably consistent estimations of \\emph{ $t^2$} have been obtained by comparing Balmer and Paschen to [\\ion{O}{3}] temperatures, and O$^{++}$ and Ne$^{++}$ ionic abundances obtained from collisionally excited and recombination lines. The chemical composition of NGC 3576 is compared with those of other Galactic \\ion{H}{2} regions and with the one from the Sun. A first approach to the gas-phase Galactic radial abundance gradient of C as well as of the C/O ratio has been made. ",
        "introduction": "}  NGC 3576 ---also known as Gum38a--- comprises the western part of the RCW 57 complex \\citep{rodg60} and corresponds to a bright knot embedded in a large system of diffuse emission gas filaments \\citep{girardi97}. This knot is one of the most luminous Galactic \\ion{H}{2} regions in the infrared \\citep{figue02}. It was thought that most of the ionization of NGC 3576 is due to two O stars (HD 97319 and HD 97484) and two B stars (HD 974999 and CPD--60$^{\\circ}$2641) which are the main visual components of the OB association \\citep{hum78}; but recent infrared observations suggest that the main ionizing sources of this \\ion{H}{2} region are behind heavily obscuring clouds \\citep{bor97, figue02}. It is located in Carina at a distance of 2.7 kpc \\citep{russ03} and at a Galactocentric distance of 7.4 kpc (assuming a Galactocentric solar distance of 8.0 kpc).  Previous abundance determinations for NGC 3576 are those by \\citet{girardi97}, based on the analysis of collisionally excited lines (hereafter CELs); \\citet{tsa03} based on CELs and some recombination lines (hereafter RLs) of C$^{++}$ and O$^{++}$; and \\citet{simpson95} based on far-infrared data and photoionization models.  The temperature fluctuations problem \\citep{peim67} is, nowadays, a  much-discussed topic in astrophysics of gaseous nebulae \\citep{liu02, liu03, este02, tor03}. Traditionally, the abundance studies for \\ion{H}{2} regions have been based on determinations from CELs, which are strongly dependent on the temperature variations over the observed volume of the nebula. Alternatively, RLs are almost independent of such variations and are, in principle, more precise indicators of the true chemical abundances of the nebula. Several authors have obtained O$^{++}$/H$^+$ from \\ion{O}{2} recombination line intensities for the brightest \\ion{H}{2} regions of the Galaxy (Peimbert, Storey \\& Torres-Peimbert 1993; Esteban et al. 1998, hereafter EPTE; Esteban et al. 1999a, hereafter EPTGR; Esteban et al. 1999b; Tsamis et al. 2003), for extragalactic \\ion{H}{2} regions \\citep{este02b, peim03, tsa03} and for planetary nebulae \\citep{liu00, liu01, ruiz03, peim04}, and all of them have found that the abundance determinations from RLs are systematically larger than those obtained using CELs. The CEL abundances depend strongly on the adopted temperature while the RL abundances are almost independent of it. In the presence of temperature inhomogeneities the temperature derived from the [\\ion{O}{3}] diagnostic lines, $T_e$(\\ion{O}{3}), is considerably higher than the average one and than those temperatures derived from the Balmer and Pashen continua. For \\ion{H}{2} regions the differences in both the abundances derived from CELs and RLs and the temperatures derived from CELs and recombination processes, can be consistently accounted for by assuming a \\te\\ (mean square temperature variation over the observed volume) in the range 0.020--0.044.  \\citet[][ see also Rubin et al. 2003]{odell03} have used a different method to show that there are temperature inhomogeneities in \\ion{H}{2} regions: these authors have determined the columnar temperature along $1.5 \\times 10^6$ lines of sight in the Orion nebula, the distribution of temperatures of their sample supports the $t^2$ values derived from other methods. The origin of temperature fluctuations is still controversial and a serious challenge to our knowledge of the physics and structure of ionized nebulae.  We have taken long-exposure high-spectral-resolution spectra with the VLT UVES echelle spectrograph to obtain accurate measurements of very faint permitted lines of heavy element ions in NGC 3576. We have determined the physical conditions and the chemical abundances of NGC 3576 with high accuracy, including important improvements over previous determinations. We have considered C$^{++}$ and O$^{++}$ abundances obtained from several permitted lines  of \\ion{C}{2} and \\ion{O}{2}, avoiding the problems of line blending, including several 3d-4f transitions which are very useful for abundance determinations because they are free of optical depth effects (Liu et al. 1995 and references therein). We have also derived O$^+$ and Ne$^{++}$ abundances from RLs for the first time in this nebula. We have computed $t^2$ values from the determination of the Balmer and Paschen temperatures, which coincide with the \\te\\ that produces the agreement between the ionic abundances obtained from CELs and RLs. Finally, we have determined helium abundances taking into account a large number of singlet lines of \\ion{He}{1}.  In \\S\\S~\\ref{obsred} and~\\ref{lin} we describe the observations and the data reduction procedure. In \\S~\\ref{phiscond} we obtain temperatures and densities using several diagnostic ratios; also, in this section, we determine \\te\\ from different line intensity ratios and temperature determinations. In \\S~\\ref{helioabund} we briefly analyze the recombination spectra of \\ion{He}{1} and derive the He$^{+}$/H$^{+}$ ratio. In \\S~\\ref{ionic} ionic abundances are determined based on RLs, as well as on CELs. In \\S~\\ref{abuntot} the total abundances are determined. In \\S\\S~\\ref{discus} and ~\\ref{conclu} we present the discussion and the conclusions, respectively. ",
        "conclusions": "}  We present echelle spectroscopy in the 3100-10400 \\AA\\ range for the \\ion{H}{2} region NGC 3576 (Gum38a). We have measured the intensities of 461 emission lines; 170 of them are permitted lines of heavy elements. This is the most complete list of emission lines obtained for this object and one of the largest collections ever taken for a Galactic \\ion{H}{2} region.  We have derived physical conditions of the nebula making use of many different line intensities and continuum ratios. The chemical abundances have been derived for a large number of ions and different elements. We find excellent agreement between the C$^{++}$/H$^+$ ratio obtained from the brightest \\ion{C}{2} RL, $\\lambda$4267 \\AA\\ and others corresponding to 3d-4f transitions of this ion. All these transitions are ---in principle--- excited by pure recombination and give a precise determination of the C$^{++}$ abundance. We find also a good agreement between the O$^{++}$/H$^+$ ratios derived from RLs of multiplets 1, 4, 10, 20 and 3d-4f, which are case-independent transitions and produced largely by recombination. Alternatively, abundances derived for N$^{++}$ for different multiplets show differences as high as a factor of 3. These differences probably are due to fluorescence effects. Finally, we have also determined abundances of O$^+$ and Ne$^{++}$ from RLs for the first time in this object.  We have obtained an average $t^2$=0.038$\\pm$0.009 both by comparing the O$^{++}$ and Ne$^{++}$ ionic abundances derived from CELs to those derived from RLs, and by comparing the electron temperatures determined from ratios of CELs to those obtained from the Balmer and Paschen continua. It is remarkable that the four individual values obtained are almost coincident. The adopted average value of $t^2$ has been used to correct the ionic abundances determined from CELs.  We have estimated the C/H, O/H, and C/O Galactic radial abundance gradients making only use of determinations based on RLs of \\ion{H}{2} regions, obtaining values of -0.090, -0.061, and -0.029, respectively. These estimation is based in four objects covering a rather narrow interval of galactocentric distances (from 6 to 9 kpc).  We would like to thank R. Kisielius and P. J. Storey for providing us with their  latest calculations of effective recombination coefficients for Ne, D.P. Smits for providing us unpublished atomic calculations for He, and L. Carigi for testing our results with her chemical evolution models. We would also like to thank an anonymous referee for his/her valuable comments. This work has been partially funded by the Spanish Ministerio de Ciencia y Tecnolog\\'{\\i}a (MCyT) under project AYA2001-0436. MP received partial support from DGAPA UNAM (grant IN 114601). MTR received partial support from FONDAP(15010003), a Guggenheim Fellowship and Fondecyt(1010404). MR acknowledges support from Mexican CONACYT project J37680-E."
    },
    "0404/astro-ph0404236_arXiv.txt": {
        "abstract": "{  Using the Gaussian+Schechter composite luminosity functions measured from the ESO-Sculptor Survey \\citep{lapparent03aI}, and assuming that these functions do not evolve with redshift out to $z\\sim1$, we obtain evidence for evolution in the late spectral class containing late-type Spiral (Sc+Sd) and dwarf Irregular (dI) galaxies. There are indications that the Sc+Sd galaxies are the evolving population, but we cannot exclude that the dI galaxies also undergo some evolution. This evolution is detected as an increase of the Sc+Sd+Im galaxy density which can be modeled as either $n(z)\\propto1+3(z-0.15)$ or $n(z)\\propto(1+z)^2$ using the currently favored cosmological parameters $\\Omega_m=0.3$ and $\\Omega_\\Lambda=0.7$; the uncertainty in the linear and power-law evolution rates is of order of unity.  For $\\Omega_m=1.0$ and $\\Omega_\\Lambda=0.0$, the linear and power-law evolution rates are $\\sim4\\pm1$ and $\\sim2.5\\pm1$ respectively.  Both models yield a good match to the ESS $BVR_\\mathrm{c}$ redshift distributions to $21-22$\\mag~and to the number-counts to $23-23.5$\\mag, which probe the galaxy distribution to redshifts $z\\sim0.5$ and $z\\sim1.0$ respectively.  \\hspace{0.5cm}  The present analysis shows the usefulness of the joint use of the magnitude and redshift distributions for studying galaxy evolution. It also illustrates how Gaussian+Schechter composite luminosity functions provide more robust constraints on the evolution rate than pure Schechter luminosity functions, thus emphasizing the importance of performing realistic parameterizations of the luminosity functions for studying galaxy evolution.  \\hspace{0.5cm}  The detected density evolution indicates that mergers could play a significant role in the evolution of late-type Spiral and dwarf Irregular galaxies. However, the ESO-Sculptor density increase with redshift could also be caused by a $\\sim1$\\mag~brightening of the Sc+Sd+dI galaxies at $z\\sim0.5$ and a $\\sim1.5-2.0$\\mag~brightening at $z\\sim1$, which is compatible with the expected passive brightening of Sc galaxies at these redshifts. Distinguishing between luminosity and density evolution is a major difficulty as these produce the same effect on the redshift and magnitude distributions.  The detected evolution rate of the ESO-Sculptor Sc+Sd+dI galaxies is nevertheless among the range of measured values from the other existing analyses, whether they provide evidence for density or luminosity evolution. ",
        "introduction": "}  Since the availability of the deep optical numbers counts, the excess at faint magnitudes has provided the major evidence for galaxy evolution at increasing redshifts \\citep{tyson88,lilly93,metcalfe95}. Using models of the spectro-photometric evolution of galaxies \\citep{guiderdoni90,bruzual93}, either passive luminosity evolution or more complex effects have been suggested to explain the faint number-count excess \\citep{guiderdoni91,broadhurst92,metcalfe95}. Although the excess objects where initially envisioned as bright early-type galaxies at high redshift, the lack of a corresponding high redshift tail in the redshift distribution \\citep{lilly93b} consolidated the interpretation in terms of evolution of later type galaxies, namely Spiral and/or Irregular/Peculiar galaxies \\citep{campos97b}.  Here, we report on yet another evidence for evolution of the late-type galaxies, derived from the ESO-Sculptor Survey (ESS hereafter). The ESS provides a nearly complete redshift survey of galaxies at $z\\la0.5$ over a contiguous area of the sky.  A reliable description of galaxy evolution require proper identification of the evolving galaxy populations and detailed knowledge of their luminosity functions.  In this context, the ESS sample has the advantage to be split into 3 galaxy classes which are based on a template-free spectral classification \\citep{galaz98}, and which are dominated by the giant morphological types E+S0+Sa, Sb+Sc, and Sc+Sd+Sm respectively \\citep[][Paper~I hereafter]{lapparent03aI}. In Paper~I, we have performed a detailed measurement of the shape of the luminosity functions (LF hereafter) for the 3 ESS spectral classes. The spectral-type LFs show marked differences among the classes, which are common to the $B$, $V$, $R_\\mathrm{c}$ bands, and thus indicate that they measure physical properties of the underlying galaxy populations.  The analysis of the ESS LFs in Paper~I also provides a revival of the view advocated by \\citet{binggeli88}: a galaxy LF is the weighted sum of the \\emph{intrinsic} LFs for each morphological type contained in the considered galaxy sample; in this picture, differences in LFs mark variations in the galaxy mix rather than variations in the intrinsic LFs \\citep{dressler80,postman84,binggeli90,ferguson91,trentham02a,trentham02b}. Local measures show that giant galaxies (Elliptical, Lenticular, and Spiral) have Gaussian LFs, which are thus bounded at both bright and faint magnitudes, with the Elliptical LF skewed towards faint magnitudes \\citep{sandage85b,jerjen97b}. In contrast, the LF for dwarf Spheroidal galaxies may be ever increasing at faint magnitudes to the limit of the existing surveys \\citep{sandage85b,ferguson91,jerjen00,flint01a,flint01b,conselice02}, whereas the LF for dwarf Irregular galaxies is flatter \\citep{pritchet99} and may even be bounded at faint magnitudes \\citep{ferguson89b,jerjen97b,jerjen00}. In Paper~I, by fitting the ESS spectral-type LFs with composite functions based on the Gaussian and Schechter LFs measured for each morphological type in local galaxy groups and clusters \\citep{sandage85b,jerjen97b}, we confirm the morphological content in giant galaxies of the ESS classes, and we detect an additional contribution from dwarf Spheroidal (dE) and dwarf Irregular galaxies (dI) in the intermediate-type and late-type classes respectively.  We then suggest that by providing a good match to the ESS spectral-type LFs, the local intrinsic LFs may extend to $z\\sim0.5$ with only small variations.  In the following, we report on the measurement of the amplitude of the LFs for the 3 ESS spectral-type LFs, and on the detection and measurement of redshift evolution for the late-type galaxies. \\sct~\\ref{ess} lists the main characteristics of the ESS spectroscopic survey.  \\sct~\\ref{sp_kcor} recalls the definition of the ESS spectral classes and the technique for deriving the corresponding K-corrections and absolute magnitudes.  \\sct~\\ref{shape} shows the measured composite fits of the ESS LFs for the 3 spectral classes. In \\sct~\\ref{phistar}, we describe the various techniques for measuring the amplitude of the LF (\\sct~\\ref{phistar_equat}) and the associated errors (\\sct~\\ref{errors}); we then apply these techniques to the ESS and show the detected evolution in the late-type galaxies (\\sct~\\ref{evol}).  In \\sct~\\ref{counts}, we use the ESS magnitude number-counts to derive improved estimate of the late-type galaxy evolution rate in the $B$, $V$, and $R_\\mathrm{c}$ bands.  We then examine in \\sct~\\ref{nz} the redshift distributions for the 3 spectral classes in the 3 filters, and we verify that the measured evolution rates for the late-type galaxies match the ESS expected redshift distributions. Then, in \\sct~\\ref{evol_comp}, we compare the detected evolution in the ESS LF with those derived from other existing redshift surveys which detect either number density evolution (\\sct~\\ref{dens_comp}) or luminosity evolution (\\sct~\\ref{lum_comp}). Finally, \\sct~\\ref{concl} summarizes the results, discusses them in view of the other analyses which detect evolution in the late-type galaxies, and raises some of the prospects. ",
        "conclusions": "}  Using the Gaussian+Schechter composite LFs measured for the ESO-Sculptor Survey, we obtain evidence for evolution in the late spectral-type population containing late-type Spiral (Sc+Sd) and dwarf Irregular galaxies.  This evolution is detected as an increase of the galaxy density $n(z)$ which can be modeled as $n(z)\\propto1+P_{0.15}(z-0.15)$ with $P_{0.15}\\sim3\\pm1$ or as $n(z)\\propto(1+z)^\\gamma$ with $\\gamma\\sim2\\pm1$ using the currently favored cosmological parameters ($\\Omega_m$,$\\Omega_\\Lambda$)=(0.3,0.7); for ($\\Omega_m$,$\\Omega_\\Lambda$)=(1.0,0.0), $P_{0.15}\\sim4\\pm1$ and $\\gamma\\sim2.5\\pm1$.  Both models yield a good match of the ESS $BVR_\\mathrm{c}$ redshift distributions to $21-22$\\mag~and the number-counts to $23-23.5$\\mag, which probe the galaxy distribution to redshifts $z\\sim0.5$ and $z\\sim1.0$ respectively.  Using \\emph{both} the redshift distributions and the number-counts allows us to lift part of the degeneracies affecting faint galaxy number counts: the redshift distributions allow to isolate the evolving populations, whereas the faint number counts provide better constraints on the evolution rate. These results are based on the hypothesis that the shape of the LF for the ESS late-type class does not evolve with redshift out to $z\\sim1$.  Examination of the other existing redshift surveys to $z\\sim0.5-1.0$ indicates that a wide range of number-density evolution rates have been obtained. The evolution rate of the Sc+Sd+dI galaxies detected in the ESS is among the range of measured values, with some surveys having weaker of higher evolution rates.  The most similar survey to the ESS, the CNOC2, yields a twice larger increase in the number density of late-type Spiral and Irregular galaxies at $z\\sim1$.  A priori, density evolution indicates that mergers could play a significant role in the evolution of late-type Spiral and Irregular galaxies.  \\citet{lefevre00} detect a $\\sim20$\\% increase in the fraction of galaxy mergers from $z\\sim0$ to $z\\sim1$, which can be modeled as $\\propto(1+z)^{3.2}$; interestingly, examination of their \\fg 1 indicates that a significant fraction of the merger galaxies have a Spiral or Irregular structure.  The ESS density increase for the Sc+Sd+dI galaxies could also be caused by a $\\sim1$\\mag~brightening of these galaxy populations at $z\\sim0.5$ and a $\\sim1.5-2.0$\\mag~brightening at $z\\sim1$ (depending on the filter and cosmological parameters). This luminosity evolution is compatible with the expected passive brightening of Sc galaxies at increasing redshifts \\citep{poggianti97}.  \\citet{driver01} also shows that the Hubble Deep Field \\citep{williams96} bi-variate brightness distributions for Elliptical, Spiral, and Irregular galaxies are all consistent with passive luminosity evolution in the 3 redshift bins $0.3-0.6$, $0.6-0.8$, $0.8-1.0$.  The ESS brightening at $z\\sim1$ agrees with the value measured from the CFRS blue galaxies \\citep{lilly95}, but is twice smaller than that measured by \\citet{cohen02} for emission-line dominated galaxies.  In all analyses of the redshift and magnitude distributions, the major difficulty is to distinguish between luminosity and density evolution, as these produce the same net effect on the redshift and magnitude distributions.  Interpretation of density and luminosity evolution of a galaxy population is also complicated by possible variations in the star formation rate with cosmic time: \\citet{lilly98} evaluate an increase in the star formation rate of galaxies with large disks by a factor of $\\sim3$ at $z\\sim0.7$, which shows as an increase of the luminosity density at bluer wavelengths. Using PEGASE \\citep{fioc97}, \\citet{rocca99} also show that the Sa-Sbc galaxies have a star formation rate which varies more rapidly in the interval $0\\la z\\la 1$ than for the E/S0 or Sc-Im galaxies.  The ESS suggests that the Sc+Sd galaxies are an evolving population, but evolution in the dI galaxies cannot be excluded. Whether the Spiral galaxies, or the Irregular galaxies, or both populations contribute significantly to the excess number-counts is still unclear from the various existing analyses: using photometric redshifts, \\citet{liu98} detected a significant amplitude increase and brightening of the $U$ LF of Sbc and bluer galaxies in the redshift interval $0\\la z \\la0.5$, together with an excess population of starburst galaxies at $z\\ga0.3$ which are absent at $z\\la0.3$; galaxy number counts in the near infrared \\citep{martini01,totani01}, which are less sensitive to current star formation, also allow some moderate number evolution of the Spiral galaxies and/or Irregular galaxies (with $\\gamma\\sim1$, see above).  Evolutionary effects are also detected at higher redshifts: \\citet{driver98} show that at $1\\la z \\la3$, the redshift distributions of Sabc galaxies to $I>25$ and of Sd/Irr galaxies to $I>24$ require some number and luminosity evolution \\citep[but see][]{driver01}. This could be consistent with the recent deep optical and infrared observations which favor mild luminosity evolution of the overall galaxy population \\citep{pozzetti03,kashikawa03}. Nevertheless, none of these surveys allow one to isolate one evolving population among the Spiral and Irregular/Peculiar galaxies.  In contrast, some surveys favor Irregular/Peculiar galaxies as a major contributor to the excess count: an excess population of Irregular/Peculiar galaxies was directly identified as the cause for a strong deviation from no-evolution in the $I$ and $K$ number-counts per morphological type \\citep{glazebrook95,huang98}; a population of gently-evolving starbursting dwarves was also invoked to explain these excess objects \\citep{campos97a}. Using morphology of galaxies obtained from Hubble Space Telescope images, \\citet{im99} provide further evidence for a marked increase in the relative abundance of Irregular and Peculiar galaxies which they interpret as starbursting sub-$L^*$ E/S0 and Spiral galaxies. \\citet{totani98} also detect such an excess population at $z\\ga0.5$. Using again Hubble Space Telescope imaging for galaxy morphology, \\citet{brinchmann98} estimate a $\\sim 30$\\% increase in the proportion of galaxies with an irregular morphology at $0.7\\la z\\la0.9$ \\citep[see also][]{abraham96,volonteri00}.  Other surveys identify Spiral galaxies as the evolving population: morphological number counts based on Hubble Space Telescope images indicate that the Spiral counts rise more steeply than the no-evolution model \\citep{abraham96}. By complementing space-based images with ground-based spectroscopic redshifts, \\citet{brinchmann98} detect a $\\sim1$\\mag~brightening of the Spiral galaxies by $z\\sim1$. \\citet{schade96} also measure a $\\sim1.6$\\mag~brightening of the central surface brightness of galaxy disks at $0.5\\la z \\la1.1$.  We emphasize that measuring a reliable evolution rate requires a realistic parameterization of the intrinsic LFs of each galaxy population. We show here that the ESS Gaussian+Schechter composite LFs provide more robust constraints on the evolution rate than pure Schechter LFs, as a small change in the faint-end slope has a large incidence onto the number-counts. In contrast, variations in the faint-end slope of the dE and dI galaxies, the only populations for which the Schechter faint-end is poorly determined have a smaller impact on the adjustment on the number counts.  \\citet{totani01} also show that separating the E and dE galaxies into a Gaussian+Schechter LF does yield a better agreement of the $K$ number-counts at $K\\ga22.5$ than a pure Schechter function for the joint class of E+dE.  By using better measures of the intrinsic LFs for the various galaxy types, one should be able to obtain improved measurements of the galaxy evolution rates, and eventually to use the faint number counts to probe the cosmological parameters. \\citet{koo93} and \\citet{gronwall95} had already suggested that passive luminosity evolution is sufficient to match the $B$ number counts, thus implying that the number counts could be used to probe the curvature of space. \\citet{pozzetti96} and \\citet{metcalfe01} then showed that passive luminosity evolution with a low value of $\\Omega_m$ provide good fits to UV-optical-near-infrared galaxy counts.  The recent adjustments of very deep optical and near-infrared galaxy counts \\citep{totani00,nagashima01,totani01,nagashima02} which probe the galaxy distribution to $z\\sim3$ further confirm that the Einstein-de Sitter cosmology with ($\\Omega_m$,$\\Omega_\\Lambda$)=(1.0,0.0) is excluded at a high confidence level \\citep[see also][]{totani97,he00}; whereas by using the currently favored values of $\\Omega_m=0.3$ and $\\Omega_\\Lambda=0.7$, these fits constrain the evolution rate of galaxies in the hierarchical clustering picture, with only some mild number evolution of Sbc/Sdm galaxies allowed \\citep{totani01}.  Under-going deep redshift surveys to $z\\ga1$ raise new prospects along this line. The present ESS analysis shows the usefulness of using both the magnitude \\emph{and} redshift distributions for studying galaxy evolution.  By obtaining the redshift distributions per galaxy type to $z\\ga1$ over large volumes which average out the large-scale structure, spectroscopic redshift survey such as the VIRMOS \\citep{lefevre03b} and DEEP2 \\citep{davis03} projects should provide improved clues on the evolving galaxy populations at $z\\sim1$ and better constrain the nature of this evolution.  If these surveys confirm that all galaxy types only experience mild luminosity/density evolution, complementing the deep infrared galaxy counts with morphological information might allow one to lift the degeneracy in the faint infrared number-counts and to confirm whether a low matter-density universe is favored. This will however require a detailed knowledge (i) of the luminosity-size relation for distant galaxies of various morphological type (used to model the selection effects caused by the cosmological dimming in surface brightness, see \\citealt{totani00}), and (ii) of the interstellar and intergalactic extinction."
    },
    "0404/gr-qc0404101_arXiv.txt": {
        "abstract": "We give a surface integral derivation of post-1-Newtonian translational equations of motion for a system of arbitrarily structured bodies, including the coupling to all the bodies' mass and current multipole moments.   The derivation requires only that the post-1-Newtonian vacuum field equations are satisfied in weak-field regions between the bodies; the bodies' internal gravity can be arbitrarily strong. In particular black holes are not excluded. The derivation extends previous results due to Damour, Soffel and Xu (DSX) for weakly self-gravitating bodies in which the post-1-Newtonian field equations are satisfied everywhere. The derivation consists of a number of steps:  (i) The definition of each body's current and mass multipole moments and center-of-mass worldline in terms of the behavior of the metric in a weak-field region surrounding the body.  (ii) The definition for each body of a set of gravitoelectric and gravitomagnetic tidal moments that act on that body, again in terms of the behavior of the metric in a weak-field region surrounding the body.  For the special case of weakly self-gravitating bodies, our definitions of these multipole and tidal moments agree with definitions given previously by DSX.  (iii) The derivation of a formula, for any given body, of the second time derivative of its mass dipole moment in terms of its other multipole and tidal moments and their time derivatives.  This formula was obtained previously by DSX for weakly self-gravitating bodies.  (iv) A derivation of the relation between the tidal moments acting on each body and the multipole moments and center-of-mass worldlines of all the other bodies.  A formalism to compute this relation was developed by DSX; we simplify their formalism and compute the relation explicitly.  (v) The deduction from the previous steps of the explicit translational equations of motion, whose form has not been previously derived. ",
        "introduction": "\\subsection{Background and motivation}   For slow motion sources in weak gravitational fields, general relativity can be accurately described in terms of a post-Newtonian approximation scheme.  This approximation scheme is extremely useful in applications and is very well developed.  Reviews of post-Newtonian theory can be found in Refs.\\ \\cite{damour,dsxI,blanchet} and in the textbook by Will \\cite{Will}.  There are several different types of equations that arise in post-Newtonian theory.  First, one has continuum field equations, which are usually specialized to gravity coupled to perfect or imperfect fluids.  These have been derived up to post-2.5-Newtonian order \\cite{chandrasekhar}.  At post-1-Newtonian order, they have been extended beyond general relativity to encompass the class of theories of gravity described by the parameterized post-Newtonian framework \\cite{Will}.   A second type of equation of motion applies to systems consisting of $N$ interacting, extended bodies moving under their mutual gravitational interactions, in the limit where the bodies' sizes are small compared to their mutual separations.  For such systems one has ``point particle'' equations of motion.  Such equations were first derived \\cite{damour} at post-1-Newtonian order by Lorentz and Droste \\cite{ld}, and later independently by Einstein, Infeld and Hoffmann (EIH) \\cite{eih}.  They were also independently derived by Petrova \\cite{petrova} using a method devised by Fock \\cite{fock}.  These equations are usually called the EIH equations. In recent years the advent of gravitational wave astronomy \\cite{detectors,CutlerThorne} has spurred renewed interest in such equations of motion.  For coalescing binary systems, the waveforms of the emitted gravitational waves are expected to carry a great deal of information, and full exploitation of the expected observations will require accurate theoretical models of the waveforms \\cite{CutlerThorne}.  This requirement has prompted the computation of point-particle equations of motion (as well as radiation reaction effects) to higher and higher post-Newtonian orders. Most recently the coalescence waveform's phase has been computed up to post-3.5-Newtonian order \\cite{frenchgw}; see also Refs.\\ \\cite{americangw,japanesegw}. At post-5-Newtonian order and higher, the concept of point-particle equations of motion will break down due to effects related to the finite sizes of the bodies \\cite{Damour83}.  However an argument due to Damour \\cite{Damour83} indicates that the point-particle equations should be well defined at lower orders, up to and including post-4.5-Newtonian.   A third type of equation of motion applies to systems of $N$ interacting bodies whose sizes cannot be neglected.  These equations consist of the point particle equations of motion supplemented by tidal interaction terms.  In principle, if one included tidal interactions to all multipole orders, and in addition coupled the equations of motion to a dynamical description of the internal degrees of freedom in each body, one would obtain a complete description of the system, equivalent to that provided by the continuum equations of motion (up to radiative effects).   For a system of bodies of typical size $\\sim R$, of typical mass $\\sim M$, and with typical separations $\\sim D$, the force $F$ that acts on one of the bodies can be written schematically as\\footnote{These scalings apply to generic bodies; if the bodies are spherically symmetric the scalings are of course altered.} \\begin{eqnarray} F &\\sim& \\frac{M^2}{D^2} \\bigg\\{ 1 + O \\left( \\frac{M}{D} \\right) + O\\left( \\frac{M^2}{D^2} \\right) + \\ldots \\nn \\\\ \\mbox{} && + O\\left[ \\left( \\frac{R}{D} \\right)^l \\right] + O\\left[ \\frac{M}{D} \\left( \\frac{R}{D} \\right)^l \\right] + \\ldots \\bigg\\}. \\label{eq:eom_schematic} \\end{eqnarray} Here we use geometric units with $G = c=1$.  The terms inside the curly brackets are as follows.  On the first line, the 1 is the usual Newtonian force between two point particles and the second and third terms are the post-1-Newtonian and post-2-Newtonian point-particle corrections.  On the second line, the first term is the correction due to Newtonian tidal couplings; the minimum value of $l$ allowed is $l=2$ corresponding to quadrupolar coupling.  The second term describes the post-1-Newtonian tidal couplings.  Here the minimum allowed value $l$ is lower than in the Newtonian case due to gravitomagnetic interactions which have no Newtonian analogs.  This minimum value is $l=1/2$, corresponding to spin-orbit couplings (assuming that the bodies internal velocities are maximal, $v \\sim \\sqrt{M/R}$).  The purpose of this paper is to compute in detail the post-1-Newtonian tidal interaction terms in Eq.\\ (\\ref{eq:eom_schematic}), for all values of $l$, for a system of $N$ bodies.  The explicit form of these terms has not been derived before, although there is a substantial literature on this topic \\cite{dsxI,dsxII,dsxIII,dsxIV,dsxV,skpw,Futamasenew,Kopeikin0,Brumberg1989,Brumberg1991,Klioner,KlionerII}.  There are a number of motivations for this computation. First, as described in Refs.\\ \\cite{dsxI,skpw}, in the area of celestial mechanics future experiments and observations in the solar system will provide very high precision data.  For example, there are current plans to increase the accuracy of lunar laser ranging from the current centimeter level to the millimeter level \\cite{lunar}.  The future astrometric missions SIM (Space Interferometry Mission) and GAIA (Global Astrometric Interferometer for Astrophysics) are expected to measure angles to an accuracy of a few microarcseconds, as compared to the current accuracy of milliarcseconds. In the radio, VLBI (Very Long Baseline Interferometry) observations currently can yield precisions of order 10 microarcseconds \\cite{vlbi}. Also, the proposed future laser astrometric test of relativity (LATOR) mission \\cite{LATOR} would be sensitive to post-2-Newtonian effects, and therefore would likely require detailed modeling of post-1-Newtonian tidal effects.  Second, gravitational wave measurements of coalescing binary compact stars will likely have some ability to detect finite size effects for sufficiently strong signals \\cite{Rasio}.  Although post-Newtonian tidal effects will in many cases be small compared to Newtonian tidal effects, there are some situations where the post-1-Newtonian effects dominate.  An example is the gravitomagnetic resonant excitation of Rossby modes in neutron stars that are spinning at $\\sim 100 \\, {\\rm Hz}$, which could be detectable with LIGO for moderately strong detected inspirals \\cite{Racine1}.   \\subsection{Tidal coupling in post-Newtonian theory}   The textbook treatment of post-1-Newtonian gravity \\cite{Will} is inadequate for the treatment of tidal interactions for several reasons, as explained by Damour et. al. \\cite{dsxI}.  First, the standard treatment uses a single global coordinate system. Although one can write down the continuum equations of motion for a given body in that coordinate system, it is very difficult to separate out the gravitational influences of the other bodies from the self-field of the body, since the fractional distortions of the coordinate system produced by the other bodies can be large even when the tidal distortion of the star is negligible.  The development of approximation schemes such as linear perturbations about an equilibrium state is hindered by the fact that the equilibrium state is not described in the usual way in the global coordinates.  This difficulty has been comprehensively addressed in a series of papers by Brumberg and Kopeikin (BK) \\cite{Kopeikin0,Brumberg1989,Brumberg1991} and by Damour, Soffel and Xu (DSX) \\cite{dsxI,dsxII,dsxIII,dsxIV}. These authors developed a detailed theory of post-1-Newtonian reference frames, in which each body has associated with it a coordinate system naturally adapted to that body. DSX also developed a formalism to compute translational equations of motion including the coupling to all the mass and current multipole moments of each body\\footnote{The BK and DSX formalisms have recently been generalized to the parameterized post-Newtonian framework for scalar-tensor theories of gravity by Kopeikin and Vlasov\\protect{\\cite{Kopeikin}}.}. They applied their formalism to compute equations of motion including spin and quadrupole couplings.  In this paper, we extend the DSX results in two ways.  First, by simplifying their formalism we are able to compute the explicit form of the translational equations of motion, including all the multipole couplings. Second, we give a derivation that is valid for strongly self-gravitating objects as well as weakly self-gravitating objects\\footnote{That is, we show that the dominant fractional errors scale as $O(M^2/D^2)$; global post-Newtonian methods \\protect{\\cite{dsxI,dsxII,dsxIII}} show only that these errors are $O(M^2/R^2)$ or smaller.}.  We need only assume that the post-1-Newtonian field equations are satisfied in a weak field region surrounding each body.  The bodies' internal gravitational fields can be arbitrarily strong; in particular our assumptions do not exclude black holes.  By contrast, DSX assumed the global validity of the post-1-Newtonian continuum field equations, and so their derivation applies only to weakly self-gravitating objects. Our result also generalizes existing derivations of the Newtonian \\cite{Futamasenew} and post-1-Newtonian \\cite{ifa,thornehartle} equations of motion for strongly self-gravitating objects that incorporate only a few low order multipoles. Similar derivations to higher post-Newtonian orders including monopole terms only can be found in Refs.\\ \\cite{damourprl,ifa1}.   One of the key ideas underlying our derivation is that the equations of motion are determined entirely by the local field equations in weak field regions between the bodies.  This was originally pointed out by Weyl and by Einstein and Grommer; see Thorne and Hartle \\cite{thornehartle} and references therein. Each body is surrounded by a vacuum, weak field region called a ``buffer region'' \\cite{thornehartle}, and the quantities entering into the equations of motion are defined in terms of the behavior of the metric in those buffer regions (see Fig.\\ \\ref{fig1}). In particular, our multipole moments are defined in terms of the behavior of the metric in the buffer regions.  Our definition of multipole moments is thus more general than the definition in terms of integrals over sources used by DSX. However, our multipole moments do coincide with those of DSX in the case of weakly self-gravitating bodies.   \\begin{figure} \\begin{center} \\epsfig{file=fig1.eps,width=8.5cm} \\caption{An illustration of our assumptions for a system of $N$ bodies. Each body is surrounded by a strong field region which is excluded from our analysis.  Surrounding these strong field regions are weak field buffer regions.  Each body's center of mass worldline and mass and current multipole moments are defined in terms of the behavior of the metric in that body's buffer region.  We assume that the vacuum post-1-Newtonian field equations are satisfied in all the buffer regions and in the regions of space between the buffer regions.} \\label{fig1} \\end{center} \\end{figure}    Our derivation consists of a number of steps:  (i) The definition of each body's current and mass multipole moments and center-of-mass worldline in terms of the behavior of the metric in that bodies' buffer region.  (ii) The definition for each body of a set of gravitoelectric and gravitomagnetic tidal moments that act on that body, again in terms of the behavior of the metric in that bodies' buffer region. For the special case of weakly self-gravitating bodies, our definitions of these multipole and tidal moments agree with definitions given previously by DSX.  (iii) The derivation of a formula, for any given body, of the second time derivative of its mass dipole moment in terms of its other multipole and tidal moments and their time derivatives.  This formula was obtained previously by DSX for weakly self-gravitating bodies.  (iv) A derivation of the relation between the tidal moments acting on each body and the multipole moments and center-of-mass worldlines of all the other bodies.  A formalism to compute this relation was developed by DSX; we simplify their formalism and compute the relation explicitly.  (v) The deduction from the previous steps of the explicit translational equations of motion, whose form has not been previously derived.    \\subsection{Results for equations of motion}\\label{sec:resultseom}   We next describe our results for the equations of motion.  We label each body by an integer $A$, with $1 \\le A \\le N$.  We use a harmonic coordinate system $(t,x^i)$ that covers all of spacetime except for the strong field regions near each body. The position of body $A$ in this coordinate system is parameterized by a function $x^i = \\cmz^A_i(t)$ called the ``center of mass worldline''.  This function is defined precisely in Sec.\\ \\ref{sec:configvars} below.  It does not correspond to an actual worldline in spacetime; rather it parameterizes the location of the local asymptotic rest frame (see below) attached to the $A$th body.  That is, it is encoded in the behavior of the metric in a weak field region surrounding body $A$ in the same way that the actual center-of-mass worldline of a weakly self-gravitating body would be encoded.   Associated with each body $A$ is a coordinate system $(s_A, y_A^i)$ which is defined only in that body's buffer region, and which is adapted to the body in the sense that it minimizes the coordinate effects of the external gravitational field due to the other bodies as much as possible.  This coordinate system is discussed in detail in Secs.\\ \\ref{sec:bodyadapted} and \\ref{sec:assumptions} below. We will call the corresponding reference frame the ``body frame'' or, following Thorne and Hartle \\cite{thornehartle}, the body's ``local asymptotic rest frame''. The details of the transformation between the body-adapted coordinates $(s_A,y_A^i)$ and the global coordinates $(t,x^i)$ are important for the purpose of deriving the translational equations of motion.  However, for the purpose of using the equations of motion, one only needs to know the following.  First, the time coordinate $s_A$ is a ``proper time'' associated with body $A$.  It corresponds to the proper time that would be measured by an observer in the local asymptotic rest frame of body $A$.  In that local asymptotic rest frame it is related to the global frame time coordinate $t$ by \\begin{equation} s_A = s_A(t), \\end{equation} where the function $s_A(t)$ is determined by a differential equation [Eqs.\\ (\\ref{eq:eomtime}) and (\\ref{eq:eomtimeexplicit}) below]. Second, the leading order relation between the spatial coordinates $y_A^i$ and $x^i$ is just a translation together with a time-dependent rotation [cf.\\ Eqs.\\ (\\ref{coordinatetransformationII}), (\\ref{eq:com3}) and (\\ref{eq:rotUdef}) below]: \\begin{equation}\\label{eq:dragging} x^i = \\cmz^A_i(t) + U^{A\\,j}_i(t) y_A^j. \\end{equation} The rotation matrix $U^{A\\,j}_i(t)$ describes dragging of inertial frames\\footnote{The time derivative of this rotation matrix is actually of post-1-Newtonian order, so to Newtonian order this is a constant matrix.  In the body of the paper we assumed that this constant, Newtonian order rotation matrix is the unit matrix.  The description of our results given here allows this constant matrix to be arbitrary; this slight generalization would be useful to describe systems that evolve for a time long enough that the accumulated rotation due to frame dragging is of order unity.}; a differential equation for its evolution is given below. The body-adapted coordinates $(s_A,y_A^i)$ rotate with respect to distant stars, while the global coordinates $(t,x^i)$ do not.   Each body $A$ has associated with it a unique set of mass multipole moments \\begin{equation} M^A_L(s_A) = M^A_{a_1 \\ldots a_l}(s_A), \\end{equation} for $l=0,1,2 \\ldots$ which are symmetric, tracefree, spatial tensors with $l$ indices, of which mass dipole $M^A_i(s_A)$ vanishes identically. It also has a unique set of current multipole moments \\begin{equation} S^A_L(s_A) = S^A_{a_1 \\ldots a_l}(s_A), \\end{equation} for $l=1,2, \\ldots$.  These quantities are functions of the body's proper time $s_A$.  In the absence of interactions with other bodies the mass monopole $M^A$ and the spin $S^A_i$ are conserved.   We will obtain below coupled equations of motion for the center of mass worldlines of all the bodies.  Appearing in these equations as unknowns will be the mass multipole moments $M^A_L(s_A)$ for $l\\ge2$, and the current multipole moments $S^A_L(s_A)$ for $l\\ge2$.  In order to obtain a closed system of equations, one would need to supplement the equations of this paper with equations determining the evolution of these multipole moments. We discuss further below various circumstances and approximations in which the evolution of the multipoles can be computed.   Next, we define the moments $\\bodyM^A_L(t)$ and $\\bodyS^A_L(t)$ to be the body's mass and current multipole moments, transformed to the non-rotating frame, and expressed as functions of the global time $t$. These moments are given by the equations [cf.\\ Eqs.\\ (\\ref{eq:calMdef}) and (\\ref{eq:calSdef}) below] \\begin{subequations} \\begin{eqnarray}\\nn \\bodyM^A_{a_1 \\ldots a_l}(t) &=& U_{a_1}^{A\\,a'_1}(t) \\ldots U_{a_l}^{A\\,a'_l}(t) M^A_{a'_1 \\ldots a'_l}[s_A(t)], \\,\\,\\,\\,\\,\\,\\, \\\\ \\label{eq:calMdef0} \\\\ \\mbox{} \\bodyS^A_{a_1 \\ldots a_l}(t) &=& U_{a_1}^{A\\,a'_1}(t) \\ldots U_{a_l}^{A\\,a'_l}(t) S^A_{a'_1 \\ldots a'_l}[s_A(t)].\\,\\,\\,\\,\\,\\,\\, \\label{eq:calSdef0} \\end{eqnarray} \\end{subequations}  We can now write down the schematic form of the equations of motion. They can be written as \\begin{subequations} \\begin{eqnarray} \\cmddotz^A_i(t) &=& {\\cal F}^A_i[ \\cmz^B_i, \\cmdotz^B_i, \\bodyM^B_L, {\\dot \\bodyM}^B_L, {\\ddot \\bodyM}^B_L, \\bodyS^B_L, {\\dot \\bodyS}^B_L ] \\nn \\\\ \\label{eq:eomschematic} \\\\ {\\dot \\bodyM}^A(t) &=& {\\cal F}^A[ \\cmz^B_i, \\cmdotz^B_i, \\bodyM^B_L, {\\dot \\bodyM}^B_L] \\label{eq:eomenergy} \\\\ {\\dot \\bodyS}^A_i(t) &=& {\\bar {\\cal F}}^A_i[ \\cmz^B_i, \\bodyM^B_L ] \\label{eq:eomspin} \\\\ \\frac{ds_A}{dt} &=& {\\bar {\\cal F}}^A[ \\cmz^B_i, \\cmdotz^B_i,\\bodyM^B_L],\\label{eq:eomtime} \\end{eqnarray} and \\begin{eqnarray} \\left[{\\dot {\\bf U}}^A \\cdot \\left({\\bf U}^A\\right)^{-1}\\right]_{ij} &=& {\\cal F}^A_{ij}[\\cmz^B_i, \\cmdotz^B_i, \\bodyM^B_L, {\\dot \\bodyM}^B_L, \\bodyS^B_L ]. \\nn \\\\ \\label{eq:Uevol} \\end{eqnarray} \\end{subequations} Here ${\\cal F}^A_i$, ${\\cal F}^A$, ${\\bar {\\cal F}}^A_i$, ${\\bar {\\cal F}}^A$ and ${\\cal F}_{ij}^A$ are functions of their argument whose specific forms are discussed below. In these equations the dependencies on the time derivatives ${\\dot \\bodyM}^B_L$, ${\\ddot \\bodyM}^B_L$ and ${\\dot \\bodyS}^B_L$ only occur for $l\\ge 2$.  Also the mass dipoles $\\bodyM^B_i$ vanish identically. Therefore, if we assume that the moments $M^A_L(s_A)$ and $S^A_L(s_A)$ are known for $l \\ge 2$, Eqs.\\ (\\ref{eq:calMdef0}) -- (\\ref{eq:Uevol}) form a closed set of evolution equations which can be solved to obtain the center of mass worldlines as well as the rotation matrices ${\\bf U}^A$ and time functions $s_A(t)$.   We remark that the three equations (\\ref{eq:eomschematic}) -- (\\ref{eq:eomspin}) by themselves form a closed set of evolution equations for the variables $\\cmz^A_i(t), \\bodyM^A(t)$, and $\\bodyS^A_i(t)$, if we assume that the moments $\\bodyM^A_L(t)$ and $\\bodyS^A_L(t)$ are known for $l \\ge 2$.  However, approximation schemes for computing the multipole moments for $l\\ge2$ usually yield the variables $M^A_L(s_A)$, $S^A_L(s_A)$ rather than the variables $\\bodyM^A_L(t)$, $\\bodyS^A_L(t)$.  This is because the moments $M^A_L(s_A)$ and $S^A_L(s_A)$ are the physical moments that would be measured by an observer in the the local asymptotic rest frame of body $A$. In such cases we must enlarge the set of variables $\\cmz^A_i(t), \\bodyM^A(t), \\bodyS^A_i(t)$ to include the rotation matrices ${\\bf U}^A(t)$ and time functions $s_A(t)$ in order to obtain a closed set of equations.   We also note that it is formally consistent to post-1-Newtonian accuracy to replace Eq.\\ (\\ref{eq:calSdef0}) with the simpler relation $\\bodyS^A_L(t) = S^A_L(t)$. Nevertheless it can be useful in some circumstances to use the more accurate relation (\\ref{eq:calSdef0}), for example for systems which evolve for sufficiently long times that the rotation matrices $U_a^{A\\,a'}$ become significantly different from unity.    We now discuss the functions ${\\cal F}^A_i$, ${\\cal F}^A$, ${\\bar {\\cal F}}^A_i$, ${\\bar {\\cal F}}^A$ and ${\\cal F}_{ij}^A$ that appear in Eqs.\\ (\\ref{eq:eomschematic}) -- (\\ref{eq:Uevol}).  The functional form of ${\\cal F}^A_i$ is one of our key results.  It is given by Eq.\\ (\\ref{fulleom}) below, with the coefficients modified according to the substitutions given in Eq.\\ (\\ref{eq:monopolesub}) and in Appendix \\ref{sec:newcoeffs}.  These modified coefficients are obtained by combining the results of this paper with those of the second paper in this series \\cite{Racine}, which we will call paper II. The functions ${\\cal F}^A$ and ${\\bar {\\cal F}}^A_i$ are standard functions that can be derived from Newtonian stress-energy conservation for weakly self-gravitating bodies, and their explicit functional forms are respectively given in section IV of paper II \\cite{Racine} and in Eq.\\ (\\ref{spinev}) below.  The validity of Eqs.\\ (\\ref{eq:eomenergy}) and (\\ref{eq:eomspin}) for strongly self-gravitating bodies is derived in paper II \\cite{Racine}. The function ${\\bar {\\cal F}}^A$ is given by [cf.\\ Eqs.\\ (\\ref{coordinatetransformationII}), (\\ref{eq:alphacans}), (\\ref{globaltidalnG}), and (\\ref{eq:calTKdef0}) below] \\begin{eqnarray} {\\bar {\\cal F}}^A &=& 1 - \\frac{1}{2} \\cmdotz^A_i \\cmdotz^A_i \\nn \\\\ \\mbox{} && - \\sum_{B\\ne A} \\sum_{k=0}^\\infty \\frac{(2k-1)!!}{k!} \\frac{ \\bodyM^B_K}{r_{BA}^{k+1}} n^{BA}_K. \\label{eq:eomtimeexplicit} \\end{eqnarray} Here $K$ is the multi-index $b_1 \\ldots b_k$, $r_{BA} = | \\cmbmz^B - \\cmbmz^A|$, ${\\bm n}^{BA} = (\\cmbmz^B - \\cmbmz^A)/ r_{BA}$, and $n^{BA}_K = n^{BA}_{b_1} \\ldots n^{BA}_{b_k}$. Finally, the function ${\\cal F}^A_{ij}$ is given by Eq.\\ (\\ref{eq:Uevol1}) below.     We next discuss various approximation schemes in which the equations of motion (\\ref{eq:eomschematic}) -- (\\ref{eq:Uevol}) can be supplemented by methods for obtaining the evolution of the mass and current multipole moments $\\bodyM^A_L(t)$ and $\\bodyS^A_L(t)$ for $l\\ge2$ in order to obtain a complete, closed set of equations.  Some examples of such approximations are as follows.  (i) The simplest case is when the effect of all the $l\\ge2$ multipoles is negligible, and one can set $\\bodyM^A_L = \\bodyS^A_L =0$ for all $l \\ge 2$.  This yields the monopole-spin truncated equations of motion discussed in Refs.\\ \\cite{thornehartle,dsxII}.  (ii) Another simple case is when the evolution of the multipoles of each body is dominated by dynamics internal to that body, and is negligibly influenced by the tidal fields of the other bodies. In this case, one can solve for the evolution of the multipoles $M^A_L(s_A)$, $S^A_L(s_A)$ of each body separately, and then insert those multipoles into the equations of motion (\\ref{eq:calMdef0}) -- (\\ref{eq:Uevol}).  This application will be valid only if the timescale over which the bodies' multipoles evolve is sufficiently long \\cite{thorne}; see Sec.\\ \\ref{sec:domain_of_validity} below for further discussion of this point.  (iii) Another useful case to consider is that of rigid bodies. As noted by Thorne and G\\\"ursel \\cite{thornegursel}, in general relativity a body's rotation can be rigid only if its angular velocity (with respect to its local asymptotic rest frame) is constant.  If the angular velocity is changing, for example due to precession, then the body cannot be rigid due to Lorentz contraction effects. However, to linear order in the body's angular velocity the motion is rigid \\cite{thornegursel}. The analysis of Thorne and G\\\"ursel can be adapted to the present context, if the bodies' rotations are slow enough that they can be idealized as rigid.  In this case, the time dependence of the mass multipole moments $M^A_L(s_A)$ for $l\\ge2$ can be parameterized in terms of a time-dependent rotation matrix ${\\cal U}_a^{A\\, {\\bar a}}(s_A)$: \\begin{eqnarray} \\;\\;\\;\\;\\;\\;\\;\\;\\; M^A_{a_1 \\ldots a_l}(s_A) &=& {\\cal U}_{a_1}^{A\\,{\\bar a}_1}(s_A) \\ldots {\\cal U}_{a_l}^{A\\,{\\bar a}_l}(s_A) M^A_{{\\bar a}_1 \\ldots {\\bar a}_l}. \\nn \\end{eqnarray} Here the moments $M^A_{{\\bar a}_1 \\ldots {\\bar a}_l}$ are constant; these are the moments in the co-rotating frame which rotates with the body. We define the angular velocity $\\Omega^A_a(s_A)$ in the usual way as ${\\dot {\\cal U}}_a^{A\\,{\\bar b}} {\\cal U}_b^{A\\,{\\bar b}} = \\epsilon_{acb} \\Omega^A_c$. Then, the co-rotating frame spin $S^A_{\\bar a} = {\\cal U}_a^{A\\,{\\bar a}} S^A_a$ is related to the co-rotating frame angular velocity $\\Omega^A_{\\bar a} = {\\cal U}_a^{A\\,{\\bar a}} \\Omega^A_a$ via \\cite{thornegursel,rigidbody} $$ S^A_{\\bar a}(s_A) = I^A_{{\\bar a}{\\bar b}} \\Omega^A_{\\bar b}(s_A), $$ where $I^A_{{\\bar a}{\\bar b}}$ is the (constant) moment of inertia tensor\\footnote{Thorne and G\\\"ursel \\protect{\\cite{thornegursel}} have shown that for fully relativistic stars, as for Newtonian stars, the moment of inertia tensor is constant, independent of the angular velocity, up to linear order in the angular velocity.}. Similarly the higher order current multipole moments are given by $$ S^A_{{\\bar a}_1 \\ldots {\\bar a}_l}(s_A) = I^A_{{\\bar a}_1 \\ldots {\\bar a}_l{\\bar b}} \\Omega^A_{\\bar b}(s_A), $$ where $I^A_{{\\bar a}_1 \\ldots {\\bar a}_l{\\bar b}}$ is a higher order generalization of the moment of inertia tensor \\cite{thornegursel}. Combining these relations with the equations of motion (\\ref{eq:calMdef0}) -- (\\ref{eq:Uevol}) yields a closed system of equations which can be solved for the center of mass worldlines, the rotation ${\\cal U}_a^{A\\,{\\bar a}}(s_A)$ of each body with respect to its local asymptotic rest frame $(s_A,y_A^i)$, and the rotation ${\\bf U}^A(t)$ of that local asymptotic rest frame with respect to distant stars.  These equations describe torqued precession of relativistic objects, generalizing the free precession equations of Thorne and G\\\"ursel \\cite{thornegursel}.  (iv) For weakly self-gravitating bodies one can use the formalism developed by DSX \\cite{dsxI,dsxII,dsxIII} to obtain a post-1-Newtonian description of the internal dynamics of each body, for example by using post-1-Newtonian stellar perturbation theory.  Coupling such a description to the equations of motion yields a closed system of equations.  (v) Lastly, for fully relativistic, spherical stars, one can compute the leading order effects of tidal interactions by combining the results of this paper with linear relativistic stellar perturbation theory using matched asymptotic expansions; see, for example, Refs.\\ \\cite{death,damour2,thorne1,flanagan}.  For example, if one is interested only in the mass quadrupoles, and one restricts attention to the dominant, fundamental $l=2$ modes with no radial nodes, then one has a relation of the form $$ M^A_{ij}(s_A) = \\int ds_A^\\prime K(s_A - s_A^\\prime) G^A_{ij}(s_A^\\prime). $$ Here $K(s_A - s_A^\\prime)$ is a Green's function which can be computed from stellar perturbation theory, and $G^A_{ij}$ is the body-frame gravitoelectric tidal moment that acts on body $A$, which is defined in Sec.\\ \\ref{sec:body_frame_moments} below and which can be computed in terms of the worldlines and multipole moments of the other bodies. Combining this relation with the equations of motion (\\ref{eq:eomschematic}) -- (\\ref{eq:Uevol}) again yields a closed system of equations, if one neglects the mass multipoles for $l \\ge 3$ and the current multipoles.     \\subsection{Domain of validity of our results} \\label{sec:domain_of_validity}  As mentioned above, the key assumption which we make in deriving our results is that the post-1-Newtonian vacuum field equations are satisfied in a weak field region between the bodies; see Sec.\\ \\ref{sec:assumptions} below for more details. We are unable to give a derivation of this assumption from first principles. However, in this subsection we discuss various physical effects which can cause our assumption to break down, and we make estimates of the sizes of these effects.  We believe that the assumption should be generally valid aside from the effects discussed in this subsection.   The first type of correction are post-2-Newtonian corrections to the metric in the weak field, vacuum region between the bodies.  These will give rise to fractional corrections of order $M^2/D^2$, where $M$ is a typical mass and $D$ a typical separation of the bodies, cf. the third term on the first line of Eq.\\ (\\ref{eq:eom_schematic}). We can estimate as follows when these corrections will be larger than the tidal coupling terms which we retain. The estimate given in the last term of Eq.\\ (\\ref{eq:eom_schematic}) can be refined by multiplying it by the dimensionless measure $\\varepsilon_l = \\bodyM_L / (M R^l)$ of the $l$th mass multipole.  Demanding that this quantity be larger than the post-2-Newtonian, point particle term in Eq.\\ (\\ref{eq:eom_schematic}) yields the criterion \\begin{equation} D \\alt \\varepsilon_l^{\\frac{1}{l-1}} \\left( \\frac{R}{M} \\right)^{\\frac{1}{l-1}} \\, R. \\end{equation} Thus, the post-2-Newtonian terms will always dominate at sufficiently large separations $D$, but for any given $l\\ge2$ there will be a range of values of $D$ for which the post-1-Newtonian tidal terms dominate, as long as $\\varepsilon_l \\agt M/R$.  In particular this will be true for generic (non-symmetric) bodies for which $\\varepsilon_l \\sim 1$. Similar estimates apply to current multipole couplings for $l\\ge2$. Thus, there is a nonempty regime in which the post-1-Newtonian tidal couplings computed here dominate over post-2-Newtonian, point-particle effects.   Note, however, that this range of values of $D$ gets smaller as the strength of internal gravity $\\sim M/R$ increases.  In the limit of $M \\sim R$ of a black hole, the post-2-Newtonian terms are always comparable to or larger than the post-1-Newtonian tidal terms. Therefore, our results cannot be applied consistently to black holes without including post-2-Newtonian and higher terms in the equations of motion.   Another type of correction, which is also formally of post-2-Newtonian order, is that due to the time dependence of the mass and current multipole moments of the individual bodies \\cite{thorne}.  The post-1-Newtonian solutions [Eqs.\\ (\\ref{basicmodelA}) -- (\\ref{reducedzeta}) below] do not exhibit the correct retarded dependence on these moments, that is, they are functions of $M_L(t)$ and $S_L(t)$ rather than $M_L(t-r)$ and $S_L(t-r)$. If these moments vary on a timescale $\\tau$, then the corresponding fractional corrections to the mass moments $M^A_L$ scale as $D^4/\\tau^4$, and the fractional corrections to the current moments $S^A_L$ scale as $D^2/\\tau^2$.  Demanding that these corrections be smaller than the post-1-Newtonian accuracies of these quantities ($\\sim M/D$ and $\\sim 1$ respectively) yields the criterion $\\tau \\gg D$ for the current moments, and the more stringent criterion \\begin{equation} \\tau \\gg D \\left( \\frac{D}{M} \\right)^{1/4} \\label{eq:criterion} \\end{equation} for the mass moments. Fractional corrections to the post-1-Newtonian tidal interactions will be of order unity for $\\tau \\sim D (D/M)^{1/4}$.  The criterion $\\tau \\gg D$ essentially says that all of the bodies lie in the near zone of the gravitational radiation produced by any one body, and not in the wave zone \\cite{thorne}, and the criterion (\\ref{eq:criterion}) is a somewhat stronger requirement than this.  To illustrate the criterion (\\ref{eq:criterion}) it is useful to consider some examples.  First, if the time evolution of a body's moments is driven by tidal interactions with other bodies, then the timescale is of order $\\tau \\sim D (D/M)^{1/2}$, and the criterion is satisfied.  Second, suppose that we have a 3-body system consisting of two black holes in a tight binary, together with a third body in orbit around the binary.  We can model such a system as a 2 body system using the formalism of this paper, treating the black hole binary as a single body\\footnote{Although our formalism cannot be consistently applied to individual black holes unless supplemented with post-2-Newtonian and higher order point particle terms, our formalism can be applied to black hole binaries treated as a single body.  This is because binaries are less compact than black holes.} whose mass and current multipole moments are evolving with time due to internal dynamics.  Then, early in the gravitational-wave driven inspiral of the black hole binary, the criterion (\\ref{eq:criterion}) will be satisfied and our results for the equation of motion will be valid.  As the black hole binary gets tighter however, eventually the orbital period will become shorter than (\\ref{eq:criterion}) and our results will no longer be applicable.  This second example illustrates that the post-1-Newtonian approximation can sometimes completely break down, even in the supposed weak-field region between the bodies.  During the final coalescence of the black holes the gravitational radiation metric perturbation will become temporarily as large as the Newtonian potential in the region between the binary and the loosely bound companion.   Our results are not applicable to such systems, in which one of the bodies emits a strong burst of gravitational radiation. Further work is required to deduce the form of the translational equations of motion in this type of situation.   There are two other assumptions made in our derivation which slightly restrict the domain of validity of our results. First, we assume that a coordinate system which covers the weak field region between the bodies can be smoothly extended to cover the bodies' interiors (see Sec.\\ \\ref{sec:pnsolns} below).  This assumption essentially restricts the spatial topology of the bodies' interiors, and excludes objects like eternal black holes, wormhole mouths and naked singularities.  It does not exclude realistic, astrophysical black holes for the reason explained in Sec.\\ \\ref{sec:pnsolns} \\cite{thornegw}. Second, in order for a body's multipole moments to be definable, it is necessary that there exist two concentric coordinate spheres surrounding the object, such that the region between the two spheres is vacuum, in a particular coordinate system centered on the body (see Sec.\\ \\ref{sec:pnsolns} below).  This assumption might break down when two bodies get within one or two radii of one another, slightly before they actually touch.    \\subsection{Organization of this paper}  The structure of the paper is as follows.  In Sec. \\ref{sec:cfe} we introduce our notations for the post-1-Newtonian continuum field equations, and following DSX we define a class of gauges (conformally Cartesian harmonic gauges) that we use throughout the paper.  Section \\ref{sec:gaugefreedom} presents a simplified version of the theory of post-1-Newtonian reference systems of Refs.\\ \\cite{dsxI,dsxII,dsxIII,Kopeikin0,Brumberg1989,Brumberg1991,Klioner}. The key result of this section is the explicit parameterization (\\ref{eq:generalcoordtransform}) of the residual gauge freedom within conformally Cartesian harmonic gauge in terms of a number of freely specifiable functions of time and one harmonic function of time and space.  Section \\ref{singlesys} is devoted to the definitions of the mass multipole moments $M_L(t)$, current multipole moments $S_L(t)$, gravitoelectric tidal moments $G_L(t)$ and gravitomagnetic tidal moments $H_L(t)$ associated with a given object and a given conformally Cartesian, harmonic coordinate system.  These definitions are given in Sec.\\ \\ref{sec:pnsolns} in terms of the general solution (\\ref{basicmodelA}) -- (\\ref{reducedzeta}) of the post-1-Newtonian field equations in a vacuum region between two concentric coordinate spheres that surround the object (the object's ``buffer region'').  Section \\ref{sec:momentsgaugetransformation} analyzes how all of these moments transform under the class of allowed gauge transformations discussed in Sec.\\ \\ref{sec:gaugefreedom}. In Sec.\\ \\ref{sec:bodyadapted} we describe gauge specializations that fix the gauge freedom completely and accordingly determine the multipole and tidal moments uniquely.  We call the resulting coordinate system a body-adapted coordinate system.  Section \\ref{sec:moments1} gives a definition of multipole moments and tidal moments associated with a given object, a given worldline and a given coordinate system.  These moments arise only in intermediate steps in the derivations of this paper and not in our final results. Finally, in Sec.\\ \\ref{sec:comparison} we compare the moment definitions used here with other definitions in the literature.  Section \\ref{eom} derives the law of motion for a single body, that is, the relation between the second time derivative of its mass dipole moment and its other multipole and tidal moments and their time derivatives.  The assumptions and result are described in Sec.\\ \\ref{sec:eomoverview}.  A general description of the surface-integral method of derivation which we use is given in Sec.\\ \\ref{sec:method_of_derivation}.  In Sec.\\ \\ref{sec:post2} we give some of the post-2-Newtonian vacuum field equations which are needed for the derivation.  Section \\ref{sec:nod} derives the single body law of motion to Newtonian order, as a warm-up exercise. Finally, the post-1-Newtonian derivation is given in Sec.\\ \\ref{sec:pnderiv}.  This derivation uses an idea due to Thorne and Hartle \\cite{thornehartle} to deduce the value of a complicated surface integral from previous weak-field computations of DSX \\cite{dsxI,dsxII}.  Section \\ref{manybody} lays the foundations for treating a system of $N$ interacting, finite-sized bodies.   Our assumptions are described and discussed in Sec.\\ \\ref{sec:assumptions}.  In Sec.\\ \\ref{sec:body_frame_moments} we define, for each body, a set of body-frame multipole and tidal moments associated with that bodies' adapted coordinate system.  These are the moments that would be measured by an observer in that bodies local asymptotic rest frame.  Section \\ref{sec:configvars} defines the configuration variables that specify the location, orientation etc.\\ within the global coordinate system of the local asymptotic rest frame which is attached to that body.  These variables include the center of mass worldline and also the time functions and rotation matrices discussed in Sec.\\ \\ref{sec:resultseom} above. In Sec.\\ \\ref{sec:globalcoords} we define for each body multipole and tidal moments associated with the global coordinate system.  These quantities appear only in intermediate steps in our computations and not in our final results.  The relation between the global-frame multipole and tidal moments and the body-frame multipole and tidal moments is computed in Sec.\\ \\ref{sec:cbf}.  Section \\ref{sec:mdefs1} defines the modified versions $\\bodyM_L$ and $\\bodyS_L$ of the body-frame multipole moments, discussed in Sec.\\ \\ref{sec:resultseom} above, which are defined with respect to a frame that is non-rotating with respect to distant stars, and which are expressed as functions of the global time coordinate.  These are the moments that appear in the final equations of motion.  Finally, Sec.\\ \\ref{explicit} gives the derivation of the complete, explicit translational equations of motion for the $N$ body system.   \\subsection{Notations and conventions} \\label{sec:notation}  Throughout this paper we use geometric units in which $G=c=1$.  We use the sign conventions of Misner, Thorne and Wheeler \\cite{mtw}; in particular we use the metric signature $(-,+,+,+)$. Greek indices ($\\mu$,$\\nu$ etc...) run from 0 to 3 and denote spacetime indices, while Roman indices ($a$, $b$, $i$,$j$, etc...) run from 1 to 3 and denote spatial indices. The spacetime coordinates will generically be denoted by $(x^0,x^i) = (t,x^i)$.  Spatial indices are raised and lowered using $\\delta_{ij}$, and repeated spatial indices are contracted regardless of whether they are covariant or contravariant indices. We denote by $n^i$ the unit vector $x^i/r$, where $r = |{\\bm{x}}| = \\sqrt{\\delta_{ij} x^i x^j}$.   When dealing with sequences of spatial indices, we use the multi-index notation introduced by Thorne \\cite{thorne} as modified slightly by Damour, Soffel and Xu \\cite{dsxI}. We use $L$ to denote the sequence of $l$ indices $a_1 a_2 \\ldots a_l$, so that for any $l$-index tensor\\footnote{ Here by ``tensor'' we mean an object which transforms as a tensor under the symmetry group SO(3) of the zeroth order spatial metric $\\delta_{ij}$, not a spacetime tensor.} $T$ we have \\begin{equation} T_L \\equiv T_{a_1a_2... \\,a_l} . \\end{equation} If $l=0$, it is understood that $T_L$ is a scalar. If $l<0$ then $T_L\\equiv 0$.  We define $L-1$ to be the sequence of $l-1$ indices $a_1 a_2 \\ldots a_{l-1}$, so that \\begin{equation} T_{L-1} \\equiv T_{a_1a_2... \\,a_{l-1}}. \\end{equation} If $l=0$, then by convention $T_{L-1} \\equiv 0$. We also define $N$ to be the sequence of $n$ indices $a_1 a_2 \\ldots a_n$, and $L-N$ to be the sequence of $l-n$ indices $a_{n+1} a_{n+2} \\ldots a_l$, so that we can write a relation like \\begin{equation} G_{a_1 ... a_l} = S_{a_1... \\, a_n} T_{a_{n+1} ... \\,a_l} \\end{equation} as $G_L = S_N T_{L-N}$, for any tensors $G$, $S$ and $T$. We define $K$, $P$ and $Q$ to be the sequences of spatial indices $b_1 b_2 \\ldots b_k$, $c_1 c_2 \\ldots c_p$, and $d_1 d_2 \\ldots d_q$, respectively. Repeated multi-indices are subject to the Einstein summation convention, as in $S_L T_L$.  We also use the notations \\begin{equation} x^L \\equiv x^{a_1a_2...\\,a_l} \\equiv x^{a_1}x^{a_2}...\\,x^{a_l} \\end{equation} and \\begin{equation} \\partial_L \\equiv \\partial_{a_1a_2...\\,a_l} \\equiv \\partial_{a_1}\\partial_{a_2} ...\\,\\partial_{a_l}. \\end{equation}   We use angular brackets to denote the operation of taking the symmetric trace-free (STF) part of a tensor. Thus for any tensor $T_L$, we define \\begin{equation} T_{<L>} \\equiv \\text{STF}_L(T_L). \\end{equation} where $\\text{STF}_L$ means taking the symmetric trace-free projection on the indices $L$. For example, if $l=2$, we have \\begin{eqnarray}\\nn T_{<L>} &=& T_{<a_1a_2>} \\\\\\nn &=& \\frac{1}{2}\\left(T_{a_1a_2} + T_{a_2a_1}\\right) - \\frac{1}{3}\\delta_{a_1a_2}T_{jj}. \\end{eqnarray} See Appendix \\ref{stf} for the general definition of STF projection, and for a collection of useful relations involving STF tensors.   Throughout this paper, symbols will generally denote functions (as is common in mathematics) rather than physical quantities (as is common in physics).  For example, in Sec.\\ \\ref{manybody} we define a mass multipole moment $M^A_L(s_A)$ which is a function of a time coordinate $s_A$.  In that section we also use a different time coordinate $t$.  Then, the notation $M^A_L(t)$ will always mean $M_L(s_A)$ evaluated at $s_a=t$, rather than $M^A_L[s_a(t)]$.  Finally, for the aid of the reader an index of symbols is provided in Table \\ref{table:symbols}.  \\bigskip ",
        "conclusions": "In this paper, we have given a surface integral derivation of the full post-1-Newtonian DSX laws of motion (\\ref{mainresult}). We have shown that these laws of motion apply to a wide class of strongly self-gravitating objects, provided that the mass and current moments are appropriately defined in terms of the asymptotic weak field metric in the buffer regions around each body. We have given an explicit form for the coupled equations of motion of the bodies' center of mass worldlines including the effects of \\textit{all} the post-Newtonian mass and current multipole couplings. To the best of our knowledge this is the first time these equations of motion have been written out explicitly. The second paper in this series will include a surface integral derivations of the evolution laws (\\protect{\\ref{eq:mainresult0}}) and (\\protect{\\ref{eq:spinresult}}) for the energy (mass monopole) and the spin $S^A_i$ \\cite{Racine}."
    },
    "0404/astro-ph0404146_arXiv.txt": {
        "abstract": "Cosmic acceleration may be the result of unknown physical processes involving either new fields in high energy physics or modifications of gravitation theory. In the latter case, such modifications are usually related to the existence of extra dimensions (which is also required by unification theories), giving rise to the so-called brane cosmology. In this paper we investigate the phenomenon of the acceleration of the Universe in a particular class of brane scenarios in which a large scale modification of gravity arises due to a gravitational \\emph{leakage} into extra dimensions. By using the most recent supernova observations we study the transition (deceleration/acceleration) epoch as well as the constraints imposed on the parameters characterizing the model. We show that these models provide a good description for the current supernova data, which may be indicating that the existence of extra dimensions play an important role not only in fundamental physics but also in cosmology. ",
        "introduction": "The current idea of a \\emph{negative-pressure} dominated universe seems to be inevitable in light of the impressive convergence of the recent observational results (see, e.g., \\cite{revde} for a review). This in turn has led cosmologists to hipothesize on the possible existence of an exotic dark component that not only would explain these experimental data but also would reconcile them with the inflationary flatness prediction ($\\Omega_{\\rm{Total}} = 1$). This extra component, or rather, its gravitational effects is thought of as the first observational piece of evidence for new physics beyond the domain of the standard model of particle physics, a conclusion that has given rise to many speculations on its fundamental origin \\cite{alcaniz}.   Alternatively, another possible route to deal with this dark pressure problem could be a modification in gravity instead of any adjustment to the energy content of the Universe. This idea naturally brings to light another important question at the interface of fundamental physics and cosmology: extra dimensions. As is well known the existence of extra dimensions is required in various theories beyond the standard model of particle physics, especially in theories for unifying gravity and the other fundamental forces, such as superstring or M theories. As suggested in Ref. \\cite{add}, extra dimensions may also provide a possible explanation for the huge difference between the two fundamental energy scales in nature, namely, the electroweak and Planck scales [$M_{Pl}/m_{EW} \\sim 10^{16}$] (see also \\cite{randall}).   In the cosmological context, the role of extra spatial dimensions as source of the dark pressure is translated into the so-called brane world (BW) cosmologies \\cite{braneS}. The general principle behind such models is that our 4-dimensional Universe would be a surface or a brane embedded into a higher dimensional bulk space-time on which gravity and only gravity could propagate. Brane world scenarios has been a topic of much interest recently. In \\cite{ss}, for instance, a class of BW models which admit a wider range of possibilities for the dark pressure than do the usual dark energy scenarios was investigated. An interesting feature of this class of models is that the acceleration of the Universe can be a transient phenomena, which could help reconcile the supernova evidence for an accelerating universe with the requirements of string/M-theory \\cite{fis}.  Another particularly interesting scenario is the one proposed by Dvali {\\it{et al.}} \\cite{dgp}, which we will refer to it as DGP model. It describes a  self-accelerating 5-dimensional BW model with a noncompact, infinite-volume extra dimension in which the whole dynamics of gravity is governed by a competition between a 4-dimensional Ricci scalar term, induced on the brane, and an ordinary 5-dimensional Einstein-Hilbert action. For scales below a crossover radius $r_c$ (where the induced 4-dimensional Ricci scalar dominates), the gravitational force experienced by two punctual sources is the usual 4-dimensional $1/r^{2}$ force whereas for distance scales larger than $r_c$ the gravitational force follows the 5-dimensional $1/r^{3}$ behavior. The theoretical consistency of the model, and in particular of its self-accelerating solution, is still a matter of debate in the current literature (see, e.g., \\cite{luty}). From the observational viewpoint, however, DGP models have been sucessufully tested in many of their predictions, ranging from local gravity to cosmological observations \\cite{deff1,alc1,alc2,lue}.  In this paper we are particularly interested in testing the viability of DGP scenarios in light of the latest supernova (SNe Ia) data, as provided recently by Riess {\\it et al.} \\cite{rnew}. The sample used, which consists of 157 \\emph{gold} events distributed over the redshift interval $0.01 \\lesssim z \\lesssim 1.7$, is the compilation of the best observations made so far by the two supernova search teams plus 16 new events observed by \\emph{HST}. In agreement with other independent analyses, it is shown that these models constitute a good explanation for the current SNe Ia observations and, hence, that the presence of extra dimensions may provide a viable alternative for the dark pressure problem. ",
        "conclusions": "The results of observational cosmology in the past years have opened up an unprecedented opportunity to establish a more solid connection between Fundamental Physics and Cosmology. Surely, the most remarkable finding among these results comes from SNe Ia observations which suggest that the cosmic expansion is undergoing a late time acceleration. Such a phenomenon has been often explained from two different ways, i.e., either by considering the presence of a negative-pressure dark component, the so-called dark energy, or by assuming the existence of extra spatial dimensions, as motivated by recent developments in particle physics. In this paper we have followed the second route. We have studied the transition (deceleration/acceleration) epoch in the context of DGP models and shown that these scenarios based on a large scale modification of gravity are in good agreement with the current SNe Ia data for values of $\\Omega_{\\rm{m}} = 0.33^{+0.10}_{-0.09}$  and $\\Omega_{r_{c}} = 0.24^{+0.05}_{-0.07}$ (95.4\\% c.l.). By assuming a SNe Ia-independent constraint on the matter density parameter, we have found $\\Omega_{\\rm{r_c}} = 0.21$, leading to an estimate of the crossover scale $r_c \\simeq 1.09 H_o^{-1}$. If flatness is imposed, then we find $\\Omega_{\\rm{m}} = 0.23$ ($\\Omega_{\\rm{r_c}} = 0.148$), which corresponds to a acceleration universe with $q_o \\simeq -0.65$ and a total expanding age of $t_o \\simeq 9.5h^{-1}$ Gyr.  In summary, what we have shown is that at least at the level of background tests (like tests involving SNe Ia measurements), DGP models constitute a viable alternative for the dark energy or dark pressure problem. This may be understood as an indication that the existence of extra dimensions play an important role not only in fundamental physics but also in cosmology."
    },
    "0404/astro-ph0404416_arXiv.txt": {
        "abstract": "We report on the methodology and first results from the Deep Lens Survey (DLS) transient search.  We utilize image subtraction on survey data to yield all sources of optical variability down to 24$^{th}$ magnitude.  Images are analyzed immediately after acquisition, at the telescope and in near--real time, to allow for followup in the case of time--critical events.  All classes of transients are posted to the web upon detection.  Our observing strategy allows sensitivity to variability over several decades in timescale.  The DLS is the first survey to classify and report all types of photometric and astrometric variability detected, including solar system objects, variable stars, supernovae, and short timescale phenomena. Three unusual optical transient events were detected, flaring on thousand--second timescales.  All three events were seen in the $B$ passband, suggesting blue color indices for the phenomena.  One event (OT 20020115) is determined to be from a flaring Galactic dwarf star of spectral type dM4.  From the remaining two events, we find an overall rate of $\\rate = 1.4~ {\\rm events}~ {\\rm deg}^{-2}~{\\rm day}^{-1}$ on thousand--second timescales, with a $95\\%$ confidence limit of $\\rate < 4.3$.  One of these events (OT 20010326) originated from a compact precursor in the field of galaxy cluster Abell 1836, and its nature is uncertain.  For the second (OT 20030305) we find strong evidence for an extended extragalactic host.  A dearth of such events in the $R$ passband yields an upper $95\\%$ confidence limit on short timescale astronomical variability between $19.5 < \\mathcal{M}_R < 23.4$ of $\\rate_R < 5.2~ {\\rm events}~ {\\rm deg}^{-2}~{\\rm day}^{-1}$.  We report also on our ensemble of astrometrically variable objects, as well as an example of photometric variability with an undetected precursor. ",
        "introduction": "\\label{sec-intro}  Characterization of the variable optical sky is one of the new observational frontiers in astrophysics, with vast regions of parameter space remaining unexplored.  At the faint flux levels reached by this optical transient search, previous surveys were only able to probe down to timescales of hours.  An increase in observational sensitivity at short timescale and low peak flux holds the promise of detection and characterization of rare, violent events, as well as new astrophysics.  The detection of transient optical emission provides a window into a range of known astrophysical events, from stellar variability and explosions to the mergers of compact stellar remnants.  Known types of catastrophic stellar explosions, such as supernovae and gamma--ray bursts (GRBs), produce prompt optical transients decaying with timescales of hours to months.  Some classes of GRBs result from the explosion of massive stars -- these hypernovae are known to produce bright optical flashes decaying with hour--long timescales due to emission from a reverse shock plowing into ejecta from the explosion. In addition, GRBs produce optical afterglows, decaying on day to week timescales, resulting from jet--like relativistic shocks expanding into a circumstellar medium.  Even more interesting are explosive events yet to be discovered, such as mergers among neutron stars and black holes.  These may have little or no high-energy emission, and hence may be discoverable only at longer wavelengths \\citep{li-pac}. Finally, there is the opportunity to find rare examples of variability, as well as the potential for discovering new, unanticipated phenomena.  One of the primary science goals of this transient search is the rate and distribution of short timescale astronomical variability. However, we catalogue and report all classes of photometric or astrometric variability with equal consideration.  A weakness in many surveys with targeted science is that serendipitous information is discarded as background -- a counterexample is the wealth of stellar variability information being gleaned from microlensing survey data. The DLS survey geometry and observation cadence are not optimal for maximizing the overall number of detected transients (e.g. \\mycite{nemiroff-03}), being driven instead by weak lensing science requirements.  However, given the current lack of constraints at short timescales and to significant depth, the potential remains for the discovery of new types of astronomical variability.  The search for variability without population bias is one of the primary goals of the Large-aperture Synoptic Survey Telescope (LSST, \\mycite{tyson-lsst02}).  Efforts such as the DLS transient search are a useful and necessary testing ground for the software design and implementation needed to reduce, in real--time, LSST's expected data flow of 20 TB nightly. ",
        "conclusions": "We have reported on the structure of and first results from our wide--field image subtraction pipeline.  An important characterization of our transient survey is the exposure $\\exposure$ at a given timescale and to a given depth.  The DLS transient search is primarily sensitive to $\\sim 1000s$ variability from $19^{th}$ to deeper than $23^{rd}$ magnitudes in $B$, $V$, and $R$.  Within this envelope of sensitivity, we have detected three short timescale optical transient events.  OT 20010326 occurred in the field of galaxy cluster Abell 1836.  One archival HST image of the cluster includes this region, and the transient precursor is present and unresolved.  This indicates a compact precursor, stellar in nature if it resides in our Galaxy.  The colors of this object are consistent at the $1 \\sigma$ level with those expected of Galactic dwarf stars (Figure~\\ref{fig-cmds}), whose flaring activity presents a known background.  However, the object would be too far out of the Galactic plane to belong to the disk population of dwarf stars, and halo subdwarfs are not known to flare. It is also possible the precursor resides in, or behind, the Abell cluster, but overall its nature remains uncertain.  OT 20020115 is identified as a Galactic M dwarf of spectral type dM4, exhibiting classical flare star activity.  Finally, the host for OT 20030305 appears consistently elliptical in the $R$ and $I$ passbands, and inconsistent with the $R$--band stellar PSF, which we consider a strong argument for a resolved host, extragalactic in nature.  We find OT 20030305 the strongest candidate so far for optically detected, short timescale cosmological variability.  The precursor or host objects for OTs 20010326 and 20030305 are definitively redder than galaxies that are known to host GRBs and XRFs.  If our events are cosmological in nature, this suggests that optical and high energy events arise from different mechanisms, or, given the dearth of GRBs from dusty star--forming galaxies, a stronger bias against GRBs and XRFs from a dusty environment.  If our lightcurves evolve analogous to the prompt stage of GRBs emission, the rapid decline coupled with their intrinsic faintness would make them difficult to monitor beyond several hours.  Finally, we emphasize the diversity of variable objects in the stellar menagerie, and thus it is not straightforward to rule out Galactic stars as the sources for our OTs.  Spectroscopic information will ultimately help to clarify their nature, and such followup observations are planned.  Our search has also yielded SN--like events that appear to have no host galaxy to significant ($> 27$) limiting magnitude -- OT 20020112 is one such example.  In addition, our catalog of SN candidates, classified as supernovae primarily due to their proximity to a host galaxy, might also contain sources with unusual temporal evolution.  Overall, the DLS transient search is well suited to explore the parameter space of OTs with small energy budgets, a population that could plausibly have escaped detection by gamma--ray and X--ray satellite missions.  This raises the possibility that the phenomena detected represent a new class of astronomical variability. Coordinated photometric followup of future optical transients is absolutely necessary to reveal fully the diversity of variability at faint optical magnitudes.  A primary goal of future variability surveys must be to enable the photometric and spectroscopic followup of detected events through timely release of information and ease of access to available data.  The wealth of information that can be gleaned from real--time synoptic transient science, only a subset of which is covered in this paper, strengthens the science case for the expansion of deep, wide astronomical surveys into the short timescale regime.  Expected future surveys like the LSST and PAN--STARRS \\citep{pan-starrs} will survey thousands of square degrees per night, and will benefit greatly from modern development in this field."
    },
    "0404/astro-ph0404285_arXiv.txt": {
        "abstract": "We present the results from an analysis of Hubble Space Telescope High Resolution Camera data for the Large Magellanic Cloud microlensing event MACHO-LMC-5. By determining the parallax and proper motion of this object we find that the lens is an M dwarf star at a distance of $578^{+65}_{-53}$pc with a proper motion of $21.39 \\pm 0.04 $ mas/yr. Based on the kinematics and location of this star it more likely to be part of the Galactic thick disk than thin disk population. We confirm that the microlensing event LMC-5 is a jerk-parallax event. ",
        "introduction": "For over a decade astronomers have been observing the Magellanic Clouds in order to determine the fraction of the dark matter in our Galaxy that may be in the form of Massive Compact Halo Objects (MACHOs).  The discovery of a significant number of microlensing events in the first two years of the MACHO project lead to an uncertain initial estimate that approximately half of the halo was composed of MACHOs (Alcock et al.~1997). With 3.7 years of additional data this estimate decreased to $\\sim 20\\%$ of the halo (Alcock et al.~2000).  While it appears these objects make up a significant fraction of the mass in the Galactic halo, little is known about their nature other than that their most probable masses lie in the range of 0.15 and 0.9$M_\\sun$  In order to obtain the most accurate information which can be gained from a microlensing event, it is useful to accurately determine the flux of the star that was lensed.  This is made difficult for sources in crowded fields such as LMC because they are usually blended with neighboring stars.  In many cases each ``object'' identified in ground-based observations consists of the blend of a number of stars (Alcock et al.~2000).  The exact location of the source star is also poorly known because of blending. To determine the locations of the sources Alcock et al.~(2001a, 2001b) analyzed the MACHO project images using Difference Image Analysis (Alcock et al.~1999, 2001a). With these positions they were able to subsequently identify and photometer the microlensing source stars in observations taken with {\\it Hubble Space Telescope} (hereafter HST) Wide Field Planetary Camera 2 (WFPC2).  Among the events discovered by the MACHO project toward the Magellanic clouds was event LMC-5.  This event had a high magnification ($\\sim$50) and was detected in the light curve of Macho object 6.5845.1091 which is located at $\\alpha= 05^{\\rm h}\\!16^{\\rm m}\\!41\\fs 1$, $\\delta =-70\\arcdeg\\! 29\\arcmin\\! 18\\arcsec$ (J2000). Gould, Bachall \\& Flynn (1997) suggested that the baseline color of this event was not consistent with an LMC source star and they proposed that the anomalous color could be attributed to the source being blended with a M dwarf in the Galactic disk. As the likelihood of finding and M dwarf within the seeing disk is small they proposed an M dwarf could be the lens in the foreground. Alcock et al.~(1997) found that the color of the source was in agreement with the colors and magnitudes of LMC stars, but the event was indeed blended with a red object.  Most of the LMC microlensing events found by Alcock et al.~(1997) were blended to some extent.  When the source star in this microlensing event was identified in HST observations it was discovered that there was a faint red star nearby.  The probability of finding an unrelated foreground M dwarf near the microlensing source star is $\\sim 1$ in 10000 (Alcock et al.~2001b).  With this in mind, it was thought very likely that this object was the lens.  The LMC-5 event shows the clear sign of microlensing parallax. In these events, the motion of the Earth during the event changes the shape of the microlensing light curve from the classical Paczy\\'nski form (Gould 1992, Alcock et al.~1995). The presence of parallax enables limits to be placed on the mass and location of the microlens.  The parallax fit for this event yielded a lens motion direction that was consistent with the red star having been the lens.  However, the solution also suggested that the lens was likely to be a sub-stellar object of 0.036M$_\\odot$ ($\\leq$ 0.097M$_\\odot$ at 3-$\\sigma$ significance) at a distance of $\\sim 200$pc.  Alcock et al.~(2001b) derived a separate distance estimate for the lens using the objects color and spectral type.  First a spectrum of the lens-source combination was obtained and it was found that the prospective lens was an M4V or M5V type star.  The V-I color of the object was determined from HST WFPC2 photometry.  This color was converted into an absolute magnitude using the $M_{V}$ vs $V-I$ relation of Reid (1991) for M dwarfs. The distance was then obtained from the observed and absolute magnitudes while the errors in distance were estimated from the dispersion in M dwarf magnitudes $V-I\\sim3$. The result was that the object was at $650 \\pm 190$pc, in stark contrast to the microlensing parallax solution.  If the red object is indeed the lens, a measurement of the parallax should confirm the distance inferred from the color of the object in the HST images.  In addition, a measurement of the magnitude and direction of the proper motion should agree with the initial estimate which assumed the red object was the lens, and that the relative separations of the two objects represented its proper motion.  To resolve the nature of the candidate lens we undertook a program of observations with HST's Advanced Camera System High Resolution Camera (hereafter ACS and HRC, respectively).  In the meantime, a new solution to the LMC-5 puzzle was proposed by Gould (2004) based on the recent work of Smith, Mao and Paczy\\'nski (2003).  By exploring the phase space of ``vector microlens parallax'' in a geocentric reference frame, Gould (2004) discovered a second solution to the microlensing parallax which varied from the original solution of Alcock et al.~(2001b) by less than 0.1 in fit $\\chi^2$.  The microlens vector parallax of this second solution differs from that of the first solution by the so called \"jerk parallax\", a vector whose direction lies perpendicular to the direction of the Earth's acceleration and whose magnitude (for LMC events) is about $(4/3) ({\\rm yr/2\\pi t_{\\rm e}}) \\sim 2.4$. Events exhibiting these so-called \"jerk-parallax degeneracies\" are expected to be rare for microlensing toward the LMC, unless the lens resides in the Galactic disk.  In the case of the LMC-5 microlensing event, Gould (2004)'s jerk-parallax solution is in agreement with the lens distance and direction estimated from the HST photometry.  The solution of Gould (2004) does not rule out the possibility that the initial solution to the microlensing fit was the correct one, since both are equally good fits to the lensing light curve.  However, when this solution is considered in combination with the other evidence from the HST data it is much more likely that the lens is a sub-stellar object not detected in the HST data. In this paper we will show with certainty that Gould (2004)'s solution is the correct one. ",
        "conclusions": "We have analyzed HRC data for LMC microlensing event LMC-5 and we find that the lens in this microlensing event is an M dwarf star.  Based on our analysis we can confirm that the jerk-parallax solution to the microlensing light curve discovered by Gould (2004) is correct.  The kinematics of this star suggest that it is most likely a part of the Galactic thick disk population rather than part of the dark halo. This is the first time that any microlens has been identified with such certainty.  However, this discovery does not affect the current estimates of the mass fraction of the Galactic dark halo in the form of MACHOs, since some microlensing events due to foreground disk stars are expected in all LMC microlensing models.  We would like to thank an anonymous referee for his many helpful suggestions. We would also like to thank Jay Anderson who generously made his results and analysis programs available to us prior to their release. Support for this publication was provided by NASA through proposal numbers GO-9394 and from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, under NASA contract NAS5-26555. This work was performed under the auspices of the U.S.~Department of Energy National Nuclear Security Administration by the University of California, Lawrence Livermore National Laboratory under contract W-7405-Eng-48.  The work done by A.~J.~D. is supported by Chilean FONDECYT grant 1030955."
    },
    "0404/astro-ph0404599_arXiv.txt": {
        "abstract": "\\baselineskip 14pt We discuss the constraints that future photometric and spectroscopic redshift surveys can put on dark energy through the baryon oscillations of the power spectrum. We model the dark energy either with a perfect fluid or a scalar field and take into account the information contained in the linear growth function. We show that the growth function helps to break the degeneracy in the dark energy parameters and reduce the errors on $w_{0},w_{1}$ roughly by 30\\% making more appealing multicolor surveys based on photometric redshifts. We find that a 200 $\\deg ^{2}$ spectroscopic survey reaching $z\\approx 3$ can constrain $w_{0},w_{1}$ to within $\\Delta w_{0}=0.21,\\Delta w_{1}=0.26$ and to $\\Delta w_{0}=0.39,\\Delta w_{1}=0.54$ using photometric redshifts with absolute uncertainty of 0.02. In the scalar field case we show that the slope $n$ of the inverse power-law potential for dark energy can be constrained to $\\Delta n=0.26$ (spectroscopic redshifts) or $\\Delta n=0.40$ (photometric redshifts), i.e. better than with future ground-based supernovae surveys or CMB data. ",
        "introduction": "Dark energy is the most elusive component of the universe: it is dark, it does not cluster and, in most models, it is subdominant in cosmic history until recently. On the other hand, it is also the most pervasive component, in that it affects background expansion, linear dynamics and possibly non-linear as well.  The characterization of DE has been so far based almost uniquely on background tests at rather low redshifts ($z<1$: Riess et al. 1998, Perlmutter et al. 1999, Tonry et al. 2003, Riess et al. 2004) or very large redshifts ($z\\approx 1000$: e.g. Netterfield et al. 2002, Halverson et al. 2002, Lee et al. 2002, Bennett et al. 2003), although tests at intermediate redshifts based on strong lensing have also appeared (Soucail, Kneib \\& Golse 2004). These tests are based essentially on estimations of luminosity $D_{L}$ or angular-diameter distances $D(z)$, i.e. on integrals of the Hubble function $H(z)$ which, in turn, contains an integral of the equation of state . If we aim at making real progress it is then necessary to cross this information with probes at intermediate $z$ and with other observables such as the linear growth function $G(z)$. The goal of this paper is to show how future observations of large scale structure (LSS) at $z$ up to 3 can set interesting limits on models of DE taking advantage of the complementarity between $H,D$ and $G$.  The method we use is based on recent proposals (Linder 2003, Blake \\& Glazebrook 2003, Seo \\& Eisenstein 2003) to exploit the baryon oscillations in the power spectrum as a standard ruler calibrated through CMB acoustic peaks. In particular, Blake \\& Glazebrook (2003) and Seo \\& Eisenstein (2003; hereinafter SE) have shown the feasibility of large (100 to 1000 square degrees) spectroscopic surveys at $z\\approx 1$ and $z\\approx 3$ to put stringent limits to the equation of state $w(z)$ and its derivative. SE have also shown that the redshift error introduced by a photometric survey (with a relative error on $1+z$ of $1\\%$) could be compensated only by a 20-fold increase of the sampled volume.  This paper has a threefold motivation. First, we extend the Fisher matrix method by taking into account the information on the DE evolution contained in the growth factor $G$ at various redshifts. While SE marginalized over $G$, we find that its marked dependence on the cosmological parameters, especially on $\\Omega _{m}$, associated with its different direction of degeneracy, are helpful in narrowing down the constraints. Secondly, we extend the analysis to a scalar field model of DE with inverse power-law potential (Ratra \\& Peebles 1988), based on the inverse power-law potential: this allows us to evaluate the growth function without further assumptions and to exploit the properties of tracking solutions (Steinhardt et al. 1999). Third, we pay special attention to the possibility of testing DE with a photometric redshift survey. Indeed the effect of the growth factor in reducing the confidence regions of the cosmological parameters makes more appealing the use of photometric redshift surveys to show any possible redshift evolution of the DE. We explore the constraints derived on the cosmological parameters from relatively wide multicolor surveys where the absolute redshift uncertainty keeps almost constant around $\\delta z=0.02-0.04$ in the redshift range $z=0-3.5$. Such surveys can be realized with the current wide field imagers on 4-8m class telescopes.  Finally, it is interesting to speculate on another aspect of the method. The baryon oscillation scale measures the angular-diameter distance $D(z)$ while the SNIa method observes the luminosity distance. As remarked recently by Bassett \\& Kunz (2003) the two quantities are related in a trivial way only in standard cosmologies: new physics like photon decay can lead to an observable breaking of their {}``reciprocity'' relation. A photometric survey can be adapted to the scope since the degradation of the error in $D(z)$ due to the redshift smearing is relatively reduced (with respect to the error in $H(z)$), as it will shown below. ",
        "conclusions": "Future deep and wide redshift surveys will have the chance to explore the little known landscape at $z\\approx 1-3$ between the recent and early universe. This is the epoch at which dark energy begins to impel its thrust to the expansion: it is therefore of extreme interest for cosmology in order to gain hold of reliable data on its evolution. As it has been shown by the recent releases of high redshift SN data (Tonry et al. 2003, Riess et al. 2004) the use of standard candles will have much to improve before it will be able to set stringent limits on the dark energy equation of state evolution. Moreover, there are still no solid proposals for standard candles visible at redshifts significantly higher than $z\\approx 1.5$. Other probes of the expansion dynamics at larger distances (gamma-ray bursts, lensing) might be promising but depend on unknown physics and modeling.  The use of the baryon oscillations in the matter power spectrum is on the contrary based on well-known phenomena that have already seen a spectacular validation in CMB observations. Moreover, it can provide a tomography of the universe expansion from $z=0$ down to $z\\approx 3.5$, virtually without loss of precision. An intrinsic advantage of this method is the combined use of $H,D$ and $G$ to reduce the degeneracies of each of these quantity. We have shown that the inclusion of $G$ reduce the errors by 30\\% roughly. We found that a 200$^{2}$ survey with absolute redshift error of $\\delta z=0.02$ can limit $w_{0},w_{1}$ by 0.39 and 0.54, respectively, assuming $w_{0}=-1,w_{1}=0$ as reference values. A spectrocopic survey can reduce the uncertainties to 0.21 and 0.26. If the dark energy is modeled as a scalar field with inverse power-law potential $n$, the limits are $n<0.26$ (spectroscopic redshift) and $n<0.40$ (redshift error $\\delta z=0.02$), along with an estimation of $\\Omega _{m0}$ to better than 3\\%. Inclusion of supernovae data further reduce the errorbars due to the different direction of degeneracy. The reduction is found to be particularly efficient in the case of photometric redshift measurements."
    },
    "0404/astro-ph0404550_arXiv.txt": {
        "abstract": "We present results based on an XMM-Newton observation of the high luminosity narrow-line QSO PHL~1092 performed in 2003 January. The $0.3-10$ keV spectrum is well described by a model which includes a power-law ($\\Gamma\\sim2.1$) and two blackbody components ($kT \\sim 130$ eV and $kT \\sim 50 $ eV). The soft X-ray excess emission is featureless and contributes $\\sim 80\\%$ to the total X-ray emission in the $0.3-10$ keV band. The most remarkable feature  of the present observation is the detection of X-ray variability at very  short time scale: the X-ray emission varied by 35\\% in about 5000 s. We find that this variability can be explained by assuming that only  the overall normalization varied during the observation. There was no evidence for any short term spectral variability and the spectral shape was similar even during the {\\it ASCA} observation carried out in 1997. Considering the high intrinsic luminosity ($\\sim$ 2 $\\times$ 10$^{45}$ erg s$^{-1}$) and the large inferred mass of the putative   black hole ($\\sim1.6\\times10^8 M_{\\sun}$), the observed time scale of variability indicates emission at close to Eddington luminosity arising from very close to the black hole. We suggest that PHL~1092 in particular (and narrow line Seyfert galaxies in general) is a fast rotating black hole emitting close to its Eddington luminosity and the X-ray emission corresponds to the high-soft state seen in Galactic black hole sources. ",
        "introduction": "\\par Narrow Line Seyfert 1 (NLS1) galaxies were first classified as a subclass of Seyfert 1  by \\citet{op85} on the basis that the FWHM of H$_\\beta$ line is less than 2000 km s$^{-1}$, the ratio of [O~III]$\\lambda 5007$ to H$_\\beta$ is less than 3, and they often show strong Fe~II emission in $4500 - 4680$ {\\rm \\AA} and $5105 - 5395$ {\\rm \\AA} regions. Though NLS1s are identified by their optical properties they have even more remarkable properties in the X-ray band as compared to other Seyfert galaxies. They show evidence for strong excess of soft X-rays (dominant below $\\sim$ 2 keV) above the hard X-ray continuum extrapolation and rapid X-ray variability (doubling times of minutes to hours).  \\par Soft X-ray excess emission above a power-law continuum is usually identified as the steepening of the X-ray continuum below $\\sim 2$ keV. This soft excess emission was first observed by {\\it HEAO-1} \\citep{pnnjwb81} and {\\it EXOSAT} \\citep{arn85,sgn85}. \\citet{bbf96} showed that AGN with steepest soft X-ray spectra in the {\\it ROSAT} band tend to lie at the lower end of the H$_\\beta$ line width distribution and hence AGNs with soft excess emission are predominantly NLS1 galaxies. Detailed X-ray study of NLS1s was carried out with {\\it ASCA} which confirmed the soft X-ray emission and, in addition, showed that the $2-10$ keV continuum slope too is steeper and anti-correlated with the FWHM of the H$_\\beta$ line \\citep{bme97,lei99b,vrwe99}. {\\it ASCA} observations also showed that many NLS1s have Fe~K$_\\alpha$ line  arising from highly ionized iron, and hence the accretion disks of NLS1s must be ionized \\citep{dew02,prob04,bif01}. NLS1s are the most extreme X-ray variable objects among the radio-quiet AGNs. NLS1s frequently exhibit rapid (doubling time scale of a few hundred seconds) and/or  large (up to a factor of 100) amplitude X-ray variability \\citep{bbf96,dsjmmn01,dbsl02}. The excess variance for NLS1s is typically an order of magnitude higher than that observed for Broad Line Seyfert 1's (BLS1).  \\par There are high-luminosity analogs of this class, the prototype of which is I~Zw~1 \\citep{phi76}. I~Zw~1 also has weak forbidden lines, narrow permitted lines, and strong Fe~II lines. One of the most extreme narrow line Fe~II QSO is PHL~1092 (B=16.7, z=0.396) \\citep{bk80,bk84,kcfzg95}. Its Fe~II$\\lambda4570/H_\\beta$ ratio is 5.3, and its H$_\\beta$ line width is  1800~km s$^{-1}$ \\citep{bk84}. PHL~1092 was rapidly variable during a {\\it ROSAT} observation, and its soft X-ray spectrum was extremely steep, modeled by a power-law with photon index more than 4 \\citep{fh96,bbfr99}. From an {\\it ASCA} observation \\citet{lei99b} found a huge soft excess (which can be described as a black body) in the PHL~1092 spectrum which is 80\\% of the total flux. An evidence for a line at around 7 keV has also been found and this line could not be explained as an ionized iron line.  \\par  Recently \\citet{dg04} made a detailed study of the XMM-Newton observations of I~Zw~1 and they showed that during the X-ray variability the power-law component varied without any change in the soft X-ray excess component. In contrast to this, PHL~1092 shows large variability and it has more than 80\\% soft X-ray excess component. To understand the X-ray spectral variability of this high luminosity object we have analyzed  an {\\it XMM-Newton} observation on this source. We present results based on this observation, including flux selected spectroscopy of high and low flux states.   \\begin{figure} \\includegraphics[angle=-90,scale=0.35]{f1.eps} \\caption{PN light curves of PHL~1092 in the 0.3 -- 0.5 keV (top panel), 0.5 -- 1 keV (middle panel) and 1 -- 10 keV (bottom panel) bands with a bin size of 512 s. The high and low flux states used for flux selected spectroscopy are shown by dotted vertical lines in the top panel of the figure. \\label{f1}} \\end{figure} ",
        "conclusions": ""
    },
    "0404/astro-ph0404544_arXiv.txt": {
        "abstract": "We report the first detection of X-ray emission from the Quintuplet star cluster and compare its X-ray emission to that of the Arches star cluster.  Four point sources are significantly detected near the core of the Quintuplet cluster with a total, absorption-corrected luminosity of $\\sim1\\times10^{33}$ ergs s$^{-1}$.  Diffuse, thermal emission is also detected near the core of the Quintuplet cluster with an absorption-corrected luminosity of $\\sim1\\times10^{34}$ ergs s$^{-1}$.  We analyze the diffuse and point-like emission from the Arches and Quintuplet and discuss the possibility that they are host to cluster wind outflows.  We also present the results of our search for X-ray emission from candidate star clusters in the Galactic center (GC) region.  We use extinction estimated by near-IR colors and X-ray spectral fits, as well as other IR properties, to determine if the candidate clusters are new, GC star clusters.  We find that three of the six candidate clusters found toward the GC are likely foreground clusters, two of the candidate clusters are not detected in the X-ray data, but have near-IR extinctions consistent with a GC location, and one of the candidate clusters has X-ray and near-IR extinctions consistent with being in the GC.  The X-ray and IR emission from the candidate clusters is compared to the known, massive, GC star clusters. ",
        "introduction": "The Galactic center (GC) region is host to some spectacular stellar clusters.  The Arches, Quintuplet, and central clusters are all young and among the densest in the galaxy.  The distribution of these clusters within 30 pc (12\\damin5 assuming a distance of 8.5 kpc) of Sgr A* may not be coincidental.  The ambient conditions of the GC region are known to be extreme (e.g., high gas and stellar densities, intense ionizing flux;  Morris \\& Serabyn 1996) and should have a significant effect on the process of star formation.  The currently known massive, GC star clusters (i.e. the Arches, Quintuplet, and IRS 16 clusters) were discovered in surveys in the near-IR (Becklin \\& Neugebauer 1975; Nagata et al.\\ 1990; Cotera et al.\\ 1996), where interstellar extinction is much less than at visible wavelengths.  Although the near-IR is an excellent regime in which to search for extincted stellar emission, the GC is particularly challenging region to survey because of confusion with stars lying along the line of sight, as well as the nonuniform distribution of extinction found there.  The Arches cluster was recently detected at X-ray wavelengths (Yusef-Zadeh et al.\\ 2002, hereafter Y02), where the emission is thought to originate from the shocked gas in the collision between individual stellar winds, as had been modeled specifically for the Arches (Cant\\'{o} et al.\\ 2000; Raga et al.\\ 2001).  These studies suggest that cluster wind emission may provide a new window into the physics of dense star clusters, although it can be difficult to separate putative cluster wind emission from unresolved stellar cluster members.  The Quintuplet cluster is another of these unusual, massive, GC star clusters.  The Quintuplet, named after its brightest five stars (Nagata et al.\\ 1990), is unusually dense and is host to several massive, windy, Wolf-Rayet stars.  However, the Quintuplet is somewhat less massive and dense than the Arches cluster, suggesting that it has been more dissolved by GC tidal forces (Kim et al.\\ 1999; Portegeis-Zwart et al.\\ 2001).  Recent theoretical and observational studies have suggested that there may be many more undiscovered massive star clusters in the GC region.  \\emph{N}-body simulations of these dense, young clusters have suggested that star clusters such as the Arches dynamically evolve rapidly, dissolving into the stellar background within 10--20 Myr (Portegeis-Zwart et al.\\ 2001;  Kim et al.\\ 1999; Kim \\& Morris 2003).  Thus it was suggested that there may be another 10--50 massive clusters distributed within the central 150 pc at various stages of dissolving into the background stellar field.  On the observational side, Dutra \\& Bica (2000, 2001; hereafter DB00, DB01) and Dutra et al.\\ (2003) searched Two-Micron All Sky Survey (2MASS) images for Arches-like objects to make a list of candidate GC clusters.  Their results were consistent with rough estimates of the number of clusters predicted by Portegeis-Zwart et al.\\ (2001).  This paper discusses the use of X-ray and IR properties of the Arches and Quintuplet clusters to examine whether the cluster candidates identified by DB00 \\& DB01 are in the Galactic center region.  In \\S \\ref{known}, we extend our previous analysis of the emission from the Arches cluster and present the first detection of X-ray emission from the Quintuplet cluster.  Section \\ref{candidate} presents our analysis of the candidate star clusters found by Dutra \\& Bica;  in particular, X-ray and near-IR extinction are used to determine whether the candidate clusters are in the Galactic center.  We conclude by comparing the candidate clusters to the known GC clusters and commenting on the population of massive star clusters in the GC. ",
        "conclusions": "\\label{discussion}  \\subsection{X-Ray Sources:  The Arches and Quintuplet Clusters} The nature of many of the X-ray point sources in the Arches and Quintuplet clusters is made clear by the compelling correlation with near-IR and radio continuum point sources.  The X-ray emission is considered to be produced in shocks formed in the winds of massive stars (O and WR type);  the more luminous of these sources are most likely to be colliding-wind binaries (Stevens \\& Hartwell 2003).  Alternatively, for mass-segregated clusters such as the Arches cluster, it is possible that X-ray emission may also arise from accretion onto an intermediate mass black hole, as recent theoretical work suggests (G\\\"{u}rkan, Freitag, \\& Rasio 2003).  The collision of winds in a binary system can greatly enhance the X-ray luminosity, depending on the type of stars in the binary and the phase of the orbit (Pollock 1987).  It has been suggested that the collective effect of these wind collisions has recently been observed in the Arches cluster, host of the more luminous X-ray sources in the GC region.  Yusef-Zadeh et al.\\ (2003) have recently found diffuse nonthermal radio emission from the core of the Arches cluster, coincident with the A1N/S X-ray source.  Although some stellar wind sources can have nonthermal radio emission, this emission is not pointlike;  the source is extended to approximately the size of the cluster core, $9\\arcsec$.  This extended nonthermal radio emission is thought to be produced by diffuse shock acceleration (e.g., Bell 1978) in the colliding winds of the cluster.  These energetic electrons may be able to Compton upscatter stellar IR radiation into the X-ray regime.  X-ray spectra from the Arches core (A1N, A1S, and A2 sources) can be adequately fitted with a one-temperature/power-law model;  in this model, the fraction of flux attributed to the power law ($\\sim1/6$) is consistent with that expected from inverse Compton scattering.  The scaling of X-ray flux with intrinsic cluster properties can give us another view of the X-ray emission mechanisms. IR observations have shown that the ratio of mass as well as luminosity between the Arches and Quintuplet is roughly $3\\colon1$ (FMM99).  In this study, the ratio of background-subtracted, X-ray flux (0.5--8 keV) from the Arches and Quintuplet is found to be roughly $11\\colon1$ ($8\\times10^{-5}\\colon7\\times10^{-6}$ counts cm$^{-2}$ s$^{-2}$);  if only the point source emission is considered, the ratio rises to $18\\colon1$.  Correcting this observed flux for absorption is difficult considering the poorly constrained \\nh measurement toward the Quintuplet.  From near-IR observations (FMM99; Figer et al.\\ 2002), extinctions toward these clusters have been measured as \\ak$^{\\rm{Arches}}=3.1$ mag and \\ak$^{\\rm{Quint}}=3.28\\pm0.5$ mag (also \\ak$^{\\rm{QPM}}=2.7$ mag for the ``Quintuplet proper members'').  The near-IR extinctions toward the Arches and Quintuplet are similar within the uncertainties (which are somewhat higher, but still within the errors of the near-IR extinction expected from the best-fit \\nh value), so the absorption-corrected X-ray luminosities should scale similarly to the observed X-ray fluxes.  The bulk of the difference in X-ray flux between the Arches and Quituplet is caused by the presence of luminous point sources in the Arches. Theoretical studies have shown how massive star clusters are subject to processes which are more likely to create X-ray luminous sources.  For example, in N-body simulations, massive stars are found to congregate near the center of their host clusters, where collisions (forming tight binaries or even mergers) are more likely (Portegies Zwart et al.\\ 1999).  Massive clusters are much more prone to these ``core collapses,'' making them more likely to be host to massive, X-ray luminous, binary systems and intermediate-mass black holes (G\\\"{u}rkan, Freitag, \\& Rasio 2003).  It may be the proliferation of such systems in the Arches that makes it so much more X-ray luminous than the Quintuplet.  As shown in \\S \\ref{quint} and Figure \\ref{quintbig}, the Quintuplet cluster has diffuse emission and a clear peak in the gas temperature near the cluster core.  In contrast to the Arches diffuse emission, the Quintuplet diffuse emission is clearly thermal, with thermal line emission evident from highly ionized species of S, Si, Ar, Ca, and Fe.  The peak in gas temperature coincident with the cluster core may be due to any or all of the following: (1) an increase in the temperature of truly diffuse emission, (2) an enhancement of nonthermal emission in the cluster core, and (3) the presence of unresolved, discrete sources.  If the emission is truly diffuse, the high temperature in the cluster core suggests that the Quintuplet could be the source of gas heating.  This would be consistent with the model of a cluster wind flow, in which the wind adiabatically cools as it escapes from the Quintuplet cluster (Raga et al.\\ 2001).  However, a nonthermal (powerlaw) spectrum has a shape similar a high temperature thermal continuum; the peak in gas temperature may indicate a larger contribution of nonthermal emission in that region.  A third alternative is that the peak in gas temperature at the cluster core could be due to unresolved cluster members.  In this case, hot stellar emission may cause the increase in gas temperature while the outer cooler gas is the genuinely diffuse, ambient gas seen throughout the GC region.  The best-fit temperature in the Quintuplet core is hotter than that expected from single O or WR stars ($<\\sim1$ keV; Chlebowski, Harnden, \\& Sciortino 1989);  this is also true of the temperature of the emission from the detected X-ray sources in the cluster core.  If this emission is due to unresolved stellar systems, a significant portion of these systems should be colliding wind binaries, which tend to be hotter and/or may emit nonthermal radiation (Pollock 1987).  While young O stars such as $\\Theta^{1}$ Ori A can exhibit hard X-ray flares (Schulz et al.\\ 2003), such hard X-rays are less common in older, less magnetically active stars such as those found in the Quintuplet cluster (with an age of $3-5$ Myr, as compared to $\\sim0.3$ Myr for $\\Theta^{1}$ Ori A).  Finally, the luminosity of this nonpoint-source emission from the Quintuplet core is within a factor of two of that of the Arches A3 source (L$_{\\rm{x}}^{Q diff}\\sim1\\times10^{34}$ ergs s$^{-1}$ compared to L$_{\\rm{x}}^{A3}\\sim1.6\\times10^{34}$ ergs s$^{-1}$ according to Y02, or L$_{\\rm{x}}^{A3}\\sim6\\times10^{33}$ ergs s$^{-1}$ if we reanalyze assuming the canonical GC \\nh value).  \\subsection{X-ray Sources and Extinctions of Candidate Clusters} \\label{canddisc} Any X-ray source detected in candidate clusters that lie in or beyond the GC would most likely be shocked-wind emission observed from O and WR stars or wind-accreting neutron stars (Muno et al.\\ 2003, and references therein).  For example, a limiting flux of $2.2\\times10^{-6}$ counts cm$^{-2}$ s$^{-1}$ (typical of our observations of the candidate clusters) could detect a 1 keV thermal source in the GC with a luminosity, L$_{\\rm{x}}\\sim5\\times10^{33}$ ergs s$^{-1}$;  this temperature and luminosity is similar to that of colliding wind binary systems (e.g., Portegies Zwart, Pooley, \\& Lewin 2002).  Nonthermal sources tend to be spectrally hard, and thus are more likely to be detected in the GC.  A $\\Gamma=2$ powerlaw would be detectable at our typical flux limit, if it had a luminosity, L$_{\\rm{x}}\\sim5\\times10^{32}$ ergs s$^{-1}$.  Candidate DB01-42 has an \\nh that is consistent with a GC location, so its X-ray sources have an origin in one of the above mechanisms.  Of the two X-ray sources in DB01-42,  one has a counterpart with $K_{s}\\sim13$ mag and another has no counterpart detected by 2MASS, down to the local detection limit of $K_{s}\\sim13$; this is consistent with Pfahl, Rappaport, \\& Podsiadlowski (2002), which finds that most neutron star companions should have intrinsic magnitudes of $K>11$ mag, or $K>13$ mag for a typical GC-like extinction.  The flux limit of this survey is adequate to detect the more luminous colliding-wind binary systems in the GC (Pollock 1987).  These systems tend to consist of early-O or WR type stars, which would have $K$-band magnitudes brighter than 13th magnitude, with the GC extinction and distance.  Spectrally soft sources in the GC would need to be dramatically more luminous (L$_{\\rm{x}}>10^{34}$ ergs s$^{-1}$) to be detected in this survey; sources at our detection limit which are nearer than the GC will be much less luminous and are more likely to be CVs, main sequence stars, or young stellar objects.  Measuring the extinction toward these clusters can help us determine where they are with respect to the GC.  As shown in \\S \\ref{candidate}, candidate DB01-42 has associated X-ray emission and possible X-ray/IR correlated sources, which allows a comparison of two independent measurements of the cluster's extinction.  The X-ray extinction for this candidate is found to be consistent with typical GC values;  the IR reddening of the X-ray/IR source is somewhat less, but consistent within the errors to a GC reddening.  Figure \\ref{jhhkfig} and Table \\ref{properties} reinforce this point by showing that the majority of the IR sources near this cluster candidate are also reddened like the GC.  DB00-1 and DB00-55 have no X-ray emission with which to measure an X-ray extinction.  Interestingly, the reddening estimated from near-IR point sources show that the candidates are extincted as if they were in the GC.  Color-color diagrams from regions near these candidates do not show the same density of highly extincted sources, validating the idea that these candidates are distinct associations and are located in the GC.  Two of the candidates with X-ray emission (DB00-6 and DB00-58) have X-ray absorbing columns inconsistent with typical GC values.  A third candidate, undetected in X-rays (DB00-5), is located in the same {\\small{H II}} region as a low-absorption candidate (DB00-6), suggesting that it too is not in the GC.  Figure \\ref{jhhkfig} shows that the X-ray/IR sources for these clusters have near-IR extinctions that are relatively small (\\av$<20$).  We can also see this by comparing the number of sources with and without GC-like reddening (Col. [6] of Table \\ref{properties}); Candidates DB00-5, DB00-6, and DB00-58 all have the highest fraction of sources with non-GC-like $H-K_s$ colors.    \\subsection{New GC clusters?} \\label{newgcc} The known GC clusters are similar in four of the diagnostics presented in Table \\ref{properties}:  density, size, reddening, and mid-IR emission.  No candidate cluster has all these characteristics to the degree seen in the known clusters.  Using the simplest of these tests, extinction, three candidate clusters were found unlikely to be in the GC, either directly (DB00-6, DB00-58), or by association with other non-GC candiates (DB00-5).  Only one of the candidates with X-ray emission is likely located in the GC:  DB01-42.  With an observed flux of $\\sim5\\times10^{-6}$ counts cm$^{-2}$ s$^{-1}$, highly extincted X-ray and near-IR emission, this candidate resembles the Quintuplet cluster (although they are much different in angular size: $12\\arcsec$ as compared to $40\\arcsec$).  The coincidence of DB01-42 inside the bow-shaped {\\small{H II}} region, G359.3--0.3, suggests that the candidate cluster may be the source of the ionization and dust heating seen in radio and mid-IR images.  Although the X-ray and near-IR emission clearly originate in the GC, the velocity of the {\\small{H II}} region is unusually small for the GC, suggesting that G359.3--0.3 may not be in the GC.  Regardless of the possible correlation with G359.3--0.3, X-ray and near-IR emission indicate that DB01-42 is likely to be a new GC cluster.  Future X-ray observations could search for diffuse emission, a tell-tale signature of young and/or massive stars.  Finally, candidates DB00-1 and DB00-55 have no X-ray emission, but their near-IR sources have GC-like reddening;  this suggests that these clusters are likely to be near the GC, but X-ray data are not sensitive enough to detect them.  The Arches and Quintuplet clusters have background-subtracted fluxes, $F_{X}(0.5$--$8$keV$)=9\\times10^{-5}$ and $7\\times10^{-6}$ counts cm$^{-2}$ s$^{-1}$, respectively;  these fluxes are larger than the X-ray flux limits of both undetected candidate clusters (see Table \\ref{properties}).  If the candidate clusters were as luminous and extincted as the Arches and Quintuplet, then the candidates should have been detected by these observations.  However, it may be possible that the candidates have X-ray luminosities like the Quintuplet (the fainter of the known clusters), but are more extincted.  Simple spectral simulations show that if candidate DB00-1 had a luminosity like the Quintuplet, but an \\nh $>1.5$ times greater, then the cluster would not be detected by current observations.  By the same analysis, DB00-55 could have the Quintuplet's flux and not be detected if \\nh was $>$3 times higher than the Quintuplet's.  It is more likely that the undetected candidate clusters are intrinsically less luminous than the Quintuplet, and thus they are probably less massive or are not host to any of the massive stars found in the Quintuplet.  These candidates are also the only two clusters not found to have any mid-IR emission, suggesting that they are somewhat older than the other clusters in this study; in fact, Dutra et al.\\ (2003) suggested this for candidate DB00-1, but not DB00-55.  One possibility is that DB00-1 and DB00-55 are stellar clusters in the process of being dissolved by the GC's strong tidal gravitational field.  However, all we can conclusively say is that these clusters seem to be significantly less luminous than the known, massive, GC clusters.  Portegies Zwart et al.\\ (2002) predicted 10--50 clusters similar to the known, massive clusters in the central 200 pc ($\\pm1$\\sdeg), at varying stages of dissolving into the stellar background.  The near-IR searches of DB00 and DB01 found many candidates and Dutra et al.\\ (2003) confirmed that most of them are stellar clusters, if not necessarily in the GC.  Of the six of these candidates which lie in the central 2\\sdeg$\\times0\\ddeg8$, we have found that:  three are \\emph{not} located in the GC, two \\emph{are} in the GC but are much less luminous than known GC clusters, and one is a reasonable candidate for a Quintuplet-like GC cluster or association.  Further X-ray and spectroscopic near-IR observations are necessary to solidify the association of these candidates with the GC and to search for the unusual stellar signatures seen in known GC clusters.  We also note that the simple model of Portegies Zwart et al.\\ (2002) predicts several more clusters within the central 200 pc;  further searches of 2MASS and DENIS data archives may prove fruitful."
    },
    "0404/astro-ph0404258_arXiv.txt": {
        "abstract": "We describe our new \"\\mlapm-halo-finder\" (\\mhf) which is based on the adaptive grid structure of the \\nbody\\ code \\mlapm. We then extend the \\mhf\\ code in order to track the orbital evolution of gravitationally bound objects through any given cosmological \\nbody-simulation - our so-called \"\\mlapm-halo-tracker\" (\\mht). The mode of operation of \\mht\\ is demonstrated using a series of eight high-resolution \\nbody\\ simulations of galaxy clusters.  Each of these halos hosts more than one million particles within their virial radii \\rvir. We use \\mht\\ as well as \\mhf\\ to follow the temporal evolution of hundreds of individual satellites, and show that the radial distribution of these substructure satellites follows a ``universal'' radial distribution irrespective of the host halo's environment and formation history. This in fact might pose another problem for simulations of CDM structure formation as there are recent findings by Taylor~\\ea (2003) that the Milky Way satellites are found preferentially closer to the galactic centre and simulations underestimate the amount of central substructure, respectively. Further, this universal substructure profile is anti-biased with respect to the underlying dark matter profile. Both the halo finder \\mhf\\ and the halo tracker \\mht\\ will become part of the open source \\mlapm\\ distribution. ",
        "introduction": "Over the last 30 years great progress has been made in the development of \\nbody\\ codes that model the distribution of dissipationless dark matter. Algorithms have advanced considerably since the first $N^2$ particle-particle codes (Aarseth 1963; Peebles 1970; Groth \\ea 1977); we have seen the development of the tree-based gravity solvers (Barnes~\\& Hut 1986), mesh-based solvers (Klypin~\\& Shandarin 1983), then the two combined (Efstathiou \\ea 1985) and multiple strands of adaptive and deforming grid codes (Villumsen 1989; Suisalu~\\& Saar 1995; Kravtsov, Klypin~\\& Khokhlov 1997; Bryan~\\& Norman 1998; Knebe, Green \\& Binney 2001). While they all push the limits of efficiency in computational resources, each code has its individual advantages and limitations. The result of such research has been highly reliable, cost effective codes. However, producing the data is only one step in the process; the ensembles of millions of (dissipationless) dark matter particles generated still require interpreting and then comparison to the real Universe. This necessitates access to analysis tools to map the phase-space which is being sampled by the particles onto ``real'' objects in the Universe; traditionally this has been accomplished through the use of ``halo finders''. Halo finders mine \\nbody\\ data to find locally over-dense gravitationally bound systems, which are then attributed to the dark halos we currently believe surround galaxies. Such tools have lead to critical insights into our understanding of the origin and evolution of structure and galaxies. To take advantage of sophisticated \\nbody\\ codes and to optimise their predictive power one needs an equally sophisticated halo finder.  Over the years, halo-finding algorithms have paralleled the development of their partner \\nbody\\ codes. We briefly outline the major halo finders currently in use:  The Friends-of-Friends (FOF) (Davis \\ea 1985; Frenk \\ea 1988) algorithm uses spatial information to locate halos. Specifying a linking length \\llink\\ the finder links all pairs of particles with separation equal to or less than \\llink\\ and calls these pairs ``friends''. Halos are defined by groups of friends (friends-of-friends) that have at least one of these friendship connections. Two such advantages of this algorithm are its ease of interpretation and its avoidance of assumption concerning the halo shape. The greatest disadvantage is its simple choice of linking length which can lead to a connection of two separate objects via so-called linking ``bridges''. Moreover, as structure formation is hierarchical, each halo contains substructure and thus the need for different linking lengths to identify ``halos-within-halos''.  There have been many variants to this scheme which attempt to overcome some of these limitations (Suto, Cen \\& Ostriker 1992; Suginohara \\& Suto 1992; van Kampen 1995; Okamoto \\& Habe 1999; Klypin \\ea 1999).  DENMAX (Bertschinger \\& Gelb 1991; Gelb \\& Bertschinger 1994a) and SKID (Weinberg, Hernquist \\& Katz 1997) are similar methods in that they both calculate a density field from the particle distribution, then gradually move the particles in the direction of the local density gradient ending with small groups of particles around each local density maximum. The FOF method is then used to associate these small groups with individual halos. A further check is employed to ensure that the grouped particles are gravitationally bound.  The two methods differ through their calculation of the density field. DENMAX uses a grid while SKID applies an adaptive smoothing kernel similar to that employed in Smoothed Particle Hydrodynamics techniques (Lucy 1977; Gingold~\\& Monaghan 1977; Monaghan 1992). The effectiveness of these methods is limited by the method used to determine the density field (G\\\"otz, Huchra \\& Brandenberger 1998).  A similar technique to the above is the Bound Density Maxima (BDM) method (Klypin \\& Holtzman 1997; Klypin \\ea 1999). In this scheme a smoothed density is derived by smearing out the particle distribution on a scale \\rsmooth\\ of order the force resolution of the \\nbody\\ code used to generate the data. Randomly placed ``seed spheres'' with radius \\rsmooth\\ are then shifted to their local centre-of-mass in an iterative procedure until convergence is reached. Hence, as with DENMAX and SKID, this process finds local maxima in the density field. Bullock \\ea (2001) further refined the BDM technique by first generating a set of possible centres, ranking the particles with respect to their local density and then implementing modifications which allow for credible identification of halos-within-halos.  The Bullock~\\ea (2001) adaptation to BDM excels at finding halo substructure.  When one is primarily concerned with distinct halos, all the mentioned methods perform exceedingly well. All efforts to refine and enhance those halo finding algorithms are due to the fact that \\nbody\\ codes overcame overmerging only recently (Klypin~\\ea 1999) and are capable of finding satellites galaxies within dark matter host halos.  It is therefore crucial to reliably identify ``halos-within-halos''. In fact, one of the remaining problems for simulations of CDM structure formation is that high-resolution simulations nowadays predict far greater substructure (in total) than observed (Klypin~\\ea 1999; Moore~\\ea 1999). Results from gravitational microlensing suggest that the majority of substructure which does exist has to be close to the inner regions (Dalal \\& Kochanel 2002) which thus far has not been confirmed by such simulations. There are recent claims that although the overmerging problem has disappeared in the outer regions of the halo, the inner regions might still suffer from it (Taylor, Silk~\\& Babul 2003). As these latter semi-analytic models do not suffer from such numerical problems, they find that such substructure does exist in the inner regions.  The question though arises as to whether there still remains an overmerging problem in the simulations or if current halo finding algorithms actually do break down at those scales. As we will discuss later, it becomes more difficult to locate peaks in the central region (if at all present) of the host halo due to a simple lack of contrast.  In this paper we present a new method for identifying gravitationally bound objects in \\nbody\\ code output that uses the adaptive meshes of \\mlapm\\ (Knebe et~al. 2001). This new code excelled at finding ``halos-within-halos'' revealing more substructure in the inner regions of the host halo.  In its native form, our new algorithm works naturally ``on-the-fly'', but it has also been constructed with the flexibility necessary to handle a single temporal output from any \\nbody\\ code. Our analysis software will become part of the publicly available \\mlapm\\ distribution\\footnote{\\texttt{http://astronomy.swin.edu.au/MLAPM/}}. The outline of the paper is as follows. In Section~\\ref{Computation} we introduce the cosmological models used to frame our discussion of the mode of operation of the new halo finder and tracker. A more detailed scientific analysis of this data set can be found in Paper~II of this series (Gill et~al. 2004a; hereafter, \\GKGD). In \\Sec{mhf} we introduce the new halo finder ``\\mlapm-halo-finder'' (\\mhf), describing its function, advantages, and limitations. \\Sec{analmhf} provides a brief analysis of the satellites found by \\mhf. In Section~\\ref{mht} we introduce the ``\\mlapm-halo-tracker'' (\\mht) which augments the halo finder by incorporating the ability to track the temporal evolution of satellites. Analysis of the halos tracked with \\mht\\ is described in \\Sec{analmht}. We next compare the two methods with other publicly available halo finding algorithms, such as FOF and SKID, in Section~\\ref{compare}. We conclude with a summary and our conclusions in \\Sec{conclusions}.  This paper is the first in a series of three based upon the suite of simulations described herein. Paper~II (\\GKGD) investigates the satellite environments and their dynamical properties, while Paper~III (Gill et~al. 2004b) will investigate the tidal streams and debris from the disrupting satellites.   \\begin{table*} \\label{HaloDetails} \\caption{Summary of the eight host dark matter halos. Distances are measured in \\hMpc, velocities in \\kms, masses in 10$^{14}$\\hMsun, and the age in Gyrs.} \\begin{tabular}{ccccccc}\\hline Halo & \\Rvir & $V_{\\rm circ}^{\\rm max}$ & \\Mvir  & \\zform & age & $N_{\\rm sat}(<r_{\\rm vir})$ \\\\  \\hline \\hline \\# 1 &  1.34 & 1125 & 2.87 & 1.16 & 8.30 & 158 \\\\ \\# 2 &  1.06 &  894 & 1.42 & 0.96 & 7.55 &  63 \\\\ \\# 3 &  1.08 &  875 & 1.48 & 0.87 & 7.16 &  87 \\\\ \\# 4 &  0.98 &  805 & 1.10 & 0.85 & 7.07 &  57 \\\\ \\# 5 &  1.35 & 1119 & 2.91 & 0.65 & 6.01 & 175 \\\\ \\# 6 &  1.05 &  833 & 1.37 & 0.65 & 6.01 &  85 \\\\ \\# 7 &  1.01 &  800 & 1.21 & 0.43 & 4.52 &  59 \\\\ \\# 8 &  1.38 & 1041 & 3.08 & 0.30 & 3.42 & 251 \\\\ \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "\\label{Summary and Conclusions} \\label{conclusions}  Computational cosmology is not only limited by crucial factors such as the dynamical range and the mass resolution, but also by its analysis tools. We emphasise, perhaps obviously to most readers, that \\nbody\\ codes simulating structure formation in the Universe can only ever be as useful as their associated analysis tools allow.  In this paper we presented two new methods for identifying gravitationally bound objects in such simulations. Both methods are based upon the open source adaptive mesh refinement code \\mlapm\\ (Knebe et~al.  2001). They both exploit the refinement hierarchy of said code and hence locate halos as well as halos-within-halos with exactly the same accuracy as \\mlapm\\ simulates their evolution. We showed the limitations of a simple snapshot analysis and how it can be overcome by taking into account the whole history of each halo. Thus not restricting the halo finding algorithm to just the spatial information but rather including the velocity information as well, hence, using the full six-dimensional information available.  Not only do we intend to implement the halo finder into the distribution of \\mlapm\\ allowing for an on-the-fly analysis saving both computational and human resources in the analysis process, but it is also our intention to have it as a stand alone program. In both cases, the implementation will be such that it can analyse a single output of any given \\nbody\\ code.  We showed that halo finding still possesses the inherent problem of overmerging in the very central regions of the host. However, by tracking satellites rather than identifying them at separate time snapshots of the simulation we learned that this problem is \\textit{not} overmerging in the conventional sense. The objects are in fact present and simulated properly, but their densities have insufficient contrast to be picked up by a simple ``finder''. Only when tracking them in time from (at best their very own) formation time were we able to quantify their existence as close as 5\\% of $r_{\\rm vir}^{\\rm host}$ to the cuspy centre. We do not intend to question the credibility of other (most excellent) halo finders such as SKID though. We rather pointed out that the results in the central region are subject to subtleties that can be most easily avoided by tracking satellites.  It has recently been pointed out by Taylor~\\ea (2003) that the radial distribution of the Milky Way's satellite galaxies does not reconcile with the predictions of semi-analytical models of galaxy formation and cosmological simulations, respectively, which has been confirmed by Kravtsov~\\ea (2004). Even though there is quite a prominent substructure population in the simulations it is mostly clustered in the outer regions of the host whereas the Milky Way satellites are preferentially found closer in.  It remains unclear if this poses a new challenge to the CDM structure formation scenario or simply reflects what we presented in this study. Namely, identifying substructure halos that lie close to the centre of the host is intrinsically challenging and might be overcome by actually \\textit{tracking} the satellites rather than \\textit{finding} them.   Finally, we further increase the statistics to suggest that anti-bias of the satellite distribution is a common property of the satellite distribution, perhaps even universal. However, this antibias is still most likely a relic of numerical overmerging. A convergence study for the satellite populations is critical in resolving this issue and reconciling the lensing observations. However, one very interesting result is the continued self-similarity in the dark matter distribution, the satellite radial distribution is common, once again perhaps \"universal\" even though the number of substructure satellites changes, reminiscent of the universal density profile of the underlying host."
    },
    "0404/astro-ph0404291_arXiv.txt": {
        "abstract": "We analyze the spectrum of an unusually-low mass white dwarf, SDSSJ123410.37-022802.9 (0335-264-52000), found in our recent, white dwarf catalog from the SDSS DR1 data release.  Two independent, model atmosphere fits result in an accurate determination of atmospheric and stellar parameters.  The more hands-on analysis yields $T_{\\rm eff}$ = 17,470$\\pm$750~K and $\\log g$ = 6.38$\\pm$0.05.  We argue that the object cannot be a main sequence A star, a horizontal branch or subdwarf B star.  Instead, it is interpreted as a very low mass white dwarf with a core composed of helium, with mass ($\\sim$0.18-0.19$M_\\odot$), similar to that published previously for the unusual companion to the millisecond pulsar J1012+5307.  The star probably remains in a binary, perhaps even with an undiscovered or dead pulsar companion.  However, the companion might be a more-ordinary star, provided Roche lobe mass transfer began shortly after the now-visible component left the main sequence.  A second SDSS low mass white dwarf candidate is also analyzed, but the spectrum for this fainter object is of poorer quality. The sample appears to include additional, similar candidates, worthy of more accurate observation.  Correct identification of any additional white dwarfs of extremely low mass requires careful observation and interpretation. ",
        "introduction": "Kleinman et al. (2004, hereafter K04) have released a catalog of 2561 white dwarf stars found in the first data release of the Sloan Digital Sky Survey (Abazajian et al. 2003, DR1; see also http://www.sdss.org/dr1).  As value added to the release, an automatic fitting procedure (autofit) was used to estimate the parameters $T_{\\rm eff}$ and $\\log~g$ for each white dwarf classified DA and DB.  Extensive human checks of the ``outliers'' in the $T_{\\rm eff}$ vs $\\log~g$ plot resulted in elimination of many invalid fits, due usually to a corrupted spectrum or flux calibration.  However, K04 noted several white dwarfs of unusually low surface gravity which appear to be valid, if usually inaccurate, fits.  While it has been known for a long time that the mass distribution of field DA white dwarfs peaks near 0.6~$M_\\odot$ (Koester, Schulz, \\& Weidemann 1980), the more precise determinations of $T_{\\rm eff}$ and $\\log g$ made possible by accurate digital spectra and model atmosphere fits to the Balmer lines demonstrated the reality of both separate high mass and low mass components to this mass distribution (Bergeron, Saffer, \\& Liebert 1992; Bragaglia, Renzini, \\& Bergeron 1995).  With derived masses $\\lesssim$0.5~$M_\\odot$, the latter were inferred to have helium cores, requiring envelope mass loss to a companion before the progenitor reached the tip of the red giant branch and could ignite the helium.  This prediction was dramatically confirmed by Marsh, Dhillon, \\& Duck (1995) and Maxted \\& Marsh (1999), who found that a large fraction of these stars and other low gravity DA white dwarfs were indeed binaries with orbital periods of hours to days.  The inferred companion was generally another white dwarf, sometimes detectable in the spectrum.  In a few cases, an infrared excess was discovered, which implied a low mass nondegenerate companion (Zuckerman \\& Becklin 1992).  The apparent helium white dwarfs in the Bergeron et al. (1992) and Bragaglia et al. (1995) samples of stars with $T_{\\rm eff}$ $<$ 50,000~K had $\\log~g$ as low as 7.22 and masses as low as 0.31~$M_\\odot$ (PG~1241-010).  Very hot, low mass white dwarfs have radii much larger than the final ``zero-temperature'' values.  For instance, in Bergeron et al.'s (1994) analysis of hot DAO stars, the lowest gravity fitted was for the star HZ34 at $\\log~g$ = 6.61 and $T_{\\rm eff}$ = 60,700~K. Probably the lowest mass white dwarf reported up to now has a most unusual evolutionary history.  It is the companion to the millisecond pulsar J1012+5307.  Presumably it spun-up the pulsar during the prior common envelope phase.  Van Kerkwijk, Bergeron, \\& Kulkarni (1996) found for the white dwarf $T_{\\rm eff}$ = 8,550$\\pm$25~K and $\\log g$ = 6.75$\\pm$0.07, from which they estimated a tiny mass of only 0.16$\\pm$0.02$M_\\odot$.  Callanan, Garnavich, \\& Koester (1998) derived $T_{\\rm eff}$ = 8670~K, $\\log~g$ = 6.34, but the same mass of 0.16$\\pm$0.02$M_\\odot$ in an independent analysis.  Several stars in the K04 DA sample fit surface gravities lower than those cited above.  For most of these mentioned in K04, Section 7.1, the spectrophotometry of these stars near 20th magnitude is too poor for accurate determinations of parameters.  Improved spectra would be recommended.  For convenience, the SDSS photometry for the stars mentioned in the cited paper is listed in Table~1.  The main purpose of this paper is to develop the case for an extremely low mass for the best-observed example, SDSSJ~123410.37-022802.9 (hereafter, SDSSJ1234), a result in which we have high confidence.  We do include analysis of one other low mass candidate, to illustrate what happens with a poorly-detected spectrum. In Section 2, independent fits to the SDSS spectrophotometry of the Balmer lines are presented.  In Section 3, we discuss the implications. ",
        "conclusions": ""
    },
    "0404/astro-ph0404128_arXiv.txt": {
        "abstract": "Variability studies represent a powerful tool to assess the emission mechanism(s) at work in the quiescent state of transient low mass X--ray binaries. Here we report on Chandra and XMM-Newton observations of two well known transient sources: Aql X-1 and Cen X-4. Long and short term variability is observed. This variability can be easily interpreted in view of an active millisecond radio pulsar emitting X--rays at the shock between a radio pulsar wind and inflowing matter from the companion star. ",
        "introduction": "Many Low Mass X--ray Binaries (LMXRBs) accrete matter at very high rates, and therefore shine as bright X--ray sources, only sporadically. Among these systems are Soft X--ray Transients (SXRTs) hosting an old neutron star (for a review see Campana et al. 1998). These systems alternate periods (weeks to months) of high X--ray luminosity, during which they share the same properties of persistent LMXRBs, to long (1--few 10 years) intervals of quiescence in which the X--ray luminosity drops by up to 5--6 orders of magnitude.  Short term variability is a powerful tool for the study of the emission mechanism(s) responsible for the SXRTs quiescent emission. A factor of 3 variability over timescales of days (Campana et al. 1997) and of $40\\%$ over 4.5 yr (Rutledge et al. 2001) has been reported in Cen X-4. Several other transient neutron star systems have also been found to be variable in quiescence by factors of 3--5 (e.g. Rutledge et al. 2000) but data have been collected over several years and with different instruments. Here we report on recent results obtained on two of the best studied sources of this class. ",
        "conclusions": "Variations in the temperature of the neutron star atmosphere might suggest that accretion onto the neutron star surface is occuring in quiescence; variations in the power law tail (photon index) together with variations in the column density should support the view of an active millisecond radio pulsar emitting X--rays at the shock between a radio pulsar wind and inflowing matter from the companion star. We tested both models with state of the art data on Aql X-1 (long term variability with Chandra) and Cen X-4 (short term variability with XMM-Newton). Despite current toughts either model can explain equally well the observational data."
    },
    "0404/hep-th0404109_arXiv.txt": {
        "abstract": "{A vacuum instability due to a massless ghost in a hidden sector can lead to an effective equation of state for dark energy that changes smoothly from $w=-3/2$ at large redshifts, to $w\\approx-1.2$ today, to $w=-1$ in the future. We discuss how this ghost can be the Goldstone boson of Lorentz symmetry breaking, and we find that this breaking in the hidden sector should occur at a scale below $\\sim$10 KeV. The normal particles that are produced along with the ghosts are then predominantly neutrinos.} \\begin{document} ",
        "introduction": "It would be disconcerting if any real undertanding of the current accelerated expansion of the universe was blocked by our profound ignorance of the cosmological constant. This motivates us to try to approach the dark energy problem in a manner quite unrelated to the cosmological constant problem. In particular we will consider a very weak vacuum instability as responsible for the origin of dark energy, which appears late in the evolution of the universe. The instability will arise through the appearance of an effective ghost, a degree of freedom with the wrong-sign kinetic term. The presence of a ghost implies that the vacuum is unstable to production of ghosts plus normal matter while conserving energy. One can view this simply as an out of equilibrium situation, where the system is trying to choose a more suitable configuration.  We would like to consider this apparent runaway behavior in more detail when its effect on cosmic evolution is also considered. We will only keep one degree of freedom of gravity, that of the scale factor $a(t)$ in the spatially flat FRW metric, $ds^2=dt^2-a(t)^2d\\mathbf{x}^2$. We will model the production of ghost and nonghost degrees of freedom as increasing densities of two perfect fluids, one ghost-like and one normal.  The respective equations of state are $p_g=w_g \\rho_g$ and $p_n=w_n \\rho_n$, where $p_g$ and $\\rho_g$ are negative and $p_n$ and $\\rho_n$ are positive. The normal fluid becomes a component of the total normal matter content of the universe, but we will continue to distinguish it from the rest of the normal matter.  If the ghost fluid was truly decoupled from everything else then it would simply dilute away with the expansion along with the normal matter. It is only when there is interaction causing nonconservation (production) of the ghost fluid can there be more dramatic effects. At the very least it can be expected that gravity itself will induce the interaction, and in particular a virtual graviton exchange can produce a pair of ghost particles and a pair of normal particles \\cite{cline}. We will also be led to consider a nongravitational interaction, but in either case we refer to the sector containing the ghost as a hidden sector.  We introduce a nonconservation of the two fluids so that $T_n^{\\mu\\nu}+T_g^{\\mu\\nu}$ is conserved, \\begin{equation} D_\\mu T_n^{\\mu\\nu}=-D_\\mu T_g^{\\mu\\nu}=\\varepsilon U^\\nu, \\label{e6}\\end{equation} where $U^\\nu$ is a velocity 4-vector and $\\varepsilon$ is a positive finite constant. $U^\\nu$ can be used to define the rest frame of the ghost fluid. But the question is how this preferred frame gets chosen in the first place, given that the underlying process is vacuum $\\rightarrow$ something. Since the vacuum of the normal sector is Lorentz invariant, we see already that there must be something intrinsically Lorentz violating about the vacuum state of the hidden sector.  We will assume that $U^\\nu$ is aligned with the cosmological rest frame. In this frame we assume a constant ghost particle production rate per unit volume $\\Gamma_g$ and an average (negative) energy per ghost particle $\\overline{E}_g$. Then $\\varepsilon=-\\overline{E}_g\\Gamma_g$, with a similar relation for the normal particles. Equations (\\ref{e6}) along with the Einstein equations yield interesting solutions for $\\delta w\\equiv w_g-w_n > 0$. The $\\delta w < 0$ case leads quickly to curvature singularities, while the $\\delta w = 0$ case is not interesting for gravity. There, the solution is simply $a(t)=1$ and $\\rho_n=-\\rho_g=\\rho_0+\\varepsilon t$ in the absence of other matter. ",
        "conclusions": "A consistent infrared modification to gravity involving Lorentz symmetry breaking results in some kind of instability. Very schematically, the mixing between the Goldstone boson $\\pi(x)$ and the nonpropagating scalar mode of gravity $\\Phi(x)$ is given by \\cite{nima} \\begin{equation} {\\cal L}=-{1\\over2}(\\nabla\\Phi)^2\\pm{1\\over2}(m\\Phi-\\dot{\\pi})^2+... \\end{equation} with $m=\\Lambda_L^2/\\sqrt{2}m_\\mathrm{Pl}$. The choice of the positive sign in this Lagrangian ensures that the Goldstone boson has positive kinetic energy, but it gives the nonstandard sign for the mass term of $\\Phi$. This nonstandard sign and the ensuing mixing leads to a classical instability, yielding the exponential growth of modes characterized by a length scale $m_\\mathrm{Pl}/\\Lambda_L^2$ and a time scale $m_\\mathrm{Pl}^2/\\Lambda_L^3$. In this paper we have instead chosen the minus sign and thus a Goldstone ghost, which does not suffer from this particular instability. There are still gravitational implications of the mixing, and they are still characterized by the same length and time scales as before, but they are now well behaved.\\footnote{At the end of section 3 we also briefly noted that the classical Goldstone ghost field can directly yield a $w<-1$ dark energy component that evolves towards $w=-1$ in the future.}  The Goldstone ghost gives rise to a different instability, that of ghost production from the vacuum. This has a finite rate because of the cutoff provided by the scale of Lorentz symmetry breaking. If it is graviton exchange that is responsible for the instability then the time scale for its effects to show up is again $m_\\mathrm{Pl}^2/\\Lambda_L^3$. It appears that the instability of the Goldstone ghost is of particular cosmological interest, since it pushes the scale factor of the whole universe towards exponential growth. The constant particle production from the vacuum balances against the dilution from the exponential growth, and thus the endpoint of this ``instability'' is well behaved. This picture is quite independent of the remaining question of why the cosmological constant vanishes.  If ghost production is the cause of the present accelerated expansion, then the universe is only part way in its transition to the steady state behavior. The effective equation of state parameter of the produced ghosts and normal particles evolves from $-3/2$ to $-1$ and is presently $w_\\mathrm{eff}\\approx -1.2$. The ghosts are basically invisible, but the normal particles produced lead to observational constraints. In fact the normal particles should be predominantly neutrinos rather than photons, and for this reason we considered a nongravitational coupling between the hidden sector and our own. In the case of a broken gauge interaction, perhaps connected to flavor physics around 1000 TeV, we showed that adequate suppression of photons relative to neutrinos can naturally occur. The energy spectrum of neutrinos will peak around the scale of Lorentz symmetry breaking, $\\Lambda_L$, which must be below 10 KeV.  $\\Lambda_L$ sets an upper bound on the energy that a propagating ghost can carry. Above this scale there is some ultraviolet completion of the infinite derivative scalar field model. We speculate that this more fundamental theory in the hidden sector could contain stable massive states with positive energy. This would have to be consistent with having a sufficiently stable ground state in the hidden sector (when considered in isolation). This may be possible as long as there is sufficient separation between the minimum mass of these states and the maximum energy (in absolute value) that can be carried by the Goldstone ghosts. Positive energy in the hidden sector may then end up in these massive states as the universe cools, and we note that the mass and number density of these states may make them of interest for dark matter.  Finally we note that the proximity of $\\Lambda_L$ to the scale of neutrino masses may suggest a common origin for these mass scales, perhaps related to the physics of flavor."
    },
    "0404/astro-ph0404152_arXiv.txt": {
        "abstract": "We have performed 2.5-dimensional general relativistic magnetohydrodynamic (MHD) simulations of collapsars including a rotating black hole. This paper is an extension of our previous paper (Mizuno et al. 2004). The current calculation focuses on the effect of black hole rotation using general relativistic MHD with simplified microphysics, i.e., we ignore neutrino cooling, physical equation of state and photodisintegration. Initially, we assume that the core collapse is failed in this star. A few $M_{\\odot}$ rotating black hole is inserted by hand into the calculation. We consider two cases, the co-rotating case and counter-rotating case with respect to the black hole rotation. Although the counter-rotating case may be unrealistic for collapsar, we perform as a maximally dragging case of magnetic field. The simulation results show the formation of a disk-like structure and the generation of a jet-like outflow near the central black hole. The jet-like outflow propagates outwardly with the twisted magnetic field and becomes collimated. We have found that the jets are generated and accelerated mainly by the magnetic field. The total jet velocity in the rotating black hole case is comparable to that of the non-rotating black hole case (Mizuno et al. 2004), $\\sim$ 0.3c. When the rotation of the black hole is faster, the magnetic field is twisted strongly owing to the frame-dragging effect. The magnetic energy stored by the twisting magnetic field is converted to kinetic energy of the jet directly rather than propagating as an Alfv\\'{e}n wave. Thus, as the rotation of the black hole becomes faster, the poloidal velocity of the jet becomes faster. In the rapidly rotating black hole case the jet-like outflow can be produced by the frame dragging effect only through the twisting of magnetic field even if there is no stellar rotation. ",
        "introduction": "It is now generally believed that ``long-soft'' gamma-ray bursts (GRBs) are a phenomenon related to the deaths of massive stars. There are direct and indirect observational evidence of the close relationship between GRBs and supernovae, such as the observed association with star-forming regions in galaxies (Vreeswijk et al. 2001; Bloom, Kulkarni, \\& Djorgovski 2002; Gorosabel et al. 2003), the ``bump'' observed in the afterglows of some GRBs (Reichert 1999; Galama et al. 2000; Bloom et al. 2002; Garnavich et al. 2003; Bloom et al. 2003), metal emission lines observed in the X-ray afterglow of GRB011211 (Reeves et al. 2002), and  the association of GRB980425 with SN 1998bw (Galama et al. 1998b). Recently, another clear evidence was found, GRB030329 was accompanied by a bright energetic supernova of Type Ic, SN 2003dh (Price et al. 2003b; Hjorth et al. 2003; Stanek et al. 2003). It has been suggested so far that some of GRBs are produced when the iron core of a massive star collapses either to a black hole (Woolsey 1993; MacFadyen \\& Woosley 1999) or to a rapidly rotating highly magnetic neutron star (Wheeler et al. 2000), both of which eventually produce a relativistic jet.  The collapsar model is one of the most promising scenarios involving massive stars (Woosley 1993; MacFadyen \\& Woosley 1999). Collapsar is a rotating massive star, lacking a hydrogen envelope. In this model, the iron core of the rotating massive star collapses to a black hole surrounded by an accretion disk. The accretion through this disk produces outflows via neutrino annihilation and/or magnetohydrorynamic (MHD) processes. They are further collimated by the passage through the stellar mantle.  The formation and propagation of relativistic flows from collapsar has been studied numerically by both Newtonian (MacFadyen \\& Woolsey 1999; MacFadyen, Woosley, \\& Heger 2001) and special relativistic hydrodynamic simulations (Aloy et al. 2000; Zhang, Woosley, \\& MacFadyen 2003). However, these previous numerical simulations of the collapsar model did not fully address the formation mechanism of the relativistic outflow.  It is suspected that magnetic fields may play an important role in the formation and acceleration of relativistic jets. In numerical simulations of gravitational collapse of massive stars the effect of stellar rotation and intrinsic magnetic fields was studied by several authors (LeBlanc \\& Wilson 1970; Symbalisty 1984; Ardeljan et al. 2000; Kotake et al. 2004; Yamada et al. 2004). Symbalisty (1984) showed the formation of high-density, supersonic jets in the combination of a rapid rotation and a strong dipole magnetic field, which was confirmed by Yamada et al. (2004).  In our previous paper (Mizuno et al. 2004) we have performed 2.5-dimensional general relativistic MHD simulations of the gravitational collapse of a rotating magnetized massive star with a non-rotating black hole at the center as a model for collapsar. We showed the formation of a disk-like structure and the generation of a mildly relativistic jet ($\\sim 0.3 c$) inside the shock wave launched at the core bounce. We have found the jet is accelerated by the magnetic pressure and the centrifugal force and is collimated by the pinching force of the toroidal magnetic field amplified by the rotation and the effect of geometry of the poloidal magnetic field. The Poynting flux transports more energy outward than the jet. However, the jet in these simulations is too slow to applicable for GRBs. In our previous paper, we put a non-rotating black hole at the center, which is not appropriate for the rotational collapse. It is natural that the iron core collapses to a rotating black hole. The rotating black hole (Kerr black hole) is useful to form the relativistic jet for some reasons. If the black hole is rotating, two kinds of energy are available for generating a jet. One is the rotational energy of the accreted matter and the other is the rotational energy of the rotating black hole itself. Blandford and Znajek (1977) investigated magnetospheres of Kerr black holes and derived a force-free solution for the electromagnetic field. Their results showed that electromagnetic energy is radiated from the black hole horizon directly. This is called Blandford-Znajek mechanism. However, the direct energy emission from the black hole horizon appears to be inconsistent with the causality at the horizon.  Punsly and Coroniti (1990) found that if the ergospheric plasma gets frozen onto large-scale magnetic field lines, it can drive a magnetic wind to infinity. Using a general relativistic MHD code, the basic mechanism of energy extraction from the Kerr black hole via magnetic field is investigated numerically (Koide et al. 2002; Koide 2003). The numerical results showed that the rotational energy of the Kerr black hole can be extracted when the magnetic field is strong enough.  Here, as a collapsar model we perform 2.5-dimensional general relativistic MHD simulations of the gravitational collapse of a rotating magnetized massive star with a rotating black hole at the center. We investigate the physics of the formation of jets, the acceleration force on the jets, and the dependence on the rotation parameter of the black hole.  This paper is extension of our previous paper (Mizuno et al. 2004). We describe the numerical method of our simulations briefly in section 2 and present our results in section 3. The summary and discussion are given in section 4. ",
        "conclusions": "We have performed general relativistic MHD simulations of collapsar, paying particular attention to the rotation of the central black hole. We considered not only the co-rotating but also the counter-rotating cases with respect to the black hole rotation to elucidate its effect on the dynamics. Our results are summarized as follows:  \\begin{enumerate} \\item The formation mechanism of the jet for the rotating black hole is the same as that for the non-rotating black hole. When stellar matter falls onto the black hole, a disk-like structure is formed in the vicinity of the black hole horizon and a jet-like outflow is formed near the black hole mainly by the magnetic field. The jet-like outflow propagates outward with twisted magnetic fields and becomes collimated. The shock waves (MHD fast shock and slow shock) are also formed near the black hole and propagate outward. Figure {\\ref{fig11}} shows the schematic picture of our simulation results.  \\item The total velocity of jet for the rotating black hole is comparable to that for the non-rotating black hole, $\\sim 0.3c$. However, the poloidal velocity of the jet in the case of rotating black hole is about 3 times higher than that in the case of non-rotating black hole. The poloidal velocity is the dominant component of the jet. Because the magnetic field is more strongly twisted due to the frame dragging effect of the rotating black hole, more magnetic energy is converted to the kinetic energy of the jet rather than to the propagation of Alfv\\'{e}n waves.  \\item In the co-rotating case, the kinetic energy flux is comparable to the Poynting flux. In the counter-rotating case, on the other hand, the kinetic energy flux is about twice as large as the Poynting flux. This is because the jet for the rotating black hole has a higher density.  \\item As the rotation parameter of the black hole increases, the poloidal velocity of the jet and the magnetic twist increase gradually and the toroidal velocity of the jet decreases. These results are related to how effective the frame-dragging is. As the rotation of black hole is faster, the magnetic field is twisted more strongly owing to the frame-dragging effect. The magnetic energy stored by the twisted magnetic field is converted to the kinetic energy of the jet directly rather than leading to the propagation of Alfv\\'{e}n waves. Thus, the poloidal velocity of the jet becomes faster as the rotation of the black hole becomes faster.  \\item In rapidly rotating black hole case the jet-like outflow can be formed by the frame dragging effect through the twisting of magnetic field even if there is no stellar rotation.  \\end{enumerate}  We discuss the application of our results to the central engine of GRBs. Using the simulation data, we can estimate the kinetic energy of the jet, $E_{\\mathrm{jet}}$. If the mass of the disk formed in the initial accretion phase is comparable to the central black hole mass, e.g., several solar masses (almost all stellar matter except the central core are actually blown away by the presupernova), the density of the disk can be estimated from \\begin{equation} \\pi r_{\\mathrm{d}}^{2} H \\rho_{\\mathrm{d}} = 4 M_{\\odot} \\simeq 10^{34} \\mathrm{g} \\label{}, \\end{equation} where $r_{\\mathrm{d}}$ is the outer radius, $H$ the thickness and $\\rho_{\\mathrm{d}}$ the density of the disk, respectively. $r_{\\mathrm{d}}$ and $H$ can be estimated from the simulation data as $r_{\\mathrm{d}} = 10 r_{\\mathrm{S}} \\simeq 10^{7} \\mathrm{cm}$, $H = 0.1 r_{\\mathrm{d}} = 1 r_{\\mathrm{S}} \\simeq 10^{6} \\mathrm{cm} $, respectively. Therefore, the density of the disk is $ \\rho_{\\mathrm{d}} \\simeq 3 \\times 10^{13} \\mathrm{g/cm^{3}} $. Because the jet-like outflow for the rotating black hole is generated very close to the central black hole, the density of the jet is comparable to the density of the disk. Thus, using the estimated density of the disk, the kinetic energy of the jet is given as $E_{\\mathrm{jet}} \\simeq 10^{54}$ ergs. Although it may somewhat overestimate, this is large enough energy to explain the standard energy ($\\sim 10^{51}$ ergs) of GRBs . However, the maximum jet velocity in our simulations including a rotating black hole was about $0.3 c$. This velocity is too slow compared to the velocity inferred for the jet of GRBs. We have to consider other acceleration mechanisms. This is the most difficult and fundamental problem in all types of ultrarelativistic outflow, e.g. AGN jets, pulsar winds and GRBs. In order to obtain not only a large energy extraction but also the observed large bulk Lorentz factors, the Poynting flux must be converted to the kinetic energy. From steady solution of relativistic outflow in ideal MHD, it is found that  the energy conversion from the Poynting flux to the kinetic energy flux occurs and the outflow is highly accelerated, if the magnetic field lines diverge with radius more quickly than in the monopole field (Begelman \\& Li 1994; Takahashi \\& Shibata 1998; Daigne \\& Drenkhahn 2002). However, this solution is not self-consistent because the geometry of the magnetic field is not solved. Moreover, the relativistic outflow may not maintain the collimated structure in this situation. Some authors proposed a dissipation-induced acceleration mechanism (Spruit, Daigne, \\& Drenkhahn 2001; Drenkhahn 2002; Drenkhahn \\& Spruit 2002; Sikora et al. 2003). Although the centrifugal acceleration may also contribute to the early stages of acceleration of the flow, this is not a magneto-centrifugal acceleration process. If the magnetic fields in the outflow change the direction on sufficiently small scales, a part of the magnetic energy can be released locally by magnetic reconnections. This process has two useful effects. First, it converts Poynting flux directly into radiation, without only intermediate step with internal shocks. Second, it leads to steeper decline of magnetic pressure, which causes great acceleration of the flow and enhanced conversion of the Poynting flux to the kinetic energy. Since in the co-rotating case of our simulations, the Poynting flux is comparable to the kinetic energy flux, the dissipation-induced acceleration mechanism which might occur later on may be the key to solve the acceleration problem of our simulations.  Although it may not be a main solution of the acceleration problem, we further discuss possible acceleration mechanisms from other simulations of relativistic jets. The first possibility is a disk-jet, especially a gas-pressure-driven jet. The gas-pressure-driven jet was seen in general relativistic MHD simulations of the black hole-accretion disk system (Koide et al. 1998, 1999, 2000; Aoki et al. 2004) and has higher velocity than the magnetically driven jet. It is generated by the shock produced by the centrifugal barrier in the equatorial plane. Since, in our simulations, we initially assumed that the rotation is almost rigid, the centrifugal barrier is not so effective in the initial accretion phase. The accreted matter falls to the central black hole smoothly without the generation of a shock. We expect that if our simulations run long enough to form the accretion disk around the central black hole, the gas-pressure-driven jet will be generated. The second possible mechanism is the break out of the jet from the stellar surface. When the jet emerges from the stellar surface, the steep density gradient accelerates the jet. Some authors (Aloy et al. 2000; Zhang et al. 2003) have showed numerically that the jet is accelerated significantly by this process, and the terminal Lorentz factor becomes more than 100. Since the jet in our simulations contains large energy, this mechanism may work effectively. We will address the propagation of the jet outside the stellar surface in a forthcoming paper.  Our results can be also applied to baryon-rich outflows associated with failed GRBs. The baryon-rich outflow is a fireball with high baryonic load and mildly relativistic velocities. The jet velocity is so slow that it cannot produce GRBs. Such failed GRBs are supposed to occur at higher event rates than successful GRBs (Woosley et al. 2002; Huang, Dai, \\& Lu 2002).  Some failed GRBs may be observed as ``hypernovae''. Particularly, SN 2002ap is a candidate for such an event. Although the association with a GRB has not been found , it is inferred to have a jet with a velocity of $\\sim 0.23 c$ and estimated kinetic energy of $\\sim 5 \\times 10^{50} \\mathrm{ergs}$ (Kawabata et al. 2002; Totani 2003). The jet found in our simulations for rotating black holes has comparable velocity and large enough energy to explain the jet of SN2002ap.  In our simulations, we have neglected microphysics (photodisintegration, treatment of neutrinos etc.) and concentrated on the general relativistic magnetohydrodynamics, particularly the effect of rotating black hole.  Such treatment may be reasonable as a first step. However, microphysics are critically important for realistic GRMHD simulations of collapsar.  We will address these issues in a forthcoming paper.  We have assumed a uniform global magnetic field in the simulations. It is inferred, however, that rotating compact stars in general have a dipole-like magnetic field. Hence, it may be more likely that the rotating stars collapse with a dipole-like magnetic field. Although we think that the assumption of a uniform global magnetic field will not be so bad for the discussion of the generic aspect of collapsars dynamics, the dependence on the configuration of magnetic field is also currently being investigated and will be presented elsewhere.   \\medskip  Y. M. appreciates many helpful conversations on GRBs and relativistic effects with S. Aoki, T. Totani, R. Yamazaki, M. Takahashi, and K. Nishikawa. He also thanks K. Uehara, H. Kigure, T. Haugb\\o elle. and Y. Kato. This work was partially supported by Japan Science and Technology Cooperation (ACT-JST), Grant-in-Aid for the 21st Century COE ``Center for Diversity and Universality in Physics'' and Grants-in-Aid for the scientific research from the Ministry of Education, Science, Sports, Technology, and Culture of Japan through No.14079202, No.14540226 (PI: K. Shibata) and No.14740166. The numerical computations were partly carried out on the VPP5000 at the Astronomical Data Analysis Center of the National Astronomical Observatory, Japan (yym17b), and partly on the Alpha Server ES40 at the Yukawa Institute for Theoretical Physics of Kyoto University, Japan."
    },
    "0404/astro-ph0404478_arXiv.txt": {
        "abstract": "{We present spectral and timing results from a long (130~ks) \\xmm\\ EPIC observation of the nucleus of the Seyfert/LINER galaxy M\\,81. During the observation the X-ray flux varied by 20\\%, but there was no significant change in spectral shape. The 2--10 keV spectrum is well described by a power law continuum and three narrow Fe~K$\\alpha$ emission lines at 6.4, 6.7 and 6.96 keV. The three emission lines have equivalent widths of $39^{+13}_{-12}$, $47^{+13}_{-13}$ and $37^{+15}_{-15}$~eV respectively. The ratios of the three lines are thus more similar to those observed from the Galactic Centre region than to those typically observed from Seyfert galaxies. The high ionization lines most likely originate either from photoionized gas within 0.1 pc of the nucleus of M~81, or from a non-thermal distribution of cosmic-ray electrons interacting with the 0.2-0.6 keV thermal plasma which is found in the bulge of M~81. ",
        "introduction": "At a distance of 3.6 Mpc (Freedman et al. 1994), the Sab spiral M\\,81 is the closest galaxy showing the spectral signatures of a ``low ionization nuclear emission-line region'' (LINER). LINER characteristics are found in a significant fraction of all galaxies (between 1/5 and 1/3), and the nature of the relationship between LINERs and low-luminosity AGN (LLAGN) has been a long standing question (Ho, Filippenko \\& Sargent \\nocite{ho97} 1997, Heckman \\nocite{heckman80} 1980). The presence of a LLAGN in M\\,81 is well established: the unresolved nuclear source emits a non-stellar UV continuum, broad optical/UV emission lines, and a significant flux of non-thermal X--rays (Ho, Filippenko \\& Sargent \\nocite{ho96} 1996).  From dynamical studies, the mass of the central compact object is found to be $\\approx 4 \\times 10^6\\ {\\rm M}_{\\odot}$ (Ho \\nocite{ho99} 1999). With a bolometric luminosity of $\\sim 10^{41}$~erg~s$^{-1}$, M\\,81 is radiating at $< 10^{-3}$ of its Eddington luminosity (Ho, Filippenko \\& Sargent \\nocite{ho96} 1996).  The nuclear X--ray source has a power law spectral shape in the 0.3-10 keV energy range (Petre et al. \\nocite{petre93} 1993, Ishisaki et al. \\nocite{ishisaki96} 1996, Swartz et al. \\nocite{swartz03} 2003). The spectrum is absorbed by a small column density ($\\sim 3\\times 10^{20}$~cm$^{-2}$) of material intrinsic to M\\,81, and below 2 keV a multitude of emission lines from circumnuclear, thermal gas can be seen in the \\xmm\\ RGS spectrum of the nucleus (Page et al. \\nocite{page03a} 2003).  Above 2 keV, the most remarkable feature seen in the spectrum is an emission line at $6.7\\pm0.1$ keV, corresponding to Fe~XXV~K$\\alpha$ (Ishisaki et al. \\nocite{ishisaki96} 1996, Pellegrini et al. \\nocite{pellegrini00} 2000). Using data from {\\em BeppoSAX}, Pellegrini et al. \\nocite{pellegrini00} (2000) claim that there is an absorption edge at $\\sim 8.6$ keV, which could also arise in Fe~XXV, indicating that the nucleus is seen through a large column density ($\\sim 2\\times 10^{23}$~cm$^{-2}$) of highly ionized gas.  The 2-10 keV luminosity is known to vary in the range $(1.5-6)\\times 10^{40}$~erg~s$^{-1}$ on a timescale of years, and variations of about 30\\% have been observed on a timescale of $10^{5}$~s (Pellegrini et al. \\nocite{pellegrini00} 2000; Ishisaki et al. \\nocite{ishisaki96} 1996).  In this paper we examine the time variation and the 2-10 keV spectrum of the M\\,81 nucleus using the \\xmm\\ European Photon Imaging Camera (EPIC) data from our long (130~ks) observation which was carried out as part of the RGS consortium guaranteed time programme. ",
        "conclusions": "The 2-10 keV X-ray spectrum of the nucleus of M\\,81 consists of a power law continuum and three Fe~K$\\alpha$ emission lines at 6.4, 6.7 and 6.96 keV.  The power law component has a photon index $\\Gamma \\sim 1.9$, consistent with the continuum observed below 2 keV with the RGS. The X-ray flux varied by 20\\% during the 130~ks observation but no gross spectral variations were detected. The spectrum does not require a reflection component, and shows no evidence for absorption edges due to highly ionized Fe. The equivalent widths of the 6.4, 6.7 and 6.96 keV emission lines are $39^{+13}_{-12}$, $47^{+13}_{-13}$ and $37^{+15}_{-15}$~eV respectively.  These equivalent widths are significantly different to those normally observed in Seyfert galaxies, which typically have stronger 6.4 keV lines and weaker 6.7 and 6.96 keV lines.  However, the ratios of the three lines in M~81 are relatively similar to the those observed from the Galactic Centre region.  The nucleus of M~81 does not appear to contain as much cold, X--ray illuminated material as most Seyfert galaxies. The 6.7 and 6.96 keV emission lines might originate in photoionized gas within 0.1~pc of the nuclear continuum source. Alternatively, these emission lines might have a similar origin to the Galactic Centre Fe line emission, possibly through the interaction of non-thermal cosmic ray electrons with the 0.2-0.6 keV thermal plasma component of the M\\,81 bulge."
    },
    "0404/astro-ph0404364_arXiv.txt": {
        "abstract": "We report the discovery of a strong iron K$\\alpha$ line in the hard X--ray selected source {\\tt CXOJ 123716.7+621733} in the {\\it Chandra} Deep Field North survey at $z=$1.146. The analysis is made possible by the very deep exposure $\\sim$ 2 Ms and low background of the {\\tt ACIS} detector. The line profile seems to be inconsistent with a narrow feature. The best fit solution is achieved with a broad line. Most of the flux in the broad component originates at energies below 6.4 keV with a  shape similar to that expected from emission in the innermost regions of the accretion disk. ",
        "introduction": "The strongest emission feature in the hard X--ray spectrum of an active galactic nucleus (AGN) is the fluorescent FeK$\\alpha$ emission line at 6.4--7 keV. Since its discovery by early X--ray observations, it is by now recognized to be an ubiquitous feature in the high energy spectra of Seyfert galaxies. The line properties -- in particular the centroid energy, intensity and profile -- carry important diagnostic information about the dynamics and physics of the region where the emission originates. For example, the best fit line energy is a measure of the ionization status of the gas, the line equivalent width (EW) is strongly correlated with the amount of fluorescing material, while the detection of a broad asymmetric profile, first resolved with ASCA by Tanaka et al. (1995) in the bright Seyfert 1 galaxy MCG--6--30--15, is considered the most direct evidence of the presence of an accretion disc extending down to a few gravitational radii of the central black hole. The power of X--ray spectroscopy as a diagnostic of the AGN physics is witnessed by hundreds of papers (see Fabian et al. 2000; Reynolds \\& Nowak 2003, for extensive reviews).  The detailed study of the iron line properties requires a large collecting area and good energy resolution. For this reason, most of the results obtained so far are limited to nearby, low luminosity Seyfert galaxies (e.g. Turner et al. 2002; Fabian et al. 2002; Dewangan, Griffiths \\& Schurch 2003). X--ray observations of sizable samples of higher luminosity AGN, extending in the quasars regime ($L_X > 10^{44}$ erg s$^{-1}$), reveal the evidence of a trend whereby the strength of the iron line decreases and the centroid moves to higher energies (Iwasawa \\& Taniguchi 1993, Nandra et al. 1997a, George et al. 2000, Reeves \\& Turner 2000, Page et al. 2004). The search for iron line emission at high redshift is hampered by the statistical quality of the X--ray spectra and only a few tentative detections at the 2$\\sigma$ level have been reported. (e.g., Vignali et al. 1999; Reeves et al. 2001; Norman et al. 2002; Gandhi et al. 2004).  Deep, sensitive X--ray observations carried out with both {\\it Chandra} and XMM--{\\it Newton} now offer the opportunity to search for the presence of iron emission line with relatively large well defined samples of X--ray selected sources, down to limiting fluxes which provides the possibility to probe a much larger redshift range. Motivated by this possibility, we have started a systematic search for emission line features among the serendipitous sources in the deepest X--ray observations available to date.  Here we report the detection of a broad iron line in the X--ray source {\\tt CXOJ 123716.7+621733}, serendipitously detected in the 2 Ms exposure of the {\\it Chandra} Deep Field--North (CDF--N; Alexander et al. 2003a). The spectroscopic redshift ($z=1.146$) is based on a single {\\tt [OII]$\\lambda$3727} line (Barger et al. 2002). The compact optical morphology, the hard X--ray spectrum and the presence of a radio source coincident with the X--ray position (Richards 2000) leaves no doubt about the AGN nature of this object. On the basis of the available optical spectrum it is not possible to determine if the object is a Type 1 or a Type 2 AGN. Unfortunately, at the expected {\\tt Mg II$\\lambda$2798} position the optical spectrum is contaminated by strong atmospheric bands and night-sky emission lines. The source is classified as radio--quiet on the basis of the 1.4 GHz flux density (0.346 mJy) and optical magnitude (Barger et al. 2003). A cosmological model with $H_0$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m$ = 0.3 and $\\Omega_{\\Lambda}= 0.7$ is assumed. ",
        "conclusions": "A strong emission feature due to the K$\\alpha$ iron line has been clearly detected in the X--ray spectrum of the relatively bright, hard X--ray source {\\tt CXOJ 123716.7+621733} at $z$=1.146 in the {\\it Chandra} Deep Field North. The broad band 0.6--7 keV {\\it Chandra} spectrum is best described by a relatively hard ($\\Gamma \\simeq$ 1.5) power law plus significant intrinsic absorption $N_H \\sim 1.8 \\times$ 10$^{22}$ cm$^{-2}$. An almost equally good fit is obtained if the continuum emission is parameterized with a steeper ($\\Gamma$ = 1.8) slope plus a reflection component as commonly observed in nearby Seyfert galaxies.  The line profile is not consistent with a narrow feature, being significantly broader than the {\\tt ACIS-I CCD} resolution at that energy. Such a result does not depend on the adopted shape of the underlying continuum.  In order to reproduce both the observed width and the optical spectroscopic redshift, most of the line flux has to originate in a red wing. An adequate description of the overall line shape has been obtained with a relativistic line model. Although the quality of the data is not such to break the degeneracy between the various line parameters, the present observation suggests an almost face--on orientation of the accretion disk.  The line equivalent width (860$\\pm$260 eV) is larger than the average values measured by {\\tt ASCA} for a sample of nearby bright Seyfert galaxies. Nandra et al (1997b) report an average value of 230$\\pm$60 eV, for a relativistic line model. Plausible possibilities include iron overabundance as in MCG--6--30--15 (see e.g. Lee et al. 1999) and/or an enhanced contribution of the reflected flux possibly associated to time variability. The latter possibility appears unlikely given the lack of significant flux variability over more than two years. It must be noted that large values of the EW are not unusual among relatively faint X--ray sources (George et al. 2000) and likely to be due to a selection effect associated to the uncertainties in modelling the continuum and line spectrum.  The broad band properties of {\\tt CXOJ 123716.7+621733} (namely, the high X--ray to optical ratio, (log$f_X/f_{opt} \\simeq$ 1), the X--ray column density ($>$ 10$^{22}$ cm$^{-2}$) and the optical, near--infrared colors ($B-V=0.7$, $R-K=4.4$; Barger et al. 2003)) indicate substantial obscuration of the nuclear source. There is also no evidence of typical AGN emission lines in the low resolution optical spectrum (Barger et al. 2002).  The present results indicate that the iron K$\\alpha$ line could be the strongest AGN signature in the broad band spectra (from optical to X--rays) of distant, obscured AGN for which only a low signal--to--noise optical spectrum covering a relatively small range of wavelengths is available, and highlight the uniqueness of X--ray observations in the study of supermassive black holes at cosmologically interesting distances.  The study of relativistically broadened iron line profiles would greatly benefit from deep X--ray observations with the XMM--{\\it Newton} large collecting area telescopes. Observations with the new generation of X--ray observatories ({\\tt Constellation X} and {\\tt XEUS}) will allow systematic studies of iron line properties at high redshifts."
    },
    "0404/gr-qc0404053_arXiv.txt": {
        "abstract": "The consequences of taking the generalized Chaplygin gas as the dark energy constituent of the Universe on the gravitational waves are studied and the spectrum obtained from this model, for the flat case, is analyzed. Besides its importance for the study of the primordial Universe, the gravitational waves represent an additional perspective (besides the CMB temperature and polarization anisotropies) to evaluate the consistence of the different dark energy models and establish better constraints to their parameters. The analysis presented here takes this fact into consideration to open one more perspective of verification of the generalized Chapligin gas model applicability. Nine particular cases are compared: one where no dark energy is present; two that simulate the $\\Lambda$-CDM model; two where the gas acts like the traditional Chaplygin gas; and four where the dark energy is the generalized Chaplygin gas. The different spectra permit to distinguish the $\\Lambda$-CDM and the Chaplygin gas scenarios. \\vspace{0.7cm}  \\par KEYWORDS: dark energy, dark matter, gravitational waves \\par \\vspace{0.7cm} \\par PACS numbers: 98.80.Bp, 98.65.Dx ",
        "introduction": "\\label{intro} The most recent results from type Ia supernovae observations \\cite{riess,permultter,tonry} and the cosmic microwave background anisotropies detection \\cite{wmap} have led cosmologists and astrophysicists to the conclusion that most of the matter of the Universe interacts in a repulsive manner (as an example through a negative pressure). There are many candidates to describe this exotic fluid. The most natural candidate is a cosmological constant, but it presents a discrepancy of $120$ order of magnitude between the theoretical predictions and the observational data \\cite{cc}. A self-interacting scalar field, known as quintessence, is another proposal to explain dark energy \\cite{quint1}. However, it asks for fine tuning of microphysical parameters in order to have a suitable potential term \\cite{quint2}. Among many other possibilities, the Chaplygin gas model has receveid recently special attention \\cite{pasquier,berto,fabris}. The Chaplygin gas model is based on a perfect fluid whose pressure, besides to be negative, varies inversely with density. One of the interesting aspects of this fluid is a connection with branes in the context of string theories \\cite{hoppe,ogawa,jackiw}. Some phenomenological generalizations of this fluid have been proposed, leading to the so-called {\\it generalized Chaplygin gas model} \\cite{berto}. \\par The aim of this paper is to investigate the particular signatures of the generalized Chaplygin gas, mainly in comparison to the standard $\\Lambda$-CDM, in what concerns gravitational waves. Special attention will be given to the spectral distribution of energy density. \\par The gravitational waves are very important to Cosmology. They must have special signatures in the polarization \\cite{paul1} of the CMB anisotropies \\cite{paul2}. Moreover, since gravitational waves has decoupled from matter already in the deep early Universe, they can be a window to the primordial phase. Even if gravitational waves (and consequently the polarization of the CMB photons \\footnote{There are two polarizations modes, called $E$ and $B$. The mode $E$ has already been identified, but only the detection of the mode $B$ will allow to identify gravitational waves \\cite{kogut,kovac}.}) have not been detected directly until now, great efforts are been done in this sense and there is a hope that the next generation of experiment in space (LISA, Planck, etc.), may allow this detection in perhaps ten years. The GW spectrum frequency range of observational interest extends from $10^{-18}Hz$ to $10^{10}Hz$ and the energy density spectrum is constrained by the CMB, in units of the critical density $\\Omega_{GW}$, as \\cite{uzan} \\begin{equation} \\label{constr} \\frac{d\\Omega_{GW}}{d \\ln \\nu}\\Big|_{10^{-18}} \\leq 10^{-10} \\quad . \\end{equation} This constraint will be used later in the evaluation of the spectra. Many works have been made in order to identify specific signatures of cosmological models in the spectra of gravitational waves, for example, in the case of quintessence model \\cite{uzan} and string cosmology \\cite{ungarelli,sanchez,gasperini}. \\par The generalized Chaplygin gas is characterized by the equation of state \\begin{equation} \\label{stateeq} p=-\\frac{A}{\\rho^\\alpha}\\quad , \\end{equation} where $A$ and $\\alpha$ are constants such that $A \\geq 0$ and $0 \\leq \\alpha \\leq 1$. If the energy-momentum tensor conservation is taken into account, the relation between the generalized Chaplygin gas density and the scale factor $a$ becomes \\begin{equation} \\label{density} \\rho = \\left(A+\\frac{B}{a^{3(\\alpha + 1)}} \\right)^{\\frac{1}{\\alpha+1}}\\quad , \\end{equation} where $B$ is an arbitrary integration constant. The main properties of the relation (\\ref{density}) are: (i) it interpolates the cosmological constant phase (when the scale factor is large) and a pressureless fluid (for small values of $a$) phase; (ii) the sound velocity associated with it goes from zero, for $A=0$, to the velocity of light, being always positive.  We will analyze the gravitational waves spectrum obtained from the generalized Chaplygin gas model, which is described in section \\ref{section_model}. In section \\ref{section_h} we discuss general properties of the gravitational waves differential equation in the context of two important dark matter models: cosmological constant and Chaplygin gas. The spectrum obtained from numerical calculations on these cases are presented in section \\ref{section_spectrum} and analyzed in section \\ref{conclusions} where we also present our conclusions. ",
        "conclusions": "\\label{conclusions}  In this work we have studied the fate of gravitational waves in the context of cosmological models where dark energy is described by the generalized Chaplygin gas. We have exploited the cases where dark energy and dark matter are unified by the Chaplygin gas, as well as the cases where a pressureless dark component is present besides the Chaplygin gas. In all situations, however, we have taken into account the baryonic component as infered from the primordial nucleosynthesis. We perform also an analysis, at same time analytical and numerical, of the $\\Lambda$-CDM model and of a pure dark matter model (where pressureless matter is the only component of the Universe). In this sense, the goal of the present work was to try to identify special signature of each model in the energy spectrum as function of the frequency. \\par The spectra for all models exhibit the same shape: initially the energy decreases with frequency, becomes almost constant, and then it increases with frequency. However, from the results it comes out that the $\\Lambda$-CDM model presents a spectra with an amplitude greater then the pure dark matter model by a factor of $6$. The Chaplygin gas model interpolates these two cases. In a typical situation, for example when $\\alpha = 1$, $\\bar A = 0.5$ and there are only Chaplygin bas and baryons, the amplitude of the spectra is about $4$ times smaller than in the $\\Lambda$-CDM case. As the $\\Lambda$-CDM case is approached ($\\bar A \\rightarrow 1$), the amplitude increases, but in a very low rate, exceptly at the neighborhood of $\\bar A = 1$. \\par These results show that gravitational waves discriminate very poorly the different models. Somehow, this is natural since the gravitational waves equation is sensitive essentially to the background behavior, in contrast with density perturbations which depends strongly on the kind of matter content. This is more accentuated in the comparison of different generalized Chaplygin gas models, that is, models with different $\\alpha$. In fact, the Chaplygin gas models are more sensitive to the parameter $\\bar A$ then $\\alpha$. \\par However, it must be stressed that even if the different models based on the generalized Chaplygin gas may be quite degenerate in what concerns the behaviour of gravitational waves, it is possible to distinguish them from the $\\Lambda$-CDM model by analyzing the typical amplitude of the spectra. But, a complete study of the value of this amplitude asks for a complementation of the analysis here done by introducing the primordial spectrum of gravitational waves, for example, that one coming from the primordial inflationary phase, which allows to fix preciselly the amplitude and the initial power spectrum \\cite{grisha}. \\par \\vspace{0.5cm} {\\bf Acknowledgements:} We thank CNPq (Brazil) for partial financial support."
    },
    "0404/astro-ph0404222_arXiv.txt": {
        "abstract": " ",
        "introduction": "The many-multiplet method, which played a central role in the search for time-variability of the fine-structure ``constant\" $\\alpha$ from the Keck/HIRES data of QSO absorption lines \\cite{mwf,mwf2}, has been applied recently to the VLT-UVES sample, with more restricted selection criteria \\cite{chand,srianand}.  The wavelengths of the lines have been determined by two different ways, using either laboratory wavelengths of Mg {\\footnotesize II,I} and Si {\\footnotesize II} lines with terrestrial isotopic abundances (case 1) or with those of the dominant isotopes (case 2) \\cite{chand}.  From case 1, which the authors claim to be more robust, they report no evidence of a changing $\\alpha$, in contrast to the previous Keck/HIRES result.  Despite the warning that more careful analysis on the likely systematic errors is necessary before final conclusion is reached \\cite{cowie}, it should be worth trying to apply some of the theoretical analyses focusing on different features of the observations accepted at face value for the time being, in order to provide hopefully a guide in entangling complications involved in the phenomenological analyses.    There has been a class of theoretical models in which a cosmological scalar field (dilaton or quintessence) shows itself through a time-dependent fine-structure constant \\cite{bj}--\\cite{motbar}.  In some of them the scalar field is expected even to be responsible for the cosmological acceleration \\cite{bj}--\\cite{bento}.   On the cosmological side, the scalar field beyond the linearization regime is supposed generally to be trapped to a potential superimposed on a smooth background, like an exponential potential, thus causing nearly the constant behavior of the sufficient amount of dark energy, the energy density of the scalar field acting as an effective cosmological ``constant.\" Due to this trapping process, the scalar field $\\sigma$ may show a damped-oscillation-like behavior, as a function of the cosmic time $t$, and hence of the fractional look-back time $s = 1-t/t_0$ with $t_0$ the present age of the universe.  We may also assume that the observed coupling strength is proportional to the scalar field, thus expecting $y(s)\\equiv \\Delta\\alpha/\\alpha\\times 10^5 =K \\Delta\\sigma(s)$, with $K$ a constant.   Another potentially serious issue comes from the Oklo phenomenon \\cite{oklo}--\\cite{lamoreaux}.  The detailed analysis of the remnants of the natural reactors supposed to have taken place about 2 billion years ago at Oklo, Gabon, West Africa resulted in either the upper bound $|\\Delta\\alpha/\\alpha|\\lsim 1.0\\times 10^{-8}$\\cite{fifonhom}, or the nonzero values $\\Delta\\alpha/\\alpha \\sim (0.9\\pm 0.1)\\times 10^{-7}$ or $\\sim (0.5\\pm 0.1)\\times 10^{-7}$ for different reasons in Refs. \\cite{fifonhom} and \\cite{lamoreaux}, respectively.  It has also been argued that it is unlikely that including the effects of the strong interaction enhances the bound more than an order of magnitude \\cite{yfptb}.  Comparing these estimates with the values at the level of $10^{-5}$ as expected currently from the QSO observations, we face a difference typically of two orders of magnitude or more.  An oscillatory behavior might provide a possible reconciliation between these two kinds of effects.   As for the suggested virialization process \\cite{motbar}, also suggested in this connection, more careful analyses seem yet to be applied in view of our unique features of the scalar field with respect to the structure of matter coupling and the partially finite-range force \\cite{cup,separt}.    In view of the uncertainties both on the theoretical and the observational sides, we follow a phenomenological approach to fit the observed $\\Delta\\alpha/\\alpha$ by the assumed simple damped oscillation for $y(s)$, as was attempted in Ref. \\cite{plb} applied to the 128 data points from the measurement on the Keck/HIRES QSO absorption lines \\cite{mwf}.  We assumed that the Oklo constraint corresponds approximately to a zero of the oscillation at $s_{\\rm oklo}= 0.142$, determining the other three parameters by best fitting the QSO data.  We then found that the 3-parameter fit is nearly as good as the weighted-mean-fit as far as the QSO result is concerned \\cite{plb}.    We now improve this fit by incorporating a natural zero at $s=0$, arising from the fact that $\\Delta\\alpha$ is the difference of $\\alpha$ from today's value and should vanish at $s=0$ by definition.  This condition is met by offsetting the damped oscillation, as implemented \\cite{yfptb} by \\beq y(s) = a\\left( e^{bs}\\cos \\left(v-v_1  \\right) -\\cos\\left( v_1 \\right)  \\right), \\label{vlt_1} \\eeq where $v/s=v_1/s_1= v_{\\rm oklo}/s_{\\rm oklo}=2\\pi T^{-1}$  with $v_1$ determined by \\beq v_1 = \\tan^{-1}\\left( \\frac{ e^{-b s_{\\rm oklo}} -\\cos (v_{\\rm oklo}) }{  \\sin (v_{\\rm oklo})}  \\right), \\label{vlt_2} \\eeq limiting ourselves at the moment to the behavior that allows many oscillations.   One easily verifies that $y(s)$ vanishes at $s=0$ and $s=s_{\\rm oklo}$.  The relaxation time $b^{-1}$ and the period $T$ are both measured in units of $t_0$.   We add that $b\\approx 2.5, T\\approx 0.22$ are expected from our typical cosmological solutions \\cite{cup,prd}, as shown in the beginning of Section 4 of Ref. \\cite{plb} and Table 2 of Ref. \\cite{yfptb}.  These values will be referred to as ``reference values\" in the following, though we do not exclude other solutions.    Equation \\reflef{vlt_1}) permits us to compute \\beq \\left( \\frac{\\dot{\\alpha}}{\\alpha} \\right)_{t_0} = -\\frac{1}{t_0}y'(0)\\times 10^{-5},\\label{vlt_3} \\eeq which may be different from the averaged rate of change; the weighted-mean divided by the average $\\bar{s}$, can be compared directly with the laboratory constraint. ",
        "conclusions": "Our analysis based on the damped-oscillation fit as suggested by a possible connection with the cosmological acceleration showed that the difference between the Keck/HIRES and VLT-UVES (case 1) results goes beyond a question merely of the presence or absence of the time-variability of $\\alpha$. It appears as if the oscillation of the scalar field shows quite different periods of time dependence.  Future efforts are expected to separate systematic errors focusing particularly upon a better determination of the period, thus testing the theoretical models favoring various oscillatory behaviors \\cite{bj}--\\cite{plb},\\cite{goldb}--\\cite{bento}.    In this respect we also add that there is an exceptional example of non-oscillating behavior in the simplified version of the two-scalar model \\cite{cup,prd} obtained by removing the second scalar field \\cite{dodel}, though we then lose the aesthetic advantage of the original model, in which no eternal inflation ensues.  The scalar field continues to fall overriding maxima of the sine-Gordon potential included in addition to the background exponential potential until it is finally trapped to one of the minima.  Unlike in the two-scalar model, a mini-inflation (persistenly as one of the repeated occurrences thus lessening the gravity of the coincidence problem) occurs only with a smooth change of the scalar field as illustrated in the lower panel of Fig. 6.  Note that $\\Omega_\\Lambda$ nevertheless passes through 0.7 as shown in the upper panel, enough to cause an acceleration of the universe.  We failed to detect any oscillation of $\\sigma$ generic to the two-scalar model.  This choice, still with the scalar field trapped temporarily, might be favored if future observations disprove oscillatory $\\Delta\\alpha/\\alpha$ and if consistency between QSO and Oklo is achieved by other means.  As shown in Fig. 7, however, the universe today might happen to be entering the final eternal inflation, thus featuring an oscillation, basically as in Refs. \\cite{goldb}--\\cite{bento}.  Note also that these two different behaviors are discussed in somewhat different context in Ref. \\cite{bento}.\\\\     \\begin{figure}[tbh] \\vspace*{-10em} \\hspace*{9em} \\epsfxsize=8cm \\epsffile{mzn_1_B1.eps} \\vspace*{-2pt} \\caption{A solution in the one-scalar model$^{26}$ with the potential $V(\\sigma)= e^{-4\\zeta\\sigma}[ \\Lambda +m^4( 1+\\cos(\\kappa\\sigma))  ]$, with $\\zeta=1.5, \\Lambda = 1.65 \\times 10^{-4}, m=0.42, \\kappa =0.351$ in reduced Planckian units, shown only around the present epoch, $t_0\\approx 13.8{\\rm Gy}$ or ${\\rm Log}\\: t_0 \\approx 60.2$.  The scalar field $\\sigma$ varies only smoothly (in the lower panel), whereas $\\Omega_\\Lambda $ passes through $\\sim 0.7$ (in the upper panel), hence causing acceleration of the universe.} \\end{figure}   \\begin{figure}[bth] \\vspace*{-10em} \\hspace*{9em} \\epsfxsize=8cm \\epsffile{mzn_3_B1.eps} \\vspace*{-2pt} \\caption{A solution for the final eternal inflation for $\\Lambda = 1.54\\times 10^{-4}$ in place of the one in Fig. 6 for repeated mini-inflations.  An oscillatory behavior of $\\sigma$ is shown in the lower panel with the vertical scale magnified by $\\sim 45$ times compared with Fig. 6, nearly the same size of the ``amplitude\" as in Fig. 5.10 of Ref. 7.  For ${\\rm Log}\\: t < 59.4$, the (invisible) curve of $\\sigma$ is almost identical with the corresponding portion in  Fig. 6.} \\end{figure}"
    },
    "0404/astro-ph0404279_arXiv.txt": {
        "abstract": "{In this talk, we take a new look at the theory of disk locking, which assumes that an accreting protostar rids itself of accreted angular momentum through a magnetic coupling with the accretion disk.  We consider that differential rotation between the star and disk twists the field lines.  For large enough twist, the magnetic field lines connecting the star and disk open and disconnect. This significantly reduces the spin-down torque on the star by the disk, and so we find that disk-locking theory predicts spin periods that are much too short to account for typical observed systems.}    \\addkeyword{accretion, accretion discs} \\addkeyword{MHD} \\addkeyword{stars: formation} \\addkeyword{stars: magnetic fields} \\addkeyword{stars: pre-main-sequence} \\addkeyword{stars: rotation}  \\begin{document} ",
        "introduction": "}    The collapse of a molecular cloud naturally leads to a phase consisting of a central protostar surrounded by a centrifugally supported accretion disk \\citep[for a review, see][]{bodenheimer95}. Disk winds, MRI turbulence, and/or viscous processes remove angular momentum from the disk and results in the accretion of material with high specific angular momentum onto the star.  For typical parameters for accreting protostars (CTTS's), the accretion alone will spin the star up to near breakup speed in less than $\\sim 10^5$ years, assuming the star hadn't already formed at near breakup speed.  Since the accretion lifetime is often greater than $10^6$ yr, the stars must rid themselves of this excess angular momentum.  Further, it has been generally accepted that accreting protostars spin more {\\it slowly} than their non-accreting counterparts \\citep[e.g.,][]{bouvierea93}. The general explanation is that the presence of an accretion disk somehow regulates the stellar spin, and then after the disk is dispersed, the star spins up as it contracts toward the main sequence.  \\citet{konigl91} applied the neutron star accretion model of \\citet{ghoshlamb79} to accreting protostars and showed that a ``disk-locking'' (DL) mechanism could explain the coincidence of accretion and slow rotation, in those systems.  According to DL theory, magnetic field lines connect the star to the disk (acting as ``lever arms'') and carry torques that oppose and balance the angular momentum deposited by accretion.  CTTS's are now known to posses kilogauss-strength fields \\citep[e.g.,][]{johnskrull3ea99}, and the general DL model has been invoked by many authors.  Recent observations of CTTS's in Orion by \\citet{stassunea99}, however, show no correlation between observed rotation periods and accretion diagnostics, calling the standard DL scenario into question. Furthermore, the magnetic fields of CTTS's, while strong, are disordered \\citep{safier98, johnskrullea99}, which reduces the effectiveness of magnetic torques required for DL.  These developments prompted us to revisit the general theory of DL.  In particular, the connectivity between the star and disk is an important issue.  Much recent work has shown that the magnetic connection between the star and disk is severed when the magnetic field is highly twisted.  Here, we show that the resulting spin-down torque is significantly reduced, and the DL model cannot account for accreting stars that spin slowly (e.g., $\\sim$10\\% of breakup speed). ",
        "conclusions": "The disk-locking scenario has recently been called into question by observations, as well as by theoretical considerations.  In particular, there are four, completely independent issues: \\begin{enumerate}  \\item{\\citet{stassunea99} found no correlation between accretion parameters and spin rates of TTS in Orion.}  \\item{CTTS's apparently do not have strong {\\it dipole} fields \\citep[e.g.,][]{safier98, johnskrullea99}.}   \\item{Stellar winds are expected to open field lines that would otherwise connect to the disk \\citep{safier98}.  A disk wind would have a similar effect.}  \\item{Finally, a large portion of the magnetic field connecting the star to the disk will open up, due to the differential rotation between the two \\citep[e.g.,][]{uzdensky3ea02}.  We have shown that the resulting spin-down torque on the star by the disk is less (by at least a factor of four and possibly by orders of magnitude) than calculated by previous authors.  The predicted equilibrium spin rate is therefore much faster.}  \\end{enumerate} \\adjustfinalcols So the DL scenario does not explain the angular momentum loss of the so called slow rotators---the group originally targeted by DL theory.  We conclude that, in order for accreting protostars to spin as slowly as 10\\% of breakup speed, there must be spin-down torques acting on the star other than those carried by magnetic field lines connecting the star to the disk.  The presence of open stellar field lines naturally leads to the possibility that excess angular momentum is carried by a stellar wind along those open lines.  We plan to investigate this possibility in the near future."
    },
    "0404/astro-ph0404565_arXiv.txt": {
        "abstract": "We present the first {\\it XMM-Newton} observations of the radio-quiet quasar \\mr\\ obtained in 2000 and 2002. The EPIC-pn spectra show a power-law continuum with a slope of $\\Gamma=1.6$ at high energies absorbed by at least two warm absorbers (WAs) intrinsic to the source. The underlying continuum in the earlier observation shows a ``soft excess'' at low X-ray energies which can be modeled as an additional power-law with $\\Gamma=2.9$. The spectra also show a weak narrow iron K$\\alpha$ emission line. The high-resolution grating spectrum obtained in 2002 shows emission lines from \\ion{N}{6}, \\ion{O}{7}, \\ion{O}{8}, \\ion{Ne}{9}, and \\ion{Ne}{10}, as well as absorption lines from the low-ionization ions of \\ion{O}{3}, \\ion{O}{4}, and \\ion{O}{5}, and other confirmed and suspected weaker absorption lines. The \\ion{O}{3}---\\ion{O}{5} lines are consistent with the properties of the emission line gas observed as extended optical [\\ion{O}{3}] emission in this source. The signal-to-noise of the 2000 grating data is too low to detect any lines. We suggest a model for the high-resolution spectrum which consist of two or three warm-absorber (WA) components.  The two-components model has a high-ionization WA with a column density of $10^{21.5}$--$10^{21.8}$ cm$^{-2}$ and a low-ionization absorber with a column density of $10^{20.3}$ cm$^{-2}$. In the three-components model we add a lower ionization component that produce the observed iron M-shell absorption lines. We investigate the spectral variations in \\mr\\ over a period of 8.5 years using data from {\\it ASCA}, {\\it BeppoSAX}, and {\\it XMM-Newton}. All X-ray observations can be fitted with the above two power laws and the two absorbers. The observed luminosity variations seems to correlate with variations in the soft X-ray continuum. The 8.5 year history of the source suggests a changing X-ray absorber due to material that enters and disappears from the line-of-sight on timescales of several months.  We also present, for the first time, the entire {\\it FUSE} spectrum of \\mr . We detect emission from \\ion{N}{3}, \\ion{C}{3}, and \\ion{O}{6} and at least 4 absorption systems in \\ion{C}{3}, \\ion{H}{1}, and \\ion{O}{6}, one at $-$580 \\kms\\ and at least 3 others which are blended together and form a wide trough covering the velocity range of 0 to $-$500 \\kms . The general characteristics of the UV and X-ray absorbers are consistent with an origin in the same gas. ",
        "introduction": "\\mr\\ was the first quasar detected by X-ray observations (by {\\it Ariel~V\\/} and {\\it SAS-3}; Cooke et al. 1978, Ricker et al. 1978), and also the first quasar where a warm absorber (WA) was suggested to explain the X-ray spectrum based on {\\it Einstein} observations (Halpern 1984).  The X-ray flux of the source is variable on timescales of $\\sim 10$ days, e.g., the {\\it EXOSAT} observations reported by Pan et al. (1990). These authors found the column density of the WA to vary and to correlate with the X-ray flux. Mineo \\& Stewart (1993) combined the earlier {\\it EXOSAT} observations with a {\\it GINGA} observation from 1989 and argued that the spectrum could be described by a power law with photon index $\\Gamma\\approx 1.7$ and a WA model with a column density of $10^{22.2}$\\,\\cmii . Using this model they found the ionization parameter to be strongly correlated with the source luminosity.  A deep {\\it ROSAT}/PSPC observation from 1993 was reported by Komossa (2001) who modeled the 0.1--2.4~keV spectrum using a WA with a column density of $10^{22.6}$\\,\\cmii .  \\mr\\ was observed with {\\it ASCA} ten times during 1993 and 1996.  These observations seem to be consistent with a WA with a column density of about $10^{21.3-21.7}$\\,\\cmii\\ (e.g., Reynolds 1997, Otani et al. 1998, Reeves \\& Turner 2000, Morales \\& Fabian 2002). \\mr\\ was also observed twice with {\\it BeppoSAX} during 1998. The two observations are separated by 5 months and show identical spectral shape and flux. The WA column was found to be $10^{21.9}$~\\cmii\\ and the difference from previous observations was attributed to the motion of the absorber across our line-of-sight (Orr et al. 2001). In summary, the observed 2--10 keV flux of \\mr\\ covers the range of $(1.7$--$5.1)\\times10^{-11}$ \\ergcms\\ which translates to a 2--10 keV luminosity of $(1.7$--$5.2)\\times10^{44}$~\\ergcms\\ (H$_{0}=70$ km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{M}=0.3$, $\\Omega_{\\Lambda}=0.7$ and assuming a $\\Gamma=1.6$ power law). The Galactic hydrogen column density towards \\mr\\ has been derived from 21~cm measurements to be $2.8\\times10^{20}$~\\cmii\\ (Lockman \\& Savage 1995).  \\begin{deluxetable*}{ccccc} \\tablecolumns{5} \\tablewidth{288pt} \\tablecaption{Observation Log for \\mr \\label{ascalog}} \\tablehead{ \\colhead{Mission} & \\colhead{Sequence Number} & \\colhead{Date} & \\colhead{Time (ks)\\tablenotemark{a}} & \\colhead{Rate (count\\,s$^{-1}$)\\tablenotemark{b}} } \\startdata {\\it ASCA } & 71035000 & 1993 Nov 11 & \\phn  3.6 & $1.69\\pm0.02$ \\\\ {\\it ASCA } & 71035010 & 1993 Nov 16 & \\phn  6.9 & $1.54\\pm0.02$ \\\\ {\\it ASCA } & 71035040 & 1993 Dec 12 & \\phn  9.9 & $1.74\\pm0.01$ \\\\ {\\it ASCA } & 71035060 & 1993 Dec 14 & \\phn  5.8 & $1.47\\pm0.02$ \\\\ {\\it ASCA } & 71035020 & 1993 Dec 19 & \\phn  6.7 & $1.28\\pm0.02$ \\\\ {\\it ASCA } & 71035050 & 1993 Dec 24 & \\phn  7.5 & $1.02\\pm0.01$ \\\\ {\\it ASCA } & 74028000 & 1996 May 26 &  17.8 & $0.65\\pm0.01$ \\\\ {\\it ASCA } & 74028010 & 1996 Jun 18 &  16.5 & $0.81\\pm0.01$ \\\\ {\\it ASCA } & 74028020 & 1996 Nov 27 &  15.9 & $0.57\\pm0.01$ \\\\ {\\it ASCA } & 74028030 & 1996 Dec 09 &  20.0 & $0.61\\pm0.01$ \\\\ {\\it BeppoSAX}  & 50556001 & 1998 Jun 14 & 47.5 &$0.419\\pm0.002$\\\\ {\\it BeppoSAX}  & 505560011& 1998 Nov 12 & 47.5 &$0.431\\pm0.002$\\\\ {\\it XMM-Newton}& 0112910301 & 2000 May 29 &\\phn 3.5&$17.91\\pm0.07$\\phn\\\\ {\\it XMM-Newton}& 0012940101 & 2002 May 18 & 44.7 &$\\phn 7.11\\pm0.01$ \\enddata \\tablenotetext{a}{Total effective exposure times determined from: SIS0 selected data for {\\it ASCA}, MECS data for {\\it BeppoSAX}, and EPIC-pn data for {\\it XMM-Newton}.} \\tablenotetext{b}{Count rates determined from: SIS0 over the 0.5--10 keV band for {\\it ASCA}, MECS over the 2--10 keV band for {\\it BeppoSAX}, and EPIC-pn over the 0.2--11 keV band for {\\it XMM-Newton}.} \\end{deluxetable*}  An Fe\\,K$\\alpha$ line was first suggested in a {\\it GINGA\\/} observation of \\mr, with an equivalent width (EW) of $125^{+100}_{-105}$ eV (Mineo \\& Stewart 1993).  This was later confirmed by the {\\it ASCA} and {\\it BeppoSAX} observations with EW of $190^{+140}_{-95}$ eV (Reynolds 1997; cf., Reeves \\& Turner (2000) found EW of $79\\pm52$ eV) and $62^{+12}_{-25}$ eV (Orr et al. 2001), respectively.  The redshift of \\mr\\ was determined by using nine optical narrow emission lines (Bergeron et al. 1983) to be $z=0.06398\\pm0.00006$ (we note that a few catalogs listed incorrect values, which has resulted in a variety of values quoted in the literature). The host galaxy has a gaseous component with temperature of $\\sim 3\\times10^4$\\,K (derived from the [\\ion{O}{3}] line ratio) and indications of low abundances of Ne, O, and N (Bergeron et al. 1983). The galaxy is surrounded by a giant \\ion{H}{2} envelope which is observed via [\\ion{O}{3}]\\,{$\\lambda$}5007 emission. \\mr\\ was observed by {\\it HST} at three epochs (Monier et al. 2001 and references therein). The spectrum shows clear Ly$\\alpha$ and \\ion{C}{4} absorption. Ganguly, Charlton, \\& Eracleous (2001) found the \\ion{C}{4} doublet absorption to vary with time, suggesting an intrinsic origin for this absorption. The quasar is radio quiet with a radio flux of 16.2$\\pm$0.7 mJy (NVSS catalog -- Condon et al. 1998).  This paper presents new {\\it XMM-Newton} and {\\it Far Ultraviolet Spectroscopic Explorer (FUSE)} observations of \\mr . We also carry out an in-depth analysis of the 10 available {\\it ASCA} observations and the two {\\it BeppoSAX} observations. We describe the data in \\S~2, perform an X-ray data temporal analysis in \\S~3, and discuss the X-ray spectral analysis in \\S~4 and the UV spectral analysis in \\S~5. In \\S~6 we elaborate on the implications of our results. ",
        "conclusions": "\\subsection{The Highly Ionized Gas in \\mr} \\label{historydiss}  The new {\\it XMM-Newton} data presented in this paper clearly show the presence of ionized gas seen in both emission and absorption in \\mr . The spectral analysis is severely limited by the poor S/N of the grating observations but several interesting results clearly emerge.  In the grating data we identified X-ray emission lines from \\ion{N}{6}, \\ion{O}{7}, \\ion{O}{8}, \\ion{Ne}{9}, and \\ion{Ne}{10}. We also identify, with high certainty, absorption lines from the low ionization ions of \\ion{O}{3}, \\ion{O}{4}, and \\ion{O}{5}, as well as the signature of absorption edges due to \\ion{O}{7} and \\ion{O}{8}. Many other absorption lines are probably detected (Table~\\ref{linetable}) but their reality and intensities are highly uncertain, because of the limited S/N.  For the 2002 RGS data we suggest one of two possible models. The first model consists of two absorbers: a highly ionized absorber with a column density of $10^{21.5-21.8}$ cm$^{-2}$  and log($U_{\\rm OX})=-1.68$, and a low ionization absorber with a column density of $10^{20.3}$ cm$^{-2}$ and log($U_{\\rm OX})=-4.00$. The second possibility is a three-component model where we split the highly ionized absorber from the above model into two components: one with log($U_{\\rm OX})=-1.4$ and the other with log($U_{\\rm OX})=-2.6$. The ionized (line-of-sight) absorber and the ionized emitter in both cases are consistent with being the same gas with a global covering factor of 0.4. The highly ionized absorption lines are probably narrower than 200 \\kms , a limit which is imposed by the equivalent width of the strongest predicted lines, given the column density, the ionization parameter, and the S/N of the observations.  More interesting conclusions are derived from analyzing the historical light curve of the source combining 8.5 years of observations by {\\it XMM-Newton}, {\\it ASCA} and {\\it BeppoSAX}. The main findings are: \\begin{enumerate} \\item All X-ray observations are consistent with a two-component continuum: a high energy power law of slope $\\Gamma = 1.6$ and a low energy soft excess component with $\\Gamma = 2.9$. Both components are absorbed by the WA and by the intrinsic neutral gas. \\item The WA observed during the 6 weeks of {\\it ASCA} observations in 1993 is consistent with being a single absorber with a column density of $10^{21.5}$ \\cmii . Less conclusive results are obtained for the short timescales behavior due to the poor S/N.  The data are consistent with a scenario in which the decrease in flux caused a corresponding decrease in the ionization parameter. Such a behavior has been suggested in the past for several other sources (e.g., MCG$-$6$-$30$-15$ -- the {\\it ASCA} observation of Otani et al. 1996; NGC\\,3516 -- the {\\it Chandra} observation of Netzer et al. 2002). This interpretation is not unique and the data also supports a more complex case where the source luminosity is not simply correlated with the ionization parameter (e.g., MCG$-$6$-$30$-15$ -- Orr et al. 1997; NGC\\,3783 -- Behar et al. 2003; Netzer et al. 2003) \\item On timescales of years, the WA properties are different and our model requires that the absorbing gas properties are changing in time. For example, the {\\it ASCA} 1996 observations clearly indicate a larger column density ($10^{21.8}$ \\cmii\\ vs. $10^{21.5}$ \\cmii ) and a smaller ionization parameter (log$(U_{\\rm OX}) \\sim -2.3$ vs. $\\sim -1.5$) absorber compared with the one observed in 1993. This could indicate new material entering our line-of-sight, between 1993 and 1996, adding to or replacing the earlier gas. The two groups of observations taken in 1996, that are separated by 5 months, are also not consistent with the notion of having the same WA. Comparing these two periods we find that the luminosity is about the same while the ionization parameter dropped by about a factor of 3. We suggest, again, a physical motion of the gas which resulted in a higher column density of material at about the same distance. Thus, a real change in the absorber properties can take place over time scales of only a few months.  The comparison of the 1993 {\\it ASCA} observations and the 2002 {\\it XMM-Newton} observation suggest a different change. In 2002, the AGN flux is smaller by a factor of 2 compared with 1993, yet the column density and the derived ionization parameter are about the same as in 1993. Similarly, a comparison between the two {\\it XMM-Newton} observations indicates that the source luminosity decreased significantly from 2000 to 2002 yet the derived WA properties remained about the same. This might mean that the absorbing material properties have changed between the two epochs (the SED in both is very similar but the luminosity decrease between 2000 and 2002 was not accompanied by a corresponding decrease in ionization parameter). An alternative explanation is that the absorbing material is very far from the central source and of low enough density such that it did not respond to the continuum luminosity variations. \\item The soft excess continuum luminosity is positively correlate with the hard continuum luminosity. \\end{enumerate}  The overall picture which emerges from this study is of a changing absorber made of material that enters and disappear from the line-of-sight on timescales of several months. On shorter timescales, of several weeks, the models are consistent with a picture in which the absorbing material responds instantly to the continuum luminosity variations. Due to the data quality and the model complexity (two power laws and several absorbers) we cannot unambiguously determine those properties.  \\subsection{The ``Neutral'' Absorber}  The 2002 {\\it XMM-Newton} observation show the presence of a low-ionization/neutral absorber intrinsic to \\mr. We derived the column density of this gas ($10^{20.3}$ \\cmii) from the shape of the soft X-ray continuum below 0.6 keV. The RGS data show evidence for \\ion{O}{3} and \\ion{O}{4} absorption lines which allow us to constrain its level of ionization (log$(U_{\\rm OX})=-4.0$).  Macchetto et al. (1990) find evidence for circumnuclear gas on distances between 3 and 6 kpc and of gaseous filaments farther out at distances of 30--50 kpc. Their lower limits on the [\\ion{O}{3}] density (e.g., their table 4) implies a column densities of $\\sim10^{20}$ \\cmii . This is entirely consistent with the properties of the low ionization absorber found in our analysis. Thus, it is possible that the same gas responsible for the [\\ion{O}{3}] emission is detected in absorption via \\ion{O}{3} and \\ion{O}{4} X-ray absorption lines. This means that parts of the emission line nebula observed in \\mr\\ lie in our line of sight to the central source.  \\subsection{Ultraviolet Emission and Absorption} \\label{hiresfuse}  This paper presents the first FUV spectrum of \\mr . We detect emission from \\ion{O}{6}, \\ion{N}{3}, and \\ion{C}{3}. We also detect at least 4 absorption systems in \\ion{O}{6}, \\ion{C}{3}, and \\ion{H}{1}, three of which are blended together. The three blended systems are best seen in \\ion{C}{3} (Figure~\\ref{velspec}b) and are definitely suggested in the other absorption lines (the main absorption system at $\\sim -300$\\,\\kms\\ looks like 3 systems blended together, see Figure~\\ref{velspec}a).  The \\ion{H}{1} Lyman absorption lines are seen all the way to the Lyman edge (Figure~\\ref{fusespec}). The lower series lines have similar EWs which suggests saturation. However, the lines are not completely black which means incomplete line of sight coverage.  Figure~\\ref{velspec}a demonstrate that the \\ion{O}{6} absorption system with the largest blueshift might not be saturated since the depth of the 1038\\,\\AA\\ line is about half the depth of 1032\\,\\AA , as expected from their oscillator strength ratios. On the other hand, the three blended absorption systems have the same depth in both lines of the \\ion{O}{6} doublet, indicating that these systems are saturated. The similarity of the \\ion{H}{1} and the \\ion{O}{6} absorption profiles suggests that all blended lines of the two ions are saturated and have the same covering fraction. Our spectral analysis indicates a covering factor between $\\sim$60\\% and $\\sim$90\\% for these systems.  The total depth in all the FUV and UV absorption lines is larger than the underlying continuum, thus the broad emission lines are absorbed by the UV absorber. This indicates that the UV absorber lies outside of the BLR. Using \\mr\\ continuum flux at 5100~\\AA\\ ($(37\\pm 3)\\times 10^{-16}$ \\ergcmsA ; Bergeron et al. 1983) and the relation between the BLR size and the object's luminosity (equation 6 of Kaspi et al. 2000) we find the BLR size to be $(1.2\\pm 0.2)\\times 10^{17}$ cm. We take this to be a lower limit on the distance of the UV absorber from the central black hole.  \\subsection{The Ultraviolet --- X-ray Connection}  Several studies suggested a link between the UV and X-ray absorber in AGNs (e.g., Mathur et al 1994; Mathur, Elvis, \\& Wilkes 1995; Shields \\& Hamann 1997; Crenshaw et al. 1999). The data presented in this paper enables us to study this connection in \\mr .  As explained, the resolution and S/N of the RGS spectrum does not allow the exact measurement of the X-ray absorption systems. We can only confirm that the velocity shifts and the FWHM of the UV absorption systems are consistent with the ones observed in the X-ray absorber. We found that the FUV absorption lines (in \\ion{H}{1}, \\ion{C}{3}, and \\ion{O}{6}) are blends of at least 4 absorption systems, all blueshifted with respect to the emission line by 0 to -600 \\kms . The {\\it HST} spectra (\\S~\\ref{hst}) show a Ly$\\alpha$ absorption line which is consistent with the {\\it FUSE} \\ion{H}{1} absorption lines. It also shows a \\ion{C}{4} absorption which has a width of 400 \\kms\\ and about the same blueshift range as seen in the FUV. The resolution of the RGS spectrum at 22 \\AA\\ is $\\sim$550 \\kms\\ hence any absorption lines which are similar in width to the UV absorption lines are predicted to be one pixel wide. In the X-ray spectrum we identify \\ion{O}{3}, \\ion{O}{4}, and \\ion{O}{5} which are consistent with the widths and blueshifts of the UV absorption lines, as well as hints for absorption from \\ion{O}{6}. We identify few absorption lines from highly ionized species with lines at wavelength shorter than 19~\\AA\\ (see Table~\\ref{linetable}). Given the poor S/N, this is consistent with the absorption seen in the UV spectra. The consistency between the UV and X-ray absorption suggests that they could arise in the same gas.  We have set a lower limit of $(1.2\\pm 0.2)\\times 10^{17}$ cm on the distance of the UV absorber from the source using the BLR distance. If the UV and X-ray absorptions are the same, this is also a lower limit on the distance of the X-ray absorber."
    },
    "0404/astro-ph0404423_arXiv.txt": {
        "abstract": "In view of the recent measurement of the scintillation velocity for PSR\\,J0737-3039, we examine the complete set of constraints imposed on the pulsar B natal kicks (magnitude and orientation) and predict the most favorable pulsar kick velocity and spin tilt for both isotropic and polar kicks. Our analysis takes into account both currently unknown parameters: the orientation of the orbital plane on the sky ($\\Omega$) and the radial component of the systemic velocity ($V_{r}$). Assuming that the system's peculiar velocity is entirely due to the second supernova explosion, we find that the system may have crossed the Galactic plane multiple times since the birth of the second neutron star and that the post-supernova peculiar velocity could have been as high as $\\simeq 1200$\\,km\\,s$^{-1}$. We also confirm the absolute lower and upper limits on the physical parameters derived in our earlier study. For specific combinations of the two unknown parameters $\\Omega$ and $V_r$, however, we find much tighter constraints on the pre-supernova binary configuration and natal kicks imparted to pulsar B, as well as on the age of system. Once $\\Omega$ is measured in the coming year, it will be straightforward to use the results presented here to further constrain the natal kicks and the spin-tilt predictions. We complete our comprehensive study and derive similar constraints and spin-tilt predictions for PSR\\,B1534+12, where the only free parameter is $V_{r}$. Lastly, for PSR\\,B1913+16, we update the progenitor and kick constraints using the measured pulsar spin tilt and allowing for Roche-lobe overflow from the progenitor of the pulsar companion. ",
        "introduction": "The recent discovery of a unique, strongly relativistic binary pulsar (Burgay et al. 2003) that is also the first eclipsing, double pulsar system found in our Galaxy (Lyne et al. 2004) promises to provide us with more impressive tests of general relativity than ever before (Kramer 2004), to greatly improve our understanding of magnetospheric and pulsar emission physics (Lyutikov 2004; McLaughlin et al. 2004; Demorest et al. 2004; Granot \\& M\\'esz\\'aros 2004; Arons et al. 2004), and possibly indicates a realistic chance for the detection of gravitational waves from double-neutron-star (DNS) inspiral with the first-generation of ground-based interferometers (Kalogera et al. 2004; Kim et al. 2004).  Apart from the binary pulsar and eclipsing behavior, PSR\\,0737-3039 sets the current records for DNS properties: for example, it has the shortest spin and orbital period, smallest eccentricity, lowest magnetic field, and fastest apsidal motion. The measured current orbital parameters and masses along with an estimate of the system's age have allowed the careful account of its recent evolutionary history, since the time just before the supernova (SN) that formed pulsar B, the current companion to the millisecond pulsar A (Willems \\& Kalogera 2004; Dewi \\& van den Heuvel 2004). In Willems \\& Kalogera (2004) (hereafter Paper I), we found that the pre-SN binary\\footnote{Throughout the paper, we use the terms pre- and post-SN to refer to the instants just before and right after the SN explosion forming the second NS.} most likely consisted of the first-formed NS and a Roche-lobe-filling helium star with mass between $2.1\\,M_\\odot$ and $4.7\\,M_\\odot$ (cf.\\ Dewi \\& van den Heuvel 2004). The pre-SN orbital separation can be constrained between $1.36\\, R_\\odot$ and $1.72\\, R_\\odot$, and the magnitude of the kick velocity between 60\\,km\\,s$^{-1}$ and 1560\\,km\\,s$^{-1}$.  Most recently, scintillation observations of pulsar A have allowed the measurement of the two systemic velocity components in the reference frame of the orbital plane projected on the sky (Ransom et al.\\ 2004). This measurement gives us the first glimpse of the kinematic evolutionary history (in addition to that related to binary evolution). Therefore we extend the analysis presented in Paper I and take a comprehensive approach to constraining the formation of PSR\\,J0737-3039. The concept is similar to that adopted by Wex, Kalogera, \\& Kramer (2000) with two main differences: (i) at present there is no constraint on pulsar A's spin tilt from timing or polarization observations; (ii) the current kinematic measurements involve two free parameters, the orientation of the orbital plane on the sky ($\\Omega$) and the radial component of the systemic velocity ($V_{r}$). \\begin{deluxetable*}{llccc} \\tabletypesize{\\scriptsize} \\tablewidth{0pt} \\tablecolumns{5} \\tablecaption{Parameters of PSR\\,J0737-3039, PSR\\,B1534+12, and PSR\\,B1913+16. \\label{param}} \\tablehead{ \\colhead{Parameter} & \\colhead{Notation} & \\colhead{PSR\\,J0737-3039\\tablenotemark{a}} & \\colhead{PSR\\,B1534+12\\tablenotemark{b}} & \\colhead{PSR\\,B1913+16\\tablenotemark{c}} } \\startdata Right Ascension (J2000)             & $\\alpha$     & $07^{\\rm h} 37^{\\rm m} 51.25^{\\rm s}$ & $15^{\\rm h} 37^{\\rm m} 09.96^{\\rm s}$ & $19^{\\rm h} 15^{\\rm m} 28.00^{\\rm s}$ \\\\ Declination (J2000)                 & $\\delta$     & $-30^\\circ 39^\\prime 40.74^{\\prime \\prime}$ & $11^\\circ 55^\\prime 55.55^{\\prime \\prime}$ & $16^\\circ 06^\\prime 27.40^{\\prime \\prime}$ \\\\ Proper motion in R.A.               & $\\mu_\\alpha$ (mas\\,yr$^{-1}$) & - &   $1.34 \\pm 0.01$ & $-3.27 \\pm 0.35$ \\\\ Proper motion in Decl.              & $\\mu_\\delta$ (mas\\,yr$^{-1}$) & - & $-25.05 \\pm 0.02$ & $-1.04 \\pm 0.42$ \\\\ Galactic longitude (J2000)          & $l$          & $245^\\circ\\!\\!\\!.\\,2$ & $19^\\circ\\!\\!\\!.\\,8$ & $50^\\circ\\!\\!\\!.\\,0$ \\\\ Galactic latitude (J2000)           & $b$          & $ -4^\\circ\\!\\!\\!.\\,5$ & $48^\\circ\\!\\!\\!.\\,3$ & $2^\\circ\\!\\!\\!.\\,1$ \\\\ Proper motion in Gal.\\ long.        & $\\mu_l$ (mas\\,yr$^{-1}$) & - & -21.57\\tablenotemark{d} & -0.60 \\\\ Proper motion in Gal.\\ lat.         & $\\mu_b$ (mas\\,yr$^{-1}$) & - & -12.80\\tablenotemark{d} & -3.38 \\\\ Distance                            & $d$ (kpc)    & $\\sim 0.6$ & $1.02 \\pm 0.05$ & $8.3 \\pm 1.4$ \\\\ Characteristic age of the ms pulsar & $\\tau_c$ (Myr) & 210 & 250 & 110 \\\\ Spin-down age of the ms pulsar      & $\\tau_b$ (Myr) & 100 & 210 & 80 \\\\ Mass of the ms pulsar               & $M_A$ ($M_\\odot$) & 1.34 & 1.33 & 1.44 \\\\ Mass of the companion               & $M_B$ ($M_\\odot$) & 1.25 & 1.35 & 1.39 \\\\ Current semi-major axis             & $A_{cur}$ ($R_\\odot$) & 1.26 & 3.28 & 2.80 \\\\ Current orbital eccentricity        & $e_{cur}$    & 0.0878 & 0.274 & 0.617 \\enddata \\tablenotetext{a}{Burgay et al. 2003; Lyne et al. 2004.} \\tablenotetext{b}{Wolszczan 1991; Stairs et al. 2002; Konacki, Wolszczan, \\& Stairs 2003; Arzoumanian, Cordes, \\& Wasserman 1999.} \\tablenotetext{c}{Hulse \\& Taylor 1975; Taylor et al. 1976; Taylor, Fowler, \\& McCulloch 1979; Taylor \\& Weisberg 1982, 1989; Damour \\& Taylor 1991; Arzoumanian et al. 1999, Wex et al. 2000.} \\tablenotetext{d}{Derived quantities.} \\end{deluxetable*} In this paper, we account for the complete set of equations describing the formation of PSR\\,J0737-3039 and derive constraints on the progenitor of pulsar B and on the kicks (both isotropic and polar) imparted to it at birth. We also derive predictions for the possible orientation of pulsar~A's spin axis with respect to the current orbital angular momentum axis. We carefully analyze and present the dependence of these results on $\\Omega$ and $V_{r}$ and on the system's age. The measurement of $\\Omega$ along with proper motion components is anticipated in the coming year. We show that such measurements can be used with the results presented here to derive tighter constraints.  We also compare our predicted spin tilts with those derived based on the geometric model of Jenet \\& Ransom (2004) for pulsar B's eclipses and find excellent agreement. We note that our predicted spin tilts can also be used along with polarization observations to anticipate the most likely timing of pulsar A's disappearance due to geodetic precession as a function of the radial velocity (which in the long term will be the only remaining unknown property) of the system.  In addition to PSR\\,J0737-3039 (Section~\\ref{sec0737}), we analyze our current understanding of the evolutionary history of the other two close currently known DNS systems in the Galaxy. For PSR\\,B1534+12, we use the masses, orbital characteristics, and measured proper motion to derive progenitor and kick constraints, as well as spin tilt predictions (Section~\\ref{sec1534}). For PSR\\,B1913+16, we extend the analysis of Wex et al.\\ (2000) to take into account the possibility that the helium star progenitor of the pulsar companion could have been filling its Roche lobe. We also derive the probability distributions of the kick magnitude imparted to the last-born NS (Section~\\ref{sec1913}). We end with a summary of our main conclusions, a discussion of the implications of our analysis for the relation of kicks to spin orientations, and a comparison of our results with other recent studies. ",
        "conclusions": "In this paper, we have extended the investigation by Willems \\& Kalogera (2004) and have derived a more elaborate and tighter set of constraints for the progenitor of PSR\\,J0737-3039 just before the SN explosion that forms the second NS as well as for the magnitude and the direction of the kick imparted to this NS at birth. The additional constraints stem from the kinematic history of the system which is derived by tracing the system's motion in the Galaxy backwards in time and by assuming that its peculiar velocity results entirely from the second SN explosion. In order to trace the system's motion backwards in time, we use the scintillation velocity measurements carried out by Ransom et al. (2004), which yield two velocity components in the plane perpendicular to the line-of-sight. Since the orientation of the velocity components is unknown, their orientation $\\Omega$ in the plane perpendicular to the line-of-sight introduces the first free parameter in the problem. In order to fully describe the system's space motion, the scintillation velocity components must furthermore be supplemented with the system's unknown radial velocity $V_r$, the second free parameter in the problem.  We find that, depending on the values of $\\Omega$ and $V_r$, the system may have crossed the Galactic plane up to two times. We identify these crossings as possible birth sites of PSR\\,J0737-3039 and use the times of the crossings as estimates for the system's age. It follows that there is a wide range of $\\Omega$- and $V_r$-values for which PSR\\,J0737-3039 may be remarkably young ($\\la 20$\\,Myr). If the system has crossed the Galactic plane twice, it is at least 20\\,Myr old regardless of the values of $\\Omega$ and $V_r$.  Our kinematical analysis furthermore shows that the post-SN peculiar velocity must have been at least on the order of $\\simeq 90$\\,km\\,s$^{-1}$ and may even have been as high as $\\simeq 1200$\\,km\\,s$^{-1}$ (see Fig.~\\ref{vcm1}). This is in contrast to the work done by Piran \\& Shaviv (2004) who used the present proximity of the system to the Galactic plane to derive a statistical 95\\% confidence upper limit of $\\simeq 150$\\,km\\,s$^{-1}$ on the post-SN peculiar velocity.  \\begin{deluxetable*}{lcccc} \\tabletypesize{\\scriptsize} \\tablewidth{0pt} \\tablecolumns{5} \\tablecaption{Summary of the constraints derived for PSR\\,J0737-3039, PSR\\,B1534+12, and PSR\\,B1913+16. For PSR\\,B1534+12, and PSR\\,B1913+16 numbers between parentheses correspond to the constraints associated with detached pre-SN progenitors. \\label{sum}} \\tablehead{ \\colhead{Parameter} & \\colhead{PSR\\,J0737-3039} & \\colhead{PSR\\,B1534+12} & \\multicolumn{2}{c}{PSR\\,B1913+16} \\\\ \\cline{4-5} \\\\ \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{$\\lambda$\\tablenotemark{f}$=18^\\circ \\pm 6^\\circ$} & \\colhead{$\\lambda$\\tablenotemark{f}$=162^\\circ \\pm 6^\\circ$} } \\startdata $M_0$ ($M_\\odot$)\\tablenotemark{a}    & 2.1--4.7 & 2.1--8.0 (4.7--8.0)  & 2.1--8.0 (4.5--8.0) & 2.1--8.0 (4.5--8.0) \\\\ $A_0$ ($R_\\odot$)\\tablenotemark{b}    & 1.2--1.7 & 2.4--4.3 (2.5--4.3)  & 1.1--5.3 (2.5--5.3) & 1.1--5.3 (2.5--5.3) \\\\ $V_k$ (km\\,s$^{-1}$)\\tablenotemark{c} & 60--1660 & 15--1350 (230--1350) & 190--600 (300--600) & 580--2000 (700--1250) \\\\ $\\theta$ (degrees)\\tablenotemark{d}   & 115--180 &  95-180  (140--180)  & 95--175  (145--175) & 160--175 (165--175) \\\\ $\\xi$ (degrees)\\tablenotemark{e}      &  25--155 &  5--175  (55--125)   & 5--85 or 95--175 (75--85 or 95--105) & 75--85 or 95--105 ($\\approx$85 or $\\approx$95) \\enddata \\tablenotetext{a}{Pre-SN mass of the last-born NS's helium star progenitor.} \\tablenotetext{b}{Pre-SN orbital separation.} \\tablenotetext{c}{Kick-velocity magnitude.} \\tablenotetext{d}{Angle between the kick direction and the pre-SN orbital velocity of the last-born NS's helium star progenitor.} \\tablenotetext{e}{Angle between the kick direction and the pre-SN orbital angular momentum axis.} \\tablenotetext{f}{Misalignment angle between PSR\\,B1913+16's spin axis and the post-SN orbital angular momentum axis.} \\end{deluxetable*}   Next, we used the age estimates and peculiar velocities associated with the Galactic plane crossings to impose additional constraints on the progenitor and formation of PSR\\,J0737-3039 besides those already discussed in Willems \\& Kalogera (2004). We find that despite the two degrees of freedom, the pre-SN and kick parameters as well as the age become much more constrained for many combinations of $\\Omega$ and $V_r$ (see Figs.~\\ref{pro1}--\\ref{pro3}). An overview of the constraints on the mass $M_0$ of pulsar~B's helium star progenitor, the pre-SN orbital separation $A_0$, and the magnitude $V_k$ of the kick velocity imparted to pulsar~B at birth as a function of $V_r$ for all possible values of $\\Omega$ is given in Fig.~\\ref{overview}. The absolute limits covering the whole range of admissible $\\Omega$- and $V_r$-values are summarized in Table~\\ref{sum}.  \\begin{figure*} \\begin{center} \\resizebox{5.4cm}{!}{\\includegraphics{f22acol.eps}} \\resizebox{5.4cm}{!}{\\includegraphics{f22bcol.eps}} \\resizebox{5.4cm}{!}{\\includegraphics{f22ccol.eps}} \\caption{Overview of the constraints on the mass $M_0$ of PSR\\,J0737-3039B's helium star progenitor, the pre-SN orbital separation $A_0$, and the magnitude $V_k$ of the kick velocity imparted to pulsar~B at birth, as a function of $V_r$ for all possible values of $\\Omega$.} \\label{overview} \\end{center} \\end{figure*}  We also examined constraints on the kick direction in terms of the angle $\\theta$ between the kick direction and the pre-SN orbital velocity of pulsar~B's helium star progenitor, as well as the angle $\\xi$ between the kick direction and the pre-SN orbital angular momentum vector (which could be associated with pulsar B's spin axis). We find that $\\theta$ is always larger than $\\simeq 115^\\circ$, while $\\xi$ ranges from $\\simeq 25^\\circ$ to $\\simeq 155^\\circ$. The kick imparted to pulsar~B at birth must therefore have been directed opposite to the orbital motion and cannot have been too closely aligned with the pre-SN orbital angular momentum axis.  For each combination of $\\Omega$ and $V_r$ yielding viable progenitors for PSR\\,J0737-3039, we derived probability distributions for the magnitude of the kick velocity under the assumption that all kick directions are equally probable (Fig.~\\ref{vk1}). The distribution functions exhibit a single peak at kick velocities which increase from $\\simeq 100$\\,km\\,s$^{-1}$ to $\\simeq 1600$--1700\\,km\\,s$^{-1}$ with increasing absolute values of the radial velocity. This is in contrast to Willems \\& Kalogera (2004) who found a single most probable kick velocity of $\\simeq 150$\\,km\\,s$^{-1}$. The dependence of the position of the peak on the radial velocity found here stems from the relation between the kick magnitude and the post-SN peculiar velocity which tends to select those kick velocities which are able to explain the system's motion in the Galaxy (high space velocities cannot be explained by low kick velocities and vice versa). However, for a given radial velocity, the range of kick velocities with non-vanishingly small probabilities found here is always much more constrained than in Willems \\& Kalogera (2004).  Similarly, we derived distribution functions for the misalignment angle $\\lambda$ between pulsar~A's spin axis and the post-SN orbital angular momentum axis (Fig.~\\ref{tilt1}). Tilt angles lower than $30^\\circ$--$50^\\circ$ are clearly favored in the cases where (i) the system has crossed the Galactic plane twice in the past and (ii) the system crossed the Galactic plane once and has a current radial velocity of less than $\\simeq 500$\\,km\\,s$^{-1}$ in absolute value. Higher tilt angles must be associated with one disk crossing and much higher present-day radial velocities of $\\simeq 500$--1200\\,km\\,s$^{-1}$ in absolute value. Tilt angles close to $90^\\circ$ are furthermore strongly disfavored for any radial velocity $V_r$.  Our probability distributions are thus compatible with the tilt-angle predictions of $16^\\circ \\pm 10^\\circ$ and $164^\\circ \\pm 10^\\circ$ made by Jenet \\& Ransom (2004) and with the spin- magnetic dipole misalignment angle derived by Demorest et al. (2004) which rules out Jenet \\& Ransom's alternative solutions of $82^\\circ \\pm 16^\\circ$ or $98^\\circ \\pm 16^\\circ$.  In view of recent claims of alignment of NS kicks with NS spin axes (e.g. Romani 2004), we also considered kick-velocity and tilt-angle distributions for polar kicks where the kick direction is confined within two oppositely directed cones with an opening angle of $30^\\circ$ and with axes parallel to the pre-SN orbital angular momentum axis (i.e. $\\xi \\le 30^\\circ$).  The confinement of the kick directions strongly reduces the range of radial velocities for which viable progenitors for PSR\\,J0737-3039 may be found. Correspondingly, the range of most probable kick velocities shrinks to $\\simeq 200$--550\\,km\\,s$^{-1}$, and the range of most probable tilt angles to $15^\\circ$--$45^\\circ$ (see Figs.~\\ref{vk1b} and~\\ref{tilt1b}).  Progenitor constraints and isotropic kick-velocity and spin-tilt distributions were also derived for PSR\\,B1534+12 and PSR\\,B1913+16. Both systems have a measured proper motion with a known direction in the plane of the sky, so that the derived constraints and probability distributions only depend on the unknown radial velocity $V_r$. For PSR\\,B1913+16, the knowledge of the tilt angle $\\lambda$ (Kramer 1998; Weisberg \\& Taylor 2002) furthermore puts additional constraints on the possible progenitor systems, which are not yet available for PSR\\,J0737-3039 and PSR\\,B1534+12.  The absolute limits on the pre-SN binary and kick parameters of the two systems are summarized in Table~\\ref{sum} (the dependency of the constraints on $V_r$ is shown Figs.~\\ref{pro10} and~\\ref{pro9}). Contrary to PSR\\,J0737-3039, the progenitors of these two systems were not necessarily undergoing a mass-transfer phase at the time the second NS was born. The helium star masses $M_0$, pre-SN orbital separations $A_0$, and kick velocities $V_k$ associated with detached progenitors are, however, always much more constrained than those associated with the full set of detached and semi-detached progenitors. In particular, the possibility of pre-SN Roche-lobe overflow considerably relaxes the constraints on the progenitor of PSR\\,B1913+16 derived by Wex et al. (2000).  The constraints on the direction of the kick imparted to the second-born NS in PSR\\,B1534+12 and PSR\\,B1913+16 yield very similar conclusions as for PSR\\,J0737-3039: the kicks are generally directed opposite to the orbital motion and cannot have been too closely aligned with the pre-SN orbital angular momentum axis.  The kick-velocity and spin-tilt distributions for PSR\\,B1534+12 and PSR\\,B1913+16 (Figs.~\\ref{vk4}, \\ref{tilt3} and~\\ref{vk2}) also show qualitatively the same behavior as for PSR\\,J0737-3039. The most probable kick velocity depends on the unknown radial velocity and, depending on the value of $V_r$, can be anywhere between the derived lower and upper limits for the kick imparted to the second-born NS. Tilt angles smaller than $30^\\circ$--$50^\\circ$ are again favored when (i) the system crossed the Galactic plane more than once and (ii) the system crossed the Galactic plane once and has a current radial velocity smaller than $\\simeq 250$\\,km\\,s$^{-1}$ in absolute value for PSR\\,B1534+12 and smaller than $\\simeq 500$\\,km\\,s$^{-1}$ in absolute value for PSR\\,B1913+16.  For conclusion, we point out that the set of observational measurements and constraints available for double neutron stars systems are sufficient to greatly limit the properties of the progenitors. As a follow-up investigation it will be therefore interesting to compare these strongly constrained progenitor properties to synthesis studies of binary populations and double neutron star formation."
    },
    "0404/astro-ph0404109_arXiv.txt": {
        "abstract": "The cosmic infrared background (IRB) at wavelengths between 1 $\\mu$m and 3 $\\mu$m provides a useful probe of early star-formation prior to and during reionization. To explain the high optical depth to electron scattering, as measured by the Wilkinson Microwave Anisotropy Probe (WMAP), one requires significant star-formation activity at redshifts 10 and higher. In addition to massive stars, the IRB flux may be contributed by a population of early miniquasars. We study the relative contributions from first stars, supernovae and quasars to the IRB for reasonable star formation rates at high redshift. If miniquasars radiate efficiently at the Eddington-limit, current background measurements limit the fraction of mass in first stars that is converted to seed black holes to be roughly less than 10\\%. In the case of supernovae, though an individual supernova is much brighter than the progenitor star, due to the shorter lifetime of order few months, the fractional contribution to the IRB remains at a level of 10\\% and below when compared to the same contribution from stars. The bright supernovae may, however, be directly detectable by future large ground-based and space telescopes. ",
        "introduction": "The Wilkinson Microwave Anisotropy Probe (WMAP) has provided strong evidence for an optical depth for electron scattering of 0.17 $\\pm$ 0.04 based on the large scale polarization pattern related to rescattering of Cosmic Microwave Background (CMB) photons (Kogut et al. 2003). If the reionization process is described as instantaneous and homogeneous, the measured optical depth implies a reionization redshift of $\\sim 17 \\pm 5$ in a spatially flat universe. Interestingly, the derived redshift for reionization is at the high end of expectations in models in which stellar populations are dominant reionization sources (e.g., Cen 2003; Fukugita \\& Kawasaki 2003; Venkatesan, Tumlinson \\& Shull 2003; Wyithe \\& Loeb 2003; Sokasian et al. 2004). Thus, in some of these models, primordial stars, the so-called Population III stars, are assumed to be rather massive and hence have a large UV photon emission rate, as suggested by recent theoretical studies (e.g. Abel, Bryan \\& Norman 2002; Bromm, Coppi \\& Larson 2002; Bromm, Kudritzki \\& Loeb 2001; Schaerer 2002, 2003).  The formation of such massive stars in the early universe has many important cosmological implications (e.g., Carr, Bond, \\& Arnett 1984; Oh, Cooray \\ Kamionkowski 2003; Yoshida, Bromm \\& Hernquist 2004). Because radiation below the Lyman-limit is absorbed by neutral hydrogen, luminous sources at redshifts 10 and higher are generally seen only in the near-IR band. In particular, if Pop~III stars are formed primarily at redshifts between 10 and 30, they are expected to contribute to the infrared background (IRB) light at wavelengths between 1 and 5 $\\mu$m at present epoch (Bond, Carr \\& Hogan 1986). Recent estimates based on theoretical models suggest that a large fraction of the IRB total intensity may indeed be due to these stars (Santos et al. 2002; Salvaterra \\& Ferrara 2003; Cooray et al. 2004). A substantial IRB could arise from the Pop~III stars not only due to the direct emission associated with these stars, but also due to indirect processes that lead to free-free and Lyman-$\\alpha$ emission from the ionized nebulae, or H{\\sc ii} regions, surrounding these stars.  Observationally, about 20-40\\% of the near-IR flux (between $\\sim$ 1 $\\mu$m and 3 $\\mu$m) has been resolved by point sources (Cambr\\'esy et al. 2001; Totani et al. 2001); while J- and K-band source counts go down to AB-magnitudes of $\\sim$ 28 and 25, respectively, the cumulative surface brightness converges by a magnitude around 23, say, in the K-band (see for a review of existing data by Pozzetti \\& Madau 2000). The missing IR flux could be either due to a high-z population of galaxies, such as the proto-galaxy population related to first stars (Santos et al. 2002; Salvaterra \\& Ferrara 2003), or a population of low surface brightness galaxies at nearby redshifts. While there still remains uncertainties, one can potentially understand the presence of first luminous sources and some details related to their population, such as the number density, by characterizing the background light from optical to infrared at wavelengths of a few micron. In this wavelength range, final products of these massive stars are also expected to contribute to IRB, in addition to the direct emission from stars; Radiation from supernovae and primordial black holes, in the form of miniquasars, may be significant sources of emission. The role of final products in contributing to the IRB, relative to the progenitor stars, depends on the initial mass function of the first stars.  As the end state of a massive star above $\\sim 260 \\msun$ is a complete collapse to a black hole (Heger \\& Woosley 2002), if the first stellar population is extremely top-heavy with all mass above this limit, then one expects roughly a similar number density of black holes as that of stars. These black holes grow via accretion and/or mergers and may partly be responsible for early reionization if they radiate as miniquasars (Haiman \\& Loeb 1998; Madau et al. 2004; Ricotti \\& Ostriker 2003). The UV output from such quasars, if not absorbed heavily by the surrounding torii, can contribute both to reionization and to the IR flux as viewed today. The unresolved soft X-ray background at energies between 0.5 and 2 keV may rule out a large density of miniquasars at high redshifts as required to reionize the universe from the quasar emission alone (Dijkstra, Haiman \\& Loeb 2004) unless reionization by miniquasars, within uncertainties of models used in Dijkstra et al., is completed by a redshift of 20. Here, we suggest that if all massive stars are eventually converted to black holes and accrete efficiently at the Eddington-rate (as would be the case if miniquasars are responsible for eventual reionization of the Universe), then the IR background is overproduced. A more realistic situation could be that the mass function is broad, and only a small fraction of the total initial stellar mass is converted to black holes; current understanding of the IRB may limit the mass fraction of stars above 260 M$_{\\sun}$ to a level less than 10\\%. This constraint can be relaxed if miniquasars radiate less efficiently below the Eddington-limit; such a scenario may be viable if stellar sources dominate the UV photon production and are, thus, responsible for reionization, though some fraction of stellar mass is converted to miniquasars.   The initial mass function of the first stars could potentially be dominated by stars with masses in the range between 140 and 260 $\\msun$. In this case, the final product will be a complete disruption of the star via the pair-instability process (Heger \\& Woosley 2002). In the case of such massive stars, in general, one expects a typical supernovae to be much brighter than the star, suggesting that a modest contribution to the IRB could come from supernovae, when compared to the fractional contribution from stellar emission alone. However, though supernovae are in fact brighter than an average star, the duration over which most optical flux is emitted is relatively smaller. We show that this results in a small but non-negligible contribution to the IRB from supernovae associated with massive Pop~III stars. While the background contribution is smaller, these supernovae are likely to be detectable directly in deep IR images above 1 $\\mu$m down to AB-magnitude limits at the level of 26 and fainter.  In the next section, we discuss the contribution to IRB from the first generation of stars, supernovae, and miniquasars under the assumption that the early stellar population is dominated by very massive stars. In the case of supernovae, we will suggest that, while the relative contribution to the IRB may be lower, individual detections may be possible with deep-IR imaging. We also suggest that the IR background itself may be used to constrain the presence of early mini-quasars. Throughout the paper, we make use of a $\\Lambda$CDM cosmological model consistent with current data (e.g., Tegmark et al. 2004) with $\\Omega_m=1-\\Omega_\\Lambda=0.3$, $h=0.7$, $\\Omega_b=0.04$, the spectral index of the primordial power $n=1$, and a normalization for the matter power spectrum with $\\sigma_8=0.9$.  \\begin{figure} \\centerline{\\psfig{file=Fig1a.eps,width=3in,angle=-90}} \\centerline{\\psfig{file=Fig1b.eps,width=3in,angle=-90}} \\caption{{\\it Top panel:} The global star formation rate density at redshifts greater than 10 calculated based on the Press-Schechter (PS) mass function with a minimum virial temperature of 400 K (top line) and 10$^4$ K (bottom line). {\\it Bottom panel:} The baryon fraction in black holes as a function of redshift, under the assumption that the star-formation history related to the seed population follows the two curves shown in the top panel. We have assumed a 5\\% of the mass in stars are converted to seed black holes. The three curves in each of the star-formation models follow a model description related to the duty cycle of accretion growth following Ricotti \\& Ostriker (2003) with parameters shown on the figure (see Section~2.1 for a discussion).} \\label{fig:sfr} \\end{figure}   \\begin{figure} \\centerline{\\psfig{file=Fig2.eps,width=3.2in,angle=-90}} \\caption{The flux spectrum, $\\nu F_\\nu$, of a 100 $M_{\\odot}$ star at a redshift of 15 as observed today as a function of the observed wavelength. In addition to the stellar spectrum (solid line), we also show the nebular Lyman-$\\alpha$ emission (long-dashed line) and free-free (dashed line) of the ionized H{\\sc ii} region surrounding the star, when the UV end of the stellar-spectrum (thin solid line) is used to ionized the surrounding medium. The spectrum follows from model calculations by Santos et al. (2002). We assume all ionizing photons are absorbed by the nebula. For comparison, we also show the spectrum of a miniquasar of the same mass ($M_{\\rm BH,seed}=100M_{\\odot}$) at a redshift of 15, as viewed today (solid line labeled 'miniquasar'). The spectrum is computed using the template derived in Sazonov et al. (2004). The bump around $\\sim$ 100$\\mu$m shows a possible contribution from thermal reradiation of UV photons by dust. While the overall flux from a miniquasar is similar to that of a star, we argue that the contribution to IRB from miniquasars could dominate stars because of their larger number density; quasars live longer than massive stars.} \\label{fig:star} \\end{figure} ",
        "conclusions": "The cosmic infrared background (IRB) at wavelengths between 1 $\\mu$m and 3 $\\mu$m provides a useful probe of the global star-formation rate prior to and during the reionization. We have studied the relative contributions from first stars, supernovae and early miniquasars to the infrared background (IRB). In addition to massive stars, the IRB flux may be dominated by a population of miniquasars, especially if the first generation of stars are very massive and the end product of the stars lead to a substantial population of black holes that will effectively radiate as miniquasars by matter accretion.  We use the Press-Schechter formalism that describes dark halo formation, and combine it with a simple star-formation model to calculate the cosmic star-formation history, the supernova rate, and the blackhole formation rate at high redshift. We follow previous calculations of the source spectrum for stars, while for Population~III supernovae, we use a simple description of the peak emission and its duration. We describe the flux spectrum of miniquasars using an average spectrum as observed for luminous quasars at redshifts less than 5 based on observed data and models by Sazonov et al. (2004). We fix the freedom related to the overall normalization of this spectrum by assuming that the miniquasars radiate efficiently at the maximum allowed by the Eddington-limit. Under this assumption, we find that a mass fraction of seed black holes, relative to Pop~III stars, at the level of $\\sim 10\\%$ and below is consistent with observations. This upper limit on the seed mass fraction can naturally be extended to a higher value if miniquasars radiate below the Eddington-limit, so the limit we derive should be considered as on the lower side rather than a strict upper limit. Future spectroscopic background measurements exploiting rockets will improve the observational uncertainties and hence the constraints on the high-redshift miniquasar population.  While the integrated light from Population III supernovae to the IRB is subdominant, one can potentially consider the possibility of directly detecting individual supernovae in high resolution deep IR imaging data. Given the star formation rate density and the expected flux, we estimate the surface density of supernovae to be order few tens per sqr. degree over a year, if the star-formation rate is close to 0.5 M$_{\\odot}$ yr$^{-1}$ Mpc$^{-3}$ at $z\\sim 10$, and assuming that all stars are massive ($ > 100$  M$_{\\odot}$). For the high redshift supernovae, the typical AB magnitudes are at the level of 26 in 1.5 $\\mu$m (Heger et al. 2002). These are within reach of deep IR imaging with existing large telescopes and from space and, in the long term, clearly within reach of missions such as the James Webb Space Telescope (JWST), which can detect point sources down to a magnitude limit of 31 around the same wavelengths in a 10$^4$ sec exposure in its 16 arcmin$^2$ field of view \\footnote{http://www.ngst.nasa.gov for sensitivities and other details}.  \\vspace{0.5cm}  {\\it Acknowledgments:} This work is supported by the Sherman Fairchild foundation and DOE DE-FG 03-92-ER40701 (AC). NY thanks support from Japan Society of Promotion of Science Special Research Fellowship (02674). We thank an anonymous referee for a careful reading of the manuscript and for detailed comments that improved the discussion in this paper."
    },
    "0404/physics0404008_arXiv.txt": {
        "abstract": "We use the relativistic Hartree-Fock method, many-body perturbation theory and configuration-interaction method to calculate the dependence of atomic transition frequencies  on the fine structure constant $\\alpha=e^2/\\hbar c$. The results of these calculations will be used in the search for variation of the fine structure constant in quasar absorption spectra. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404345_arXiv.txt": {
        "abstract": "X-ray spectroscopy of cooling-flow clusters reveal an unexpected deficit of soft X-ray emission lines from gas at $\\sim 1/3$ of the ambient plasma temperature, across a wide range of X-ray luminosities and virial temperatures. We propose excess ionization from either a population of suprathermal electrons or photoionization by X-ray continuum emission from hot gas or a central AGN as a means of decoupling the thermal state of the gas from its emission line spectrum. The former effect is thought to operate in the solar corona. Because they generally become important at some fixed fraction of the cluster gas temperature, such mechanisms could in principle explain both the universality and temperature dependence of the emission line suppression, properties which none of the present-day models based on gas heating can explain. Ultimately these models cannot explain the observations; however, they have attractive and robust features which could be useful in elucidating a final solution to the soft X-ray deficit. ",
        "introduction": "Cooling flow clusters have long been a mystery. While the X-ray luminosity is consistent with gas cooling rates of up to $100-1000 {\\rm M_{\\odot} \\, yr^{-1}}$, whether such a cooling flow is indeed really operating and the fate of the cooled gas has been a matter of long-standing controversy (see \\scite{Fabian1994} for a review). X-ray spectroscopy acquired by the Reflection Grating Spectrometer on XMM-Newton has added yet another puzzle to the picture: there is a strong deficit of emission lines associated with the low temperature gas which should be present in a cooling flow. In particular, the observed differential emission measure from clusters is consistent with \\cite{petersonetal2002}: \\begin{equation} \\frac{d EM}{dT}= \\frac{5}{2}\\frac{\\dot{M}}{\\mu m_{p}} \\frac{k_{B}}{\\Lambda(T)} (\\alpha +1) \\left( \\frac{T}{T_{o}} \\right)^{\\alpha} \\label{eqn:emission_measure} \\end{equation} where $\\dot{M}$ is the inferred mass dropout rate, $T_{o}$ is the background temperature, and $\\alpha \\approx 1-2$. This conflicts with the standard isobaric cooling model, which predicts $\\alpha=0$, independent of the details of gas clumping or geometry (\\scite{Fabian1994} and references therein). The observed emission measure of low temperature gas is therefore lower than expected. In particular, while the spectra exhibits strong emission from the ambient temperature $T_{o}$ down to about $T_{o}/2$--which suggests that cooling is indeed taking place--there is signifiantly less emission from lower temperature gas than is expected. What is particularly striking is that this seems to be a universal phenomenon, across a wide range of cluster temperatures and mass deposition rates: emission from gas at $\\sim T_{o}/3$ is always strongly suppressed. This hints at some common underlying physical mechanism which scales with the plasma temperature.  A whole host of explanations have been put forth (for reviews, see \\scite{fabianetal2001}, \\scite{kaastraetal} and references therein). Many center on heating the gas, to prevent it from cooling below some minimum temperature (e.g. through AGN outflows, electron thermal conduction, mergers, or magnetic reconnection). This suffers from two generic problems: firstly, it is notoriously difficult to heat gas in a stable manner; for instance, conduction tends to either make the gas isothermal (contrary to observations), or allow cooling to continue unimpeded, unless the conduction suppression factor is excessively fine-tuned \\cite{bregmandavid,meiskin88} (though see \\scite{kim_narayan} for a countervailing claim). Secondly, such solutions are much more effective at altering the normalization of the emission measure (by reducing $\\dot{M}$), rather than its shape. Indeed, for consistency with observations, heating must either completely suppress local thermal stability or heat the gas all the way back up to the ambient temperature. If heating delays but does not halt cooling at some lower temperature, this will increase the emission measure of low temperature gas over predictions of the standard isobaric cooling model, rather than reduce it. Another possibility is to cool gas rapidly by some process which does not involve X-ray emission (e.g. turbulent mixing, dust cooling), though such proposals to date suffer from similar difficulties in that it is difficult to naturally identify a characteristic temperature scale at which such effects become important (as is required by observations). Other proposals (more similar in spirit to the route taken in this paper) seek to suppress the emission lines by resonant scattering or absorption \\cite{petersonetal2001}, or by an inhomogeneous metallicity distribution which significantly alters the predicted emission line spectrum \\cite{morrisfabian}. However, none have proven wholly successful, and the problem remains unsolved.  In this paper we pursue another avenue: breaking the expected correspondence between the thermal state of cooling gas and its emission spectrum. If the ionization stage of an atom does not correspond to its value in collisional ionization equilibrium, then the emission line spectrum ceases to be a good thermometer of the plasma temperature. From equation (\\ref{eqn:emission_measure}), the luminosity in an emission line at frequency $\\nu$ from an ionization stage $i$ is given in the standard isobaric cooling flow model (where $\\alpha=0$) by: \\begin{equation} L_{\\nu}= \\frac{5}{2}\\frac{k_{B}}{\\mu m_{H}}{\\dot{M}} \\int_{T_{min}}^{T_{max}} \\frac{\\epsilon_{\\nu}(T,Z)}{\\Lambda(T,Z)}{dT}. \\label{eqn:line_luminosity} \\end{equation} Note the inverse weighting by the cooling function: emission at lower temperatures is given progressively less weight as the gas cools rapidly through that phase. Since $dL_{\\nu} \\propto \\left[ x_{i}(T)/\\Lambda(T) \\right] dT$ (where $x_{i}$ is the ionization fraction of stage $i$), if we can delay the appearance of an ionization stage $i$ until low temperatures when cooling is rapid, the emission measure of its associated lines will be severely reduced. This could happen, for instance, if the gas falls out of collisional equilibrium; however, for conditions in cooling flow clusters $t_{\\rm rec} \\ll t_{\\rm cool}$, so this is unlikely. On the other hand, the gas can also be overionized if there is some additional source of ionization. We consider two sources: collisional ionization from a non-Maxwellian tail of suprathermal electrons, and photo-ionization by free-free emission from hot intracluster gas. Both of these effects are associated with the large reservoir of hot gas in the cluster, and the associated ionization scales naturally with the temperature of the ambient gas. Thus, the onset of excess ionization takes place at some fraction of the cluster virial temperature. This would naturally explain the universality of the low temperature cut-off and the reason why it always appears at some fixed fraction ($\\sim 1/3$) of the temperature of the hot gas, features which none of the presently popular models can explain. For instance, it is difficult to understand why heating mechanisms like AGN heating should target low temperature gas at the expense of high temperature gas, particularly when cool gas occupies a very small fraction of the cluster by volume.  Despite the above promising characteristics, we find that excess ionization fails to solve the soft X-ray cooling flow problem, for relatively general and model-independent reasons. In clusters, suprathermal electrons are much more efficient at heating the gas via Coulomb collisions rather than ionizing the metals (unlike in the solar corona, where their ionization effects have been invoked to explain anomolous ionization patterns, e.g. \\scite{scudder94}). Photoionization can certainly produce excess ionization at low temperatures with little associated photo-heating, since the the ionizing photons only target the metals, which are a rare tracer species. However, the same excess ionization reduces the efficiency of gas cooling at low temperatures which ultimately {\\it increases} the emission measure of low ionization stage emission lines. This arises because the missing emission lines are produced at temperatures where line cooling, rather than bremstrahhlung, is the dominant cooling mechanism; one therefore cannot alter the ionization properties of the gas without altering its cooling properties. Despite the failure of these two attempts, the generality and temperature scaling of such alterations to standard isobaric cooling emission spectrum models is very attractive, and we feel, under-explored. It is possible that some variant of these ideas may ultimately prove successful. ",
        "conclusions": "In this paper we have examined two means of decoupling the ionization stage of metal ions from the plasma temperature: excess ionization from a non-Maxwellian tail of suprathermal electrons, and photo-ionization either from free-free emission by hot cluster gas or from internal AGN or X-ray binaries. These are very attractive because they should be universal and also become important at some fixed fraction of the ambient plasma temperature, observed properties which none of the present-day models based on gas heating can explain. Neither of these are successful: suprathermal electrons are much more efficient at heating the gas than ionizing the metals, and photo-ionizing radiation decreases gas cooling at low temperatures, ultimately {\\it increasing} the emission measure of emission lines associated with low temperature gas. Nonetheless, they have attractive features which may prove useful in elucidating a final solution. In a companion paper, we shall examine the feasibility of solutions involving heating. In particular, we shall examine the global stability of cooling flow clusters, addressing the classic cooling flow problem of mass dropout."
    },
    "0404/astro-ph0404459_arXiv.txt": {
        "abstract": "Spectra of the bright symbiotic star AG Draconis (BD+67$^{\\circ}$922) in the wavelength range 905--1187\\,\\AA\\ obtained with the \\emph{Far Ultraviolet Spectroscopic Explorer} (\\fuse) are presented. The spectra show a number of narrow, nebular emission lines, together with a uniform continuum from the hot component of the system, and numerous interstellar absorption lines. We infer the existence of \\ion{Ne}{viii} in the AG Dra nebula through the identification of the \\ion{Ne}{vii} 973.3\\,\\AA\\ recombination line. The emission line spectrum is dominated by intense lines of \\ion{O}{vi}, but also shows weaker lines from highly-ionized ions including \\ion{Ne}{v}, \\ion{Ne}{vi}, \\ion{S}{iv}, \\ion{S}{vi}. Members of the \\ion{He}{ii} Balmer series can be identified up to n=20. Lines of \\ion{Fe}{ii} and \\ion{Fe}{iii} fluoresced by \\ion{O}{vi} \\lam1032 are identified at wavelengths 1141.172\\,\\AA\\ and 1142.429\\,\\AA, respectively. The emission lines are shown to be produced in a plasma with an electron temperature of 20--30,000~K, photoionized by the white dwarf. The \\ion{Ne}{vi} \\lam997/\\lam999 ratio shows that this ion and all others except perhaps \\ion{Ne}{vii} are formed at least 300 white dwarf radii from the white dwarf. Revised wavelengths for the \\ion{Ne}{v} $2s^22p^2$ $^3P_{0,1}$ -- $2s2p^3$ $^5S_2$ and \\ion{Ne}{vi} $2s^22p$ $^2P_J$ -- $2s2p^2$ $^4P_{J^\\prime}$ transitions are published. ",
        "introduction": "Symbiotic stars are binary stars that are identified spectroscopically as typically having three components: a giant, a hot star (usually a white dwarf), and a circumstellar nebula, ionized by the hot component's radiation field.  Accretion of material from the giant's wind onto the white dwarf can lead to eruptions that may last several years and raise the system's brightness by several magnitudes. AG Draconis (BD+67$^{\\circ}$922) is one of the brightest symbiotics at UV and  X-ray wavelengths and undergoes outbursts in 3--6 year long intervals that repeat every 15 years or so. The  observations described here took place during quiescence.  The components of AG Dra are a K0--4 giant and a  hot star, most likely a white dwarf \\citep{miko95}. For convenience we will refer to the hot component as a white dwarf for the remainder of this article. The temperature of the white dwarf  is derived through fitting a blackbody curve to the continuum observed at X-ray and UV wavelengths, although contrasting results are found: \\citet{greiner97} finding a temperature of 170,000~K from ROSAT X-ray spectra, while \\citep{miko95} finding 100,000~K from IUE spectra.  The system is a non-eclipsing binary with an orbital period of $\\approx$550~days \\citep{mein79, galis99, fekel00}. The radial velocity was found by \\citet{demed99} to be $-147.46\\pm 3.54$~\\kms\\ and they give the rotational velocity of the giant as $5.9\\pm 1.0$~\\kms. The abundances of the giant are well determined \\citep{smith96} and demonstrate that it is metal deficient and possibly a barium star. The distance to AG Dra is crucial in fixing many of the properties of the system, yet it is rather uncertain, with the \\emph{Hipparcos} satellite only yielding a lower limit of 1~kpc. \\citet{miko95} argue from a classification of the giant for a distance of 2.5~kpc, which is the value we adopt in the present work.  AG Dra is a particularly attractive target in the far ultraviolet on account of the low extinction \\citep[E$_{B-V}$=0.05][]{miko95} along its sightline, resulting from the high latitude of the system ($b=41^\\circ$). Spectra from the International Ultraviolet Explorer (\\iue) showed strong emission lines from highly ionized species including \\ion{He}{ii} \\lam1640, \\ion{C}{iv} \\lam\\lam1548,1550 and \\ion{N}{v} \\lam1238,1242 \\citep{viotti83}. The first reported observations for wavelengths below 1200\\,\\AA\\ were made by \\citet{schmid99} who identified strong emission lines from \\ion{O}{vi} \\lam1032,1038 and \\ion{He}{ii} \\lam1085.  We report here observations made with the \\emph{Far Ultraviolet Spectroscopic Explorer} (\\fuse) in the wavelength region 905--1187\\,\\AA\\ made in 2000 March and 2001 April during a quiescent period of the system. These spectra are at a higher resolution and sensitivity than any previous spectra in this range and reveal a number of new plasma diagnostics available for interpreting the nebular emission from the system. ",
        "conclusions": "\\label{sect.discussion}  The previous sections have presented \\fuse\\ spectra obtained in 2000~March and 2001~April during a quiescent phase of AG Dra's $\\approx$15 year outburst cycle. These represent the first high resolution, high sensitivity spectra obtained in the 905--1187\\,\\AA\\ spectral region for AG Dra. The high ionization states found in the \\fuse\\ bandpass reveal significant new information about the system, unavailable from earlier, longer wavelength IUE spectra. In particular, the presence of the \\ion{Ne}{vii} \\lam973 recombination line implies the existence of \\ion{Ne}{viii} in the nebula, the highest ionization species ever recorded for AG Dra. The \\ion{O}{vi} \\lam1032 line has a P Cygni profile, demonstrating a wind in the high ionization part of the nebula during quiescence. Previously P Cygni profiles had only been seen in IUE spectra obtained during outbursts. The \\ion{Ne}{vi} intercombination lines are sensitive to the radiation field of the white dwarf and their ratios imply that the \\ion{Ne}{vi} emitting region (and thus also the \\ion{O}{vi} emitting region) is at least 300 white dwarf radii (or 0.14~AU for $R_{\\rm wd}\\approx 0.10~R_\\odot$) from the white dwarf. Evidence for the \\ion{O}{vi} being formed far from the white dwarf is also found from the \\ion{Fe}{ii} and \\ion{Fe}{iii} fluorescence lines: the \\ion{O}{vi} resonance lines are capable of fluorescing iron lines in the giant's atmosphere, whereas \\ion{H}{i} Ly\\al\\ is not, suggesting \\ion{O}{vi} is formed close to the giant's surface, while Ly\\al\\ is formed much further out.  The classical view of symbiotic nebulae has material flowing out from the giant and being ionized in the vicinity of the white dwarf. High ionization species (such as \\ion{He}{ii} and \\ion{O}{vi}) are found in the region around the white dwarf \\citep[e.g.][]{fried83}, while low ionization species are formed at increasingly further distances. An alternative theory is that the strong nebular lines in high density systems such as AG Dra are formed in the illuminated part of the giant's atmosphere, helping to explain the orbital modulation of the UV line fluxes found from \\iue\\ \\citep[e.g.,][]{altamore81}. \\citet{proga96,proga98} have created non-LTE photoionization models for the ionization of a giant's atmosphere in the case where it is illuminated by a hot companion star and have demonstrated that the general features of symbiotic star spectra can be reproduced with high effective temperatures and luminosities of the companion. The large emission line fluxes of highly-ionized species can only be reproduced if the giant has a wind extending to 2--3 giant radii.  \\begin{figure}[h] \\epsscale{0.7} \\plotone{f11.eps} \\caption{Predictions for the location of the \\ion{O}{vi}, \\ion{Ne}{vi}, and \\ion{Ne}{viii} ions in relation to the white dwarf, assuming a simple model for the ionization structure of the nebula (App.~\\ref{model}), and a white dwarf black body temperature of 170,000~K.\\label{dist-o6-ne6}} \\end{figure}  \\begin{figure}[h] \\epsscale{0.7} \\plotone{f12.eps} \\caption{A schematic diagram illustrating the locations of the \\ion{O}{vi} and \\ion{Ne}{viii} emitting regions (vertical lines) suggested by the present work.\\label{cartoon}} \\end{figure}  Evidence for the \\ion{O}{vi} lines being formed in the giant's atmosphere comes from the aforementioned \\ion{Ne}{vi} result and Fe fluorescence lines. We can also consider the location of the \\ion{Ne}{vi, viii} and \\ion{O}{vi} ions in the system through ionization balance arguments. The high ionization potential of these ions means that they are produced through irradiation from the X-ray section of the white dwarf's radiation field. Fits to the measured X-ray spectrum \\citep{greiner97} are consistent with a blackbody field of temperature 170,000~K. By constructing a simple model of the ionization/recombination processes in the plasma (Appendix~\\ref{model}) we can calculate the fractional population of \\ion{Ne}{vi} and \\ion{O}{vi} as a function of distance from the white dwarf (Fig.~\\ref{dist-o6-ne6}). The distance here corresponds to the inner face of a nebular shell, and the model assumes that the radiation field is unattenuated between the white dwarf and inner face. It is clear that such a hot blackbody prevents \\ion{O}{vi} and \\ion{Ne}{vi} being formed out to distances of 0.5--1.0 AU. For the model of AG Dra presented in Fig.~9 of \\citet{miko95} this places the \\ion{O}{vi} and \\ion{Ne}{vi} regions close to the surface of the giant, and well away from the white dwarf. \\ion{Ne}{viii}, however, must be formed much closer to the white dwarf. A cartoon illustrating the structure of the AG Dra nebula is shown in Fig.~\\ref{cartoon}. A difficulty for this particular model lies in the P Cygni profile of the \\ion{O}{vi} line, since the continuum in the far ultraviolet is due to the white dwarf and not the giant: for white dwarf continuum photons to be absorbed by \\ion{O}{vi} at orbital phase 0.5 (white dwarf in front of the giant), the \\ion{O}{vi} ions are required to lie in front of the white dwarf along the observer's line of sight. We note however, that the electron scattered wings around the \\ion{O}{vi} lines significantly enhance the continuum (Fig.~\\ref{o6-broad-wings}), and that the P Cygni absorption extends down to around $3\\times 10^{-13}$~\\ecs, the level of the white dwarf continuum, and not to zero flux. We thus suggest that the \\ion{O}{vi} ions are absorbing the electron-scattered \\ion{O}{vi} photons and \\emph{not} the stellar continuum. This can be understood if the electron scattering is occurring in the densest part of the \\ion{O}{vi} region close to the giant, with the absorption of the scattering wings then occuring in the lower density, extended wind. The high density of $10^{11}$~cm$^{-3}$ derived from the analysis of the wings in Sect.~\\ref{sect.o6-wings} supports this idea.   A full understanding of the structure of the AG Dra nebula can only come through detailed modelling of the ionization and wind structure, constrained by high quality spectral data."
    },
    "0404/astro-ph0404173_arXiv.txt": {
        "abstract": "Gravitational waves and X-ray flares are expected from tidal disruption of  stars by a massive black hole. Using a relativistic smoothed particle hydrodynamics code, we investigate the fate of main sequence and Helium stars in plunge orbits passing near Schwarzschild or a Kerr black holes of  mass $\\sim 10^{5-6} M_\\odot$. We show that quadrupole gravitational waves emitted during the tidal disruption process are described reasonable well by a point particle approximation even in the strong encounter case. An additional hydrodynamic calculation based on the Godunov method indicates that shocks develop for sufficiently high tidal compressions. The shock-heating results in an X-ray flare, which for solar-type stars disrupted by $\\sim 10^6M_\\odot$ black holes is in the keV range, associated with the gravitational wave signal. The hardness and duration of the X-ray flare may serve as a diagnostic of the mass of the central black hole. ",
        "introduction": "There is strong evidence, based on stellar kinematics and gas dynamics, for the existence of massive black holes with masses from a few times $10^5$ to a few times $10^9 M_\\odot$ in the center of a substantial fraction of galaxies (e.g. Kormedy \\& Gebhardt 2001 ). Gravitational radiation caused by stars orbiting closely around such black holes is one of the most promising candidates for detection by the Laser Interferometer Space Antenna (LISA; e.g. Danzmann et al. 1998; Hils \\& Bender 1995; Sigurdsson \\& Rees 1997; Ivanov 2002; Freitag 2003). Compact stars (white dwarfs, neutron stars, or stellar mass black holes) in bound orbits relatively close to the black hole will emit gravitational waves as they spiral all the way down to the horizon. Although the detection of such waves requires a set of accurate wave templates and significant computational power to search the large dimensionality parameter space, tracking the signals over a large number of cycles boosts the strength of the signals, hence previous studies have mainly addressed the gravitational waves from such stars.  On the other hand, there is a larger reservoir of main sequence stars, which are affected by the gravitational field of the  massive black hole already at large distances, where their gravitational  wave emission is still negligible. Of these, only stars in almost radial orbits can approach the horizon, where they can emit strong gravitational waves, as well as being subject to tidal disruption. The rate of captures leading to disruption may be even larger for He stars. Since a single periastron passage near the tidal radius can disrupt the star, the detectable gravitational wave signal takes the form of a burst. However, the gravitational waves are expected also to be associated with X-ray/UV emission originating from the tidal compressional heating and disruption of the star. The detection of the photon flares provides the means for pinpointing the direction, occurrence time and possibly the redshift of the event, as well as for obtaining independent evidence for the existence and properties of the massive black hole. The temporal coincidence of gravitational wave and electromagnetic signals could enhance the confidence level of the detection of both signals (e.g. Finn, Mohanty \\& Romano 1999). Multi-window observations (gravitational and electromagnetic emissions) might allow us to set tighter constants on the properties of both the massive black hole and the star.  Here we study the gravitational and electromagnetic radiation from the tidal disruption of a main sequence star and a Helium star by a massive black hole. We use a three dimensional relativistic smoothed particle hydrodynamic (SPH) code (Laguna, Miller \\& Zurek 1993), and investigate the gravitational wave and photon emission accompanying the tidal disruption event. In \\S 2, we discuss the main features of the tidal disruption process and the photon flares expected during the disruption process. In \\S 3 we provide a rough estimate of the characteristics of the gravitational waves. In \\S 4 we discuss the numerical results obtained from relativistic SPH calculations around both Schwarzschild and maximally rotating Kerr black holes. In \\S 5 the detectability of the gravitational waves, and the constraints obtainable from photon and gravitational wave observations are discussed. Conclusions and discussions are given in \\S 6. In the appendix, we discuss the shock formation during the tidal compression, using a relativistic Lagrangian code based on the Godunov method (Kobayashi, Piran \\& Sari 1999). ",
        "conclusions": "We have numerically investigated the evolution of solar-type and He-type stars which are captured and disrupted by massive black holes with mass of $10^{5-6}M_{\\odot}$. We have shown that for both types of stars, plunge orbits which approach the horizon produce gravitational wave bursts with large amplitudes, which may be detectable by LISA up to $D\\sim 20$ Mpc. Even though the stars undergo significant disruption, the waveforms of the gravitational signal can be reasonably well approximated by a point particle approximation. We have also shown that the gravitational wave signal during the disruption process is largely insensitive to the equation of state and to the detailed inner structure of the star. For the same penetration parameter $\\beta_p$ (ratio of tidal to periastron radius), the waveforms of main-sequence and He stars differ considerably in shape, although deep penetration (e.g. $\\beta_p=10$) main-sequence waveforms have some resemblance to critical penetration ($\\beta_p=1$) He-star waveforms. For large Kerr rotation parameters, the waveform shape and polarization is sensitive to the black-hole rotation and its sign. In addition to long-term (days to years) X-ray flares, we predict also a prompt X-ray flash, of luminosity $L_x\\sim 10^{42}$ erg/s and duration $\\sim 10$ s, associated with the gravitational wave bursts. The characteristic spectral temperature can give constraints on the black hole mass.  We thank M. Rees, S. Sigurdsson, T. Alexander, S. Mahadevan and the referee for valuable comments. This work is supported by NASA NAG5-13286, NASA NAG5-10707, NSF AST 0098416, the Monell Foundation, the Merle Kingsley fund, and the Pennsylvania State University Center for Gravitational Wave Physics, funded under cooperative agreement by NSF PHY 01-14375.  \\appendix"
    },
    "0404/astro-ph0404490_arXiv.txt": {
        "abstract": "The\\unmarkedfootnote{Corresponding author: m.casse@cea.fr \\\\ $^{3}$ Institut d'Astrophysique de Paris, 98 bis Boulevard Arago, 75014 Paris, France \\\\ $^{4}$ F\\'ed\\'eration de Recherche Astroparticule et Cosmologie, Coll\\`ege de France, 11 Place Marcellin Berthelot, 75231 Paris, France } nature of Dark Matter remains one of the outstanding questions of modern astrophysics. The success of the Cold Dark Matter cosmological model argues strongly in favor of a major component of the dark matter being in the form of elementary particles, not yet discovered. Based on earlier theoretical considerations, a possible link between the recent SPI/INTEGRAL measurement of an intense and extended emission of 511 keV photons (the hallmark of positron annihilation) from the central Galaxy, and this mysterious component of the Universe, has been established advocating the existence of a {\\sl light} dark matter (LDM) particle (at variance with the neutralino, in general considered as very {\\sl heavy}).  \\vspace{-2mm}  We show that it can explain the 511 keV emission mapped with SPI/INTEGRAL without overproducing undesirable signals like high energy gamma-rays arising from $\\pi^\\circ$ decays, and radio synchrotron photons emitted by high energy positrons circulating in magnetic fields. Combining the annihilation line constraint with the cosmological one (i.e. that the relic LDM energy density reaches about 23\\% of the density of the Universe), one can restrict the main properties of the light dark matter particle. Its mass should lie between $\\,\\approx$ 1 and 100 MeV, and the required annihilation cross section, velocity dependent, should be significantly larger than for weak interactions, and may be induced by the virtual production of a new light neutral spin 1 boson $U$.  \\vspace{-2mm}  On astrophysical grounds, the best target to validate the LDM proposal seems to be the observation by SPI/INTEGRAL and future gamma ray telescopes of the annihilation line from the Sagittarius dwarf galaxy and the Palomar-13 globular cluster, thought to be dominated by dark matter.  \\vspace{-1mm} ",
        "introduction": "In the currently popular Cold Dark Matter (CDM) cosmological model with a cosmological constant $\\Lambda$, the Universe is composed of $\\,\\approx 4\\,\\%$ baryonic matter and $\\approx 23\\,\\%$ cold, collision-less dark matter, and is made flat by a cosmological constant \\citep{spergel2003}. The very nature of dark matter is one of the tantalizing and yet unanswered questions of contemporary astrophysics. While its existence has been inferred for several decades, its amount has been clearly assessed on various astrophysical scales only in the recent past. The timed questions are its distribution and above all its true nature. The precise knowledge of its density profile is likely to be decisive in answering the ultimate question of the identification of the kind of particle(s) composing dark matter. The success of the CDM cosmology, combined with dark energy ($\\Lambda$), argues strongly for a major component of the dark matter being in the form of fundamental (supersymmetric?) particles.   One possible way of probing the nature of dark matter particles is to look for their annihilation signal (annihilation of dark matter particles on their anti-particles, knowing that both could be identical). The best regions to search for this annihilation radiation is where dark matter accumulates, i.e. in deep gravitational potential wells. The annihilation rate is proportional to the square of the number density of the dark matter particles (in the case where $dm=\\overline{dm}$). The Galactic Center and its surroundings are therefore obvious targets because of their density and proximity.   A wide literature exists on the prospects of observing dark-matter-induced gamma rays from the Galactic Center and clumps of dark matter in the Galactic halo (\\cite{bertone2002} and references therein). But these gamma rays are at high energy, outside the INTEGRAL grasp, and are prime targets for high energy gamma ray satellite experiments (like GLAST), and atmospheric Cherenkov telescopes (such as HESS, VERITAS, MAGIC and CANGAROO). In the following, however, we advocate a link between dark matter and {\\it \\,low energy\\,} gamma rays, within the reach of INTEGRAL.  511 keV photons are the signature of electron-positron annihilations. Recent observations of 511 keV photons from the central region of the Galaxy, performed by the SPI spectrometer on board of the INTEGRAL satellite \\citep{ jean2003,knodlseder2003,jean2004,weidenspointner2004} confirm previously reported observations \\citep{leventhal1993, milne2002} and improve the spatial resolution. The SPI/INTEGRAL data indicate that the 511 keV emission zone is spherically symmetric, best fitted by a gaussian distribution with 9$^\\circ$ (FWHM), and inconsistent with a single point source (although it is premature to exclude multiple point sources). The large extension of the emission region around the Galactic Center combined with the surprising failure of the leading candidates (SN~Ia), as shown in \\cite{schanne2004}, opens the way to new speculations. A possible link between the 511 keV emission line mapped by SPI/INTEGRAL and a new kind of dark matter, bosonic and light, that had been discussed previously \\citep{boehm2003a,boehm2003b} was proposed by \\cite{boehm2004a}, since other, more conventional, explanations are questionable. ",
        "conclusions": "Summarizing, the 511 keV emission from the Galactic bulge observed by SPI/INTEGRAL could be the signature of light ($\\,\\approx$ 1 -- 100 MeV) dark matter particles, with a preference for the lower end of this mass interval. Such particles could annihilate throughout the Galactic bulge into positrons (and electrons) which, after being stopped, themselves annihilate into 511 keV photons. The link between the positron annihilation line from the central Galactic regions and dark matter, if real, would push INTEGRAL at the forefront of astroparticle research. It is thus essential to further test this hypothesis, and to discard, if possible, competing ones using INTEGRAL itself as well as other means. For this purpose, it is crucial to get a detailed map of the 511 keV emission in the central region of the Galaxy."
    },
    "0404/astro-ph0404203_arXiv.txt": {
        "abstract": "{ It is shown that the nonlinear kinetic theory of cosmic ray (CR)  acceleration in supernova remnants (SNRs) describes the shell-type nonthermal X-ray morphology of Cas~A, obtained in Chandra observations, in a satisfactory way. The set of empirical parameters, like distance, source size and total energy release, is the same which reproduces the dynamical properties of the SNR and the spectral characteristics of the emission produced by CRs. The extremely small spatial scales in the observed morphological structures at hard X-ray energies are due to a large effective magnetic field $B_\\mathrm{d}\\sim 500$~$\\mu$G in the interior which is at the same time not only required but also sufficient to achieve the excellent agreement between the spatially integrated radio and X-ray synchrotron spectra and their calculated form. The only reasonably thinkable condition for the production of such a large effective field strength is a very efficiently accelerated nuclear CR component. Therefore the Chandra data confirm first of all the inference that Cas~A indeed accelerates nuclear CRs with the high efficiency required for Cas~A to be considered as a member of the main class of Galactic CR sources and, secondly, that the nonlinear kinetic theory of CR acceleration in SNRs is a reliable method to determine the magnetic field value in SNRs. ",
        "introduction": "Supernova remnants (SNRs) are the main sources of energy for the Interstellar Medium (ISM). They control the physical state of the ISM, and that presumably includes the nonthermal component of Interstellar matter, often called the Galactic Cosmic Rays (CRs). Quantitative analysis of the SNR dynamics and its energetic particle production requires the determination of the effective magnetic field in SNRs. For this we note that the magnetic field decisively influences the acceleration of CRs and their subsequent dynamics in SNRs. In the following we shall present evidence for the case of Cassiopeia~A (Cas~A) that the nonlinear kinetic theory for CR acceleration in SNRs (Berezhko et al. \\cite{byk96}; Berezhko \\& V\\\"olk \\cite{bv97}; \\cite{bv00}) does not only satisfactorily explain the observed SNR dynamics and the properties of the nonthermal emission, but that it also permits to determine the value of the effective magnetic field in SNRs, as already demonstrated for the case of SN~1006 (Berezhko et al. \\cite{bkv02}; \\cite{bkv03}).  The existence of a powerful population of relativistic particles in SNRs is known from their nonthermal radiation. For some SNRs this emission is detected in a very wide range from the radio to the \\gr band. Cas~A represents such an example. It is a shell-type SNR, and a bright source of synchrotron radiation observed from radio waves (e.g. Baars et al. \\cite{baars}) to hard X-rays (Allen et al. \\cite{all97}). These emissions are direct evidence for the existence of a large number of relativistic electrons presumably accelerated at the shock, with energies reaching about 10~TeV. Direct experimental knowledge of nuclear CR production can only be obtained by high energy \\gr measurements. If protons are efficiently accelerated in Cas~A, then the $\\pi^0$-decay $\\gamma$-ray spectrum as a result of inelastic collisions with the background nuclei must extend beyond 1~TeV with a hard power-law spectrum. The detection of a corresponding if weak signal in TeV $\\gamma$-rays has been indeed reported by the HEGRA collaboration (Aharonian et al. \\cite{aha01}).  According to the analysis by Borkowski et al. (\\cite{bor96}) which is consistent with the observed dynamics of Cas~A, the supernova shock expands into a non-uniform circumstellar medium strongly modified by the intense progenitor star wind. This circumstellar environment consists of a tenuous inner bubble, created by the Wolf-Rayet wind, a dense shell of swept-up slow Red Supergiant (RSG) wind material, and a subsequent unperturbed RSG wind. The medium further out is not relevant here.  Since the shock has already passed through the dense shell and has swept up a mass of about 5$M_{\\odot}$ (Favata et al. \\cite{favata}), a detectable $\\pi^0$-decay \\gr flux is to be expected. Recently Chevalier \\& Oishi (\\cite{chevo}) have argued that a Wolf-Rayet phase and the corresponding modification of the preexisting RSG wind is not required. To reach a final conclusion about this important issue a detailed dynamical and kinetic consideration has to be made, similar to the one contained in our previous paper (Berezhko et al. \\cite{bpv03}). At the moment we can only say that the global characteristics of nonthermal radiation of Cas~A are expected to be not sensitive to these details of the structure of circumstellar medium, because they are mainly determined by the total amount of the swept up gas and the magnetic field value in the remnant. Since the observed hard X-ray synchrotron emission of Cas~A comes from the thin region near the current SN shock position, the properties of this radiation are not expected to be very sensitive to the previous history of the SNR evolution. Here we use our model (Berezhko et al. \\cite{bpv03})  which agrees with all previous data, leaving the question about the role of Wolf-Rayet stage for future scrutiny, and test its consistency with the hard X-ray morphology.  Comparing the nonthermal radio and X-ray synchrotron data with the calculated spectrum of the energetic electrons, Berezhko et al. (\\cite{bpv03})  have inferred that the existing data require very efficient acceleration of {\\it CR nuclei} at the SNR blast wave which has converted already several percent of the initial SNR energy content into nuclear CR energy despite the necessity of renormalization of the overall nuclear CR content due to the nonuniform injection of suprathermal ions (V\\\"olk et al. \\cite{vbk03}). At the same time a large interior magnetic field strength $B_\\mathrm{d} \\approx 500 \\mu$G is required. The large value of $B_\\mathrm{d}$ is expected and possible in this case, using the spatially integrated radio and X-ray synchrotron spectra to empirically determine the effective magnetic field strength and the injection rate in the theory (see also V\\\"olk \\cite{voel03}).  In fact there are very recent measurements of details of the X-ray morphology of Cas~A with the Chandra telescope by Vink \\& Laming (\\cite{vl03}).  These authors have used the detailed Chandra morphology to estimate the magnetic field strength by assuming the small spatial X-ray scale to be determined by postshock synchrotron losses. Our aim is to show that the actual B-field strength coincides with that inferred from the radio spectrum and nonlinear acceleration theory. We confirm that the field is indeed amplified in Cas~A, although it turns out that the actual value is significantly larger than estimated by Vink \\& Laming (\\cite{vl03}). ",
        "conclusions": "Compared with SN~1006 which exploded into an essentially uniform ISM, Cas~A is a very different and much more complex object. The dynamics of the SN shock and the major fraction of the emission produced by accelerated CRs are predominantly determined by the parameters of the swept-up RSG shell. They depend less sensitively on the parameters of the medium near the current shock position since that presumably consists of free RSG wind material and contains little mass. Nevertheless, since the required magnetic field in the interior downstream of the shock is so large that it strongly influences the energy spectrum of the CR electrons as a result of their synchrotron losses (Berezhko et al. \\cite{bpv03}), the hard X-ray synchrotron emission should come predominantly from the relatively thin postshock region, because deeper inside the remnant the upper cutoff energy of CR electrons becomes too low. Therefore one would expect that the nonlinear kinetic theory should not only reproduce the spectral characteristics of Cas~A in a satisfactory way, but should also be able to explain the observed fine structure of the nonthermal emission in hard X-rays.  For the comparison with the Chandra data we shall not present the details of our model. They have already been described in the above paper (Berezhko et al. \\cite{bpv03}). In order to give a clear interpretation and simple estimate of the spatial variation of the synchrotron emission near the SNR shock we shall use some simple analytical approximations, extending the discussion of SN~1006 (Berezhko et al. \\cite{bkv03}). As a second step we shall make a comparison of the Chandra data with the results of the nonlinear theory.  It was shown (Berezhko et al. \\cite{bkv03}) that due to the strong synchrotron losses the very energetic electrons occupy only a thin region behind the shock of thickness \\begin{equation} l_2^{-1}=(2l_\\mathrm{d2})^{-1}(\\sqrt{1+4l_\\mathrm{d2}/l_\\mathrm{c2}}-1), \\label{eq1} \\end{equation} where $l_\\mathrm{d}(p)=\\kappa(p)/u$ and $l_\\mathrm{c}(p)=\\tau(p) u$ are the diffusive and convective length scales respectively, $\\tau(p)= 9 m_\\mathrm{e}^2c^2/(4 r_0^2B^2p)$, is the synchrotron loss time, $m_\\mathrm{e}$ is the electron mass, $r_0$ denotes the classical electron radius, $B$ is the effective magnetic field strength, $\\kappa(p)=\\rho_\\mathrm{B} v/3$ is CR diffusion coefficient, and $\\rho_\\mathrm{B}(p)$, $p$ and $v$ are the particle gyroradius, momentum and velocity, respectively. In addition $u_2=V_\\mathrm{s}/\\sigma$ is the downstream plasma speed relative to the shock front, $V_\\mathrm{s}$ is the shock speed,  and $\\sigma$ is the total shock compression ratio.  Depending upon the ratio $l_\\mathrm{d}/l_\\mathrm{c}$ we have two different extreme cases. For electron momenta $p$ corresponding to $l_\\mathrm{c}\\gg 4l_\\mathrm{d}$ (weak losses) we have $l_2=l_\\mathrm{c}$. In this case the loss time $\\tau$ is much larger than the acceleration time $\\tau_\\mathrm{acc}\\sim \\kappa/u^2$ and therefore the losses do not influence the electron spectrum during their acceleration. The electron spectrum has a power law character up to a cutoff momentum $p_\\mathrm{max}>p$, which exceeds the considered momentum $p$. Later on, the losses in the downstream region lead to a decrease of the cutoff momentum so that at the distance $r=R_\\mathrm{s}-l_\\mathrm{c}$ it drops to the value $p_\\mathrm{max}(r)=p$. Here $R_\\mathrm{s}$ is the shock radius. In the opposite extreme case of strong losses $l_\\mathrm{c}\\ll 4l_\\mathrm{d}$, we have $l_2=\\sqrt{\\kappa_2 \\tau_2}$. which is in fact independent of momentum if particle diffusion proceeds in the Bohm limit. In this case losses become significant already during the electron acceleration and therefore the cutoff momentum $p_\\mathrm{max}$ is lower than $p$, so that the corresponding electrons belong to the exponential tail of the distribution. Since the nonthermal X-ray emission of Cas~A belongs to the steep tail of the synchrotron spectrum, Eq.(\\ref{eq1}) suggests that $l_\\mathrm{c} \\ll l_\\mathrm{d}$ for the radiating electrons.  Since the synchrotron data are given in terms of the emission frequency $\\nu$, it is useful to rewrite Eq.(\\ref{eq1}), taking into account the approximate relation $\\nu\\propto Bp^2$: \\begin{equation} l_2=\\sqrt{\\kappa_2 \\tau_2}/(\\sqrt{1+\\delta^2}-\\delta), \\label{eq2} \\end{equation} where \\begin{equation} \\delta^2=l_\\mathrm{c2}/(4l_\\mathrm{d2})=0.12[c/(r_0\\nu)][V_\\mathrm{s}/(\\sigma c)]^2. \\label{eq3} \\end{equation}  We then have the situation that the synchrotron emissivity $q_\\mathrm{\\nu}(\\epsilon_\\mathrm{\\nu},r)$ is everywhere close to zero except for a thin radial region of thickness $\\Delta r\\sim l_2$ just behind the SN shock, since the emissivity from the upstream region $r>R_\\mathrm{s}$ is insignificant.  In projection along the line of sight, the radial emissivity profile determines the remnant's surface brightness $J_\\mathrm{\\nu}$. For the X-ray energy $\\epsilon_\\mathrm{\\nu}$ it has the form \\begin{equation} J_\\mathrm{\\nu}(\\epsilon_\\mathrm{\\nu},\\rho)= 2\\int_a^0 dx q_\\mathrm{\\nu}(\\epsilon_\\mathrm{\\nu},r=\\sqrt{\\rho^2+x^2}), \\label{eq4} \\end{equation} where $\\rho$ is the distance between the center of the remnant and the line of sight. In the case of a spherical shock $a=-\\sqrt{R_\\mathrm{s}^2-\\rho^2}$.  To estimate the shape of the expected profile $J_\\mathrm{\\nu}(\\rho)$ one can consider the simple case, when the emissivity has a form \\begin{equation} q_\\mathrm{\\nu}(\\epsilon_\\mathrm{\\nu},r)=q_2(\\epsilon_\\mathrm{\\nu}) \\exp[(r-R_\\mathrm{s})/l_2] \\label{eq5} \\end{equation} for $r\\le R_\\mathrm{s}$. In the case of interest, $l_2\\ll R_\\mathrm{s}$, the main contribution to the integral (\\ref{eq4}) comes from the thin regions near the boundary $x=a$. To perform the integration along the line of sight one can therefore expand the quantity $r=\\sqrt{\\rho^2+x^2}$ in Eq.(\\ref{eq5}) in powers of $x-a$ up to the second order term which results in: \\[ J_\\mathrm{\\nu}=\\frac{2q_2R_\\mathrm{s}l_2}{\\sqrt{R_\\mathrm{s}^2-\\rho^2}} \\left\\{ 1- \\frac{R_\\mathrm{s}^2-2\\rho^2}{R_\\mathrm{s}^2z}\\right. - \\] \\begin{equation} \\hspace{0.5cm} \\left. -e^{-z} \\left[ 1- \\frac{R_\\mathrm{s}^2-2\\rho^2}{R_\\mathrm{s}^2z} +\\frac{(z+2)}{2} \\left(\\frac{2\\rho^2-R_\\mathrm{s}^2}{R_\\mathrm{s}^2}\\right) \\right] \\right\\}, \\label{eq6} \\end{equation} where $z=(R_\\mathrm{s}^2-\\rho^2)/(R_\\mathrm{s}l_2)$. This approximate expression is accurate to about one percent for the case of a small thickness of the emission layer $l_2\\ll 0.1R_\\mathrm{s}$. As it is clear from this expression the distant observer will see the profile $J_\\mathrm{\\nu}(\\rho)$ which, going from large distances to the center of the SNR, starts at $\\rho=R_\\mathrm{s}$, reaches its peak value $J_\\mathrm{max}=2q_2\\sqrt{R_\\mathrm{s}^2-\\rho_m^2}$ at $\\rho_\\mathrm{m}\\approx R_\\mathrm{s}-0.85l_2$ and smoothly drops to $J_\\mathrm{min}=2l_2q_2$ at the SNR center $\\rho=0$. Therefore the outer scale $L_1\\approx l_2/2$ characterizes the emission profile at $\\rho>\\rho_\\mathrm{m}$, whereas for the inner scale, which characterizes the profile at $\\rho<\\rho_\\mathrm{m}$, we have $L_2\\approx 13 L_1$. This means that the width $L=L_1+L_2\\approx 7l_2$ of the observed brightness profile $J_\\mathrm{\\nu}(\\rho)$ is always appreciably wider than the width $l_2$ of the emissivity layer $q_\\mathrm{\\nu}(r)$, simply due to the geometry of the system.  According to (\\ref{eq2}) the downstream magnetic field $B_\\mathrm{d}$ is \\begin{equation} B_\\mathrm{d}=[3m_e^2c^4/(4er_0^2l_2^2)]^{1/3}(\\sqrt{1+\\delta^2}-\\delta)^{-2/3}. \\label{eq7} \\end{equation} The case of strong losses $l_\\mathrm{c}\\ll 4l_\\mathrm{d}$ corresponds to $\\delta=0$. According to Vink \\& Laming (\\cite{vl03}) the total thickness of the observed profile at $\\epsilon_{\\nu}>2.7$~keV is as small as $1.5''$. However, following the above consideration, only the small fraction $1.5''/7\\approx 0.22''$ corresponds to the thickness of the downstream emission region $l_2$. For a distance $d=3.4$~kpc and an angular radius of $2.6'$ for Cas~A this gives $l_2\\approx 1.1\\times 10^{16}$~cm. In the case of strong losses $(\\delta=0)$ $B_\\mathrm{d}\\approx 470$~$\\mu$G, independently of synchrotron energy and shock speed. For  $V\\mathrm{s} = 3000$~km/s and $\\sigma=6$, corresponding to the overall nonlinear model (Berezhko et al. \\cite{bpv03}), $\\delta^2=0.054$ resulting in $B_\\mathrm{d}\\approx 550$~$\\mu$G. For a larger shock speed $V_\\mathrm{s} = 5000$~km/s, as derived from X-ray measurements and used by Vink \\& Laming, we have $\\delta^2= 0.15$ and $B_\\mathrm{d}\\approx 610$~$\\mu$G.  The corresponding magnetic field value is by a factor of about five larger than estimated by Vink \\& Laming (\\cite{vl03}). There are two reasons for such a large discrepancy (i) only a small part of the observed profile width $L$ corresponds to the thickness $l_2$ of the emission region (ii) the correct formula (\\ref{eq1}), cf. V\\\"olk et al. (\\cite{vmf81}), implies $l_\\mathrm{c2}<l_2$. Using $l_\\mathrm{c2}=L$, as was done by Vink \\& Laming, therefore doubly underestimates $B_\\mathrm{d}$, since $l_\\mathrm{c2}\\propto B_\\mathrm{d}^{-2}$.  The numerically calculated brightness profile $J_\\mathrm{\\nu}(\\rho)=\\int_{\\epsilon_1}^{\\infty}d\\epsilon_\\mathrm{\\nu} J_\\mathrm{\\nu}(\\epsilon_\\mathrm{\\nu},\\rho)$ for the X-ray energy $\\epsilon_\\mathrm{\\nu}>\\epsilon_1=2.7$~keV is shown in Fig.\\ref{f1}.  The sharpest experimental X-ray brightness profiles obtained by the Chandra observers (Vink \\& Laming \\cite{vl03}) are shown as a function of angular distance $\\psi =2.6\\times 60(\\rho-R_\\mathrm{s})/R_\\mathrm{s}$ in arcseconds, taking into account that the angular radial size of Cas~A is $2.6'$. Since the absolute values of the measurements are not known, the theoretical profile is normalized so as to obtain the best fit to the experimental profile. The very sharp part on the right side of the calculated profile $J_\\mathrm{\\nu}(\\psi)$ corresponds to the shock front position $\\rho=R_\\mathrm{s}$. Compared with Vink \\& Laming (\\cite{vl03}) all experimental points are shifted by $\\Delta \\psi =0.3''$ to the left. One can see that the experimental values agree quite well with our calculations, which yielded a postshock value of the magnetic field strength at the current evolutionary epoch (see Berezhko et al. \\cite{bpv03}) of $B_\\mathrm{d}=480$~$\\mu$G. \\begin{figure} \\centering \\includegraphics[width=7.5cm]{xfcasa.eps} \\caption{The X-ray emissivity for X-ray energies $2.7<\\epsilon_\\mathrm{\\nu}<9$~keV as a function of angular radial distance. The Chandra measurements are shown (Vink \\& Laming \\cite{vl03}).} \\label{f1} \\end{figure} Taking into account that the Chandra angular resolution is $0.5''$ we conclude that the actual downstream magnetic field in Cas~A is indeed $B_\\mathrm{d}\\approx 500$~$\\mu$G with an uncertainty of about 30\\% which encompasses the above simple estimates of the local field.  There are a number of reasons for a broadening of the observed profile, compared with the ideal case of a spherical shell. First of all, it is clear from the observations that the actual shock front deviates from a strictly spherical form. The wavy shape of the shock front can be produced as the result of small scale density inhomogeneities of the ambient ISM. Any small-scale distortion of the spherical emission shell leads to a broadening of the observed brightness profile. This explains why the observed X-ray radial profile is slightly broader than the theoretical profile (see Fig.1). Therefore it is hard to believe that the actual magnetic field in Cas~A could be significantly lower than the above value.  Since at the current epoch the supernova shock propagates through the free RSG wind, where the required effective magnetic field would be so strong that the Alfv\\'en speed is $c_\\mathrm{a}\\approx 160$~km/s (Berezhko et al. \\cite{bpv03}), this field exceeds the preexisting wind field by at least an order of magnitude, because the wind should presumably be superalfvenic, that is $c_\\mathrm{a}<V_\\mathrm{w}$, where $V_\\mathrm{w}\\sim 10$~km/s is the RSG wind velocity. From this argument, the required strong downstream magnetic field is inevitably the result of considerable amplification near the supernova shock. As argued for the case of SN~1006 (Berezhko et al. \\cite{bkv03}) such a strong magnetic field amplification can only be produced nonlinearly by a very efficiently accelerated {\\it nuclear CR component}. In this case their number, consistent with all existing data, is so high that they are able to strongly excite magnetohydrodynamic waves, and thus to amplify the magnetic field, and at the same time to permit efficient CR scattering on all scales, approaching the Bohm limit (Bell \\& Lucek \\cite{bluc01}). The same large effective magnetic field is required by the comparison of our selfconsistent theory with the synchrotron observations.  For such a high downstream magnetic field $B_\\mathrm{d}\\approx 0.5$~mG, which leads to strong synchrotron losses of the CR electron component, the leptonic contribution to the TeV \\gr emission is so small that it is far below the detected HEGRA flux (Berezhko et al. \\cite{bpv03})."
    },
    "0404/astro-ph0404035_arXiv.txt": {
        "abstract": "We present photometric and spectroscopic observations of SN 2000ch, an unusual and extremely luminous variable star located in the galaxy NGC 3432. The object was discovered on 2000 May 3.2 during the course of the Lick Observatory Supernova Search, at an unfiltered magnitude of about 17.4. Pre-discovery images obtained in 1997, 1998, and 2000 April show the object with $R = 19.2-19.5$ mag.  Optical spectra obtained beginning on 2000 May 6 show a smooth, flat continuum and strong, broad hydrogen Balmer emission lines at wavelengths consistent with the catalogued redshift of NGC 3432, strengthening the association of the variable with the galaxy. Photometric monitoring reveals a complex and erratic light curve over a time span of $\\sim$10 days. Subsequent optical spectra over the next $\\sim3$ months continued to show strong Balmer emission lines with a mean full-width at half-maximum intensity $\\sim 1550$ km s$^{-1}$ and a distinct red asymmetry.  A spectrum obtained 9 months after the outburst is similar to the previous spectra, but the integrated flux in H$\\alpha$ is nearly half that observed during the outburst. The object's photometric behavior, spectrum, and luminosity suggest that it is a very massive and luminous variable star and might be related to some luminous blue variable stars such as $\\eta$ Carinae and SN 1997bs in NGC 3627.  The brightest apparent magnitude implies an absolute magnitude of $M_V \\approx -12.7$ at the distance of NGC 3432, a value which is comparable to $\\eta$ Carinae during its outburst in the mid-nineteenth century. ",
        "introduction": "Recent systematic searches for supernovae (SNe) in external galaxies such as the Lick Observatory Supernova Search (LOSS; Li et al. 2000; Filippenko et al. 2001) have revealed not only a plethora of new SNe but many other luminous variable objects as well.  Some of these include a class of peculiar type II supernovae (SNe~II), designated SNe IIn, which exhibit hydrogen emission lines consisting of a broad complex base underlying a stronger, sometimes slightly blueshifted, narrow emission component with no trace of the broad P-Cygni profiles that are characteristic of normal SNe~II (Schlegel 1990; Filippenko 1991).\\footnote{Sometimes, however, a very narrow, low-velocity absorption component is visible; see SNe 1994W and 1994ak in Figure 14 of Filippenko (1997), for example.} SNe~IIn probably originate from the explosion of massive stars in which the high-speed ejecta and hard radiation field interact with the dense circumstellar envelope, created previously by extreme mass-loss from the progenitor.  It has become increasingly clear that some objects cataloged previously as peculiar SNe~II are actually SNe~IIn while others are not genuine SNe, defined to be the final demise of a star at the end of its life. The latter objects are now believed to result from outbursts of luminous, massive stars such as the class of luminous blue variables (LBVs).  Historical examples include SN 1954J in NGC 2403 (Humphreys \\& Davidson 1994; Smith, Humphreys, \\& Gehrz 2001), SN 1961V in NGC 1058 (Goodrich et al.  1989; Filippenko et al. 1995; Van Dyk et al. 2002), and SN 1997bs in NGC 3627 (Van Dyk et al. 2000).  An excellent review of the observational properties and challenges posed by LBVs is given by Humphreys \\& Davidson (1994).  LBVs are evolved, massive, and unstable hot stars that undergo sporadic, sometimes violent, mass-loss events and whose nature is not well understood. The group is not entirely homogeneous since there are stars which are luminous, blue, and slightly variable that are not LBVs.  In general, well-studied LBVs have bolometric absolute magnitudes brighter than $-9.5$, suggesting masses in excess of 50 M$_\\sun$. These stars exhibit relatively large mass-loss rates of 10$^{-5}$ to 10$^{-4}$ M$_\\sun$ yr$^{-1}$ which may be observed as an infrared excess from circumstellar ejecta.  An exceptional circumstellar envelope is the ``homunculus'' surrounding $\\eta$ Carinae.  LBVs exhibit eruptions of 1-2 mag on time scales of 10--40 years and may have superoutbursts of more than 2 mag on time scales of hundreds to thousands of years. Their optical spectra typically show emission lines of H and \\ion{He}{1}, as well as both permitted and forbidden lines of \\ion{Fe}{2}.  Often these lines are accompanied by P-Cygni profiles with typical expansion velocities of 100--200 km s$^{-1}$.  At minimum light, the spectra are typical of hot supergiants with $T_{\\sc eff} \\approx 12,000--30,000$ K, while at maximum the spectra resemble cooler supergiants and are characteristic of an optically thick extended photosphere with $T_{\\sc eff} \\approx 7000-8000$ K.  Examples in the Galaxy include P Cygni, $\\eta$ Carinae, S Doradus, and AG Carinae.  Their eruptions can be dramatic: $\\eta$ Carinae underwent an outburst between 1830 and 1860 in which it briefly reached $M_{\\sc Bol} \\approx -14$ mag.  In this paper we describe the discovery and evolution of a new luminous variable star designated SN 2000ch (Papenkova \\& Li 2000; Wagner et al. 2000a,b) which might be related to LBVs.  The object exhibits erratic photometric variations on relatively fast time scales compared to classical LBVs, but it has similar optical spectra when compared to some LBVs in outburst.  During the outburst, its luminosity and size appear comparable to that of $\\eta$ Carinae, while in quiescence it remains a luminous object and sufficiently bright for detailed study. ",
        "conclusions": "The spectral features observed early in the outburst of SN 2000ch are reminiscent of those observed in some M31 classical novae.  This suggested to Wagner et al. (2000a) that SN 2000ch might be related to a class of super-bright classical novae (Della Valle 1991), which have peak absolute magnitudes $\\sim$1 mag brighter than typical Galactic and M31 novae ($M_V \\approx -9$ mag). However, a bright-nova interpretation is inconsistent with the high luminosity ($10^{6-7}\\ L_\\sun$), the relatively low expansion velocity, and the fast timescale that we observed. In the study presented by Della Valle (1991), no classical nova has been observed in the Galaxy, M31, or the Virgo Cluster with a luminosity comparable to that of SN 2000ch.  Based on the early data presented here, Filippenko (2000) suggested that the variable star in NGC 3432 could be a subluminous SN~IIn (Schlegel 1990; Filippenko 1991) and thus perhaps related to the class of luminous blue variables (LBVs). Nearly all SNe IIn show hydrogen emission lines consisting of a broad complex base and a strong, narrow, and slightly blueshifted component superposed on a blue continuum.  Broad P-Cygni absorption components are not present (Schlegel 1990).  In addition, some SNe IIn have been observed in spiral galaxies in and near \\ion{H}{2} regions or large massive star-forming complexes.  Their absolute magnitudes at discovery are typically $-17.5$ to $-18.0$ and thus comparable to most SNe II (Schlegel 1990).  The light curves of SNe~IIn, however, are not well characterized.  Recently, it has become apparent that the SN~IIn class is quite heterogeneous (Filippenko 1997). It consists of objects which are certainly {\\it bona-fide} SNe~II in which core collapse has led to the destruction of the star. But as additional objects have been discovered whose characteristics resemble SNe IIn, it appears that some are in fact not true core-collapse SN but are the result of superoutbursts of luminous, massive stars.  Some of these stars are luminous blue variables on their way to becoming Wolf-Rayet stars. Objects in this class include $\\eta$ Carinae (Davidson \\& Humphreys 1997), SN 1954J in NGC 2403 (Humphreys \\& Davidson 1994), SN 1961V in NGC 1058 (Goodrich et al. 1989; Filippenko et al. 1995; Van Dyk et al. 2002), SN 1997bs in NGC 3627 (Van Dyk et al.  2000), SN 1999bw in NGC 3198 (Filippenko, Li, \\& Modjaz 1999), and SN 2002kg in NGC 2403 (Schwartz, Filippenko, \\& Chornock 2003).  All these objects are subluminous with respect to more normal SNe~IIn, but otherwise exhibit similar spectral characteristics, apart from the very high velocities seen in some SNe~IIn (e.g., SN 1988Z; Filippenko 1997).  SN 2000ch in NGC 3432 is clearly similar to these other subluminous objects and thus probably related to the LBV class.  Its spectrum is similar to that of SN 1997bs (Van Dyk et al. 2000), although SN 2000ch lacked the strong \\ion{Fe}{2} emission observed in SN 1997bs.  The Balmer line widths are similar in the two objects, implying comparable outflow velocities.  In addition, the maximum absolute magnitude of $M_V \\approx -13.8$ mag for SN 1997bs compares reasonably well with the value of $M_V = -12.7$ mag for SN 2000ch. While their basic spectral properties and maximum luminosities are similar, their host environments are quite different. SN 2000ch is located in an \\ion{H}{2} region or large star-forming complex at the edge of NGC 3432 (see Figure 2) while SN 1997bs does not appear to reside in a comparable environment (Van Dyk et al. 2000).  The most striking difference is in their light curves.  SN 2000ch was a very fast, erratic event characterized by a rapid rise of $\\sim2$ mag from quiescence, a local minimum lasting $\\sim$15~d, and finally a return back to quiescence. In contrast, the light curve of SN 1997bs reveals a gradual decline of $\\sim$6 mag over $\\sim$260~d. Smith, Humphreys, \\& Gehrz (2001) point out that the fast photometric variations exhibited by SN 2000ch are reminiscent of the oscillations in brightness reported for SN 1954J and $\\eta$ Carinae prior to and during their giant eruptions.  In summary, SN 2000ch appears to be related to the LBV phenomenom but with rapid photometric variations near maximum brightness.  While SN 1997bs has faded below detectability in NGC 3627, SN 2000ch is quite luminous in quiescence ($R \\approx 19.4$ mag) and can be studied in detail using modern instruments on large telescopes.  It is brighter than the progenitor of SN 1954J (Smith et al. 2001) and thus can be photometrically monitored frequently with small telescopes or observed spectroscopically at relatively high spectral resolution.  We encourage further monitoring and detailed observations of this unusual outburst in an effort to discern its nature and similarity to the class of LBVs."
    },
    "0404/astro-ph0404217_arXiv.txt": {
        "abstract": "{Quasi-hydrostatic cooling of the intracluster gas is studied.  In the quasi-hydrostatic model, work done by gravity on the inflow gas with $dP \\neq 0$, where $P$ is the gas pressure, is taken into account in the thermal balance.  The gas flows in from the outer part so as to compensate the pressure loss of the gas undergoing radiative cooling, but the mass flow is so moderate and smooth that the gas is considered to be quasi-hydrostatic. The temperature of the cooling gas decreases toward the cluster center, but, unlike cooling flows with $dP = 0$, approaches a constant temperature of $\\sim 1/3$ the temperature of the non-cooling ambient gas.  This does not mean that gravitational work cancels out radiative cooling, but means that the temperature of the cooling gas appears to approach a constant value toward the cluster center if the gas maintains the quasi-hydrostatic balance. We discuss the mass flow in quasi-hydrostatic cooling, and compare it with the standard isobaric cooling flow model. We also discuss the implication of $\\dot{M}$ for the standard cooling flow model. ",
        "introduction": "The intracluster gas, i.e. hot gas in clusters of galaxies, which undergoes radiative cooling is thought to flow in toward the cluster center to maintain pressure equilibrium (see Fabian \\cite{Fabian94} for a review and references therein).  According to the standard cooling flow picture, the mass deposition rate $\\dot{M}$ can be estimated from the differential luminosity.  X-ray observations showed that $\\dot{M}$ inferred from the surface brightness profile was over 100~$M_{\\sun}$~yr$^{-1}$ for so called cooling flow clusters.  However, firm evidence has not been found for such a large amount of cooled gas that should be detected at longer wavelengths if it existed in the cluster core (e.g., Edge \\cite{Edge01}).  {\\it ASCA} and recent {\\it Chandra} and {\\it XMM-Newton} observations (e.g., David et al. \\cite{David01}; Kaastra et al. \\cite{Kaastra01}; Allen et al. \\cite{Allen01}; Peterson et al. \\cite{Peterson03}) with better spatial and spectral resolution reveal, however, that $\\dot{M}$ is smaller by an order of magnitude than what was estimated before from the surface brightness. Moreover, spectroscopic analysis shows that the gas temperature decreases toward the cluster center, as expected for cooling flows, but down to only half or one-third the temperature of the non-cooling ambient gas.  This implies the lack of soft X-ray emission, or no appreciable contribution of lower temperature  gas.  To explain the observations, various theoretical models have been proposed (see e.g., B\\\"{o}hringer et al. \\cite{Boehringer02}; Peterson et al. \\cite{Peterson03}).  Many of those consider the central active galaxy (with an active galactic nucleus (AGN)) as a heating source to compensate radiation loss (e.g., Churazov et al. \\cite{Churazov01}; Ruszkowski \\& Begelman \\cite{Ruszkowski02}; Kaiser \\& Binney \\cite{Kaiser03}).  This idea is interesting from the aspect of history of activity (Reynolds, Heinz \\& Begelman \\cite{Reynolds02}), since many cooling flow clusters possess cD galaxies in their central regions.  However, some fine tuning is required. While the radiative cooling rate is determined by the local values of temperature and density, the conductive or convective heating rate including hot bubbles from AGN depends on the scale length of the temperature or pressure as well. This could make the processes complicated, and the parameters such as diffusion coefficient or mixing length remain open (Narayan \\& Medvedev \\cite{Narayan01}; Voigt et al. \\cite{Voigt02}).  In a standard analysis of X-ray data, an isobaric cooling flow model is applied to estimate $\\dot{M}$ (Fabian \\cite{Fabian94}). However, analysis based on surface brightness and analysis based on spectrum give inconsistent results, as mentioned above, in spite of both being based on the same physical concept. This may suggest some problem underlying the standard cooling flow model. In the standard isobaric model, work done by gravity on the inflow gas is not properly taken into account.  This is not an additional heating source but is naturally expected for the gas under gravity.  If the inflow is induced and regulated smoothly by local radiative cooling under gravity, quasi-hydrostatic structure may be attained.  Such a simple picture is thought to be a starting point for considering an alternative heat source.  In the present paper we examine quasi-hydrostatic intracluster gas undergoing radiative cooling, taking proper account of the work done by gravity.  In Sect. 2 we present our quasi-hydrostatic cooling model. We compare our model with the standard isobaric cooling flow model in Sect. 3, and give concluding remarks in Sect. 4. ",
        "conclusions": "In the quasi-hydrostatic model, we consider the ideal condition that a change is quasi-static.  The pressure loss due to radiative cooling is immediately compensated by continuous, smooth, spherical mass inflow from the outer part.  Neither turbulence nor overshooting occurs during the process.  The inflow is so moderate as not to disturb the hydrostatic balance significantly.  Thus, if the local inflow rate is perfectly controlled by the local cooling rate, the gas is quasi-hydrostatic and $\\dot{M}_r \\sim \\tilde{M}_r/\\tilde{t}_{cool}$ is attained; we will investigate this point further by means of hydrodynamical calculations elsewhere.  As mentioned in Sect. 2, `quasi-hydrostatic' means not static but steady-state.  Accordingly, although the temperature of the cooling gas approaches a constant value toward the cluster center (Fig. 1), this does not mean that gravitational work can cancel out radiative cooling.  This temperature profile is no more than a result of the steady inflow to maintain the hydrostatic balance. It should be noted that the gravitational potential or $M^*_r$ is assumed to be little affected by the gas inflow.  At the cluster center (or close to a cD galaxy) where the inflow mass is accumulating, the cooling time can be too short for the actual flow rate to follow the cooling rate of the gas inside, i.e. $\\tilde{M}_r/\\tilde{t}_{cool} > \\dot{M}_r$. Then the quasi-hydrostatic balance would break, and the gas temperature inside decreases below $\\sim (1/3)\\beta T_0$, yet the temperature in the steady flow region will continue to follow Eq. (\\ref{dlnT/dlnr}).  In the opposite case, if the cooling rate at a given radius is smaller than the inflow rate as $\\tilde{M}_r/\\tilde{t}_{cool} < \\dot{M}_r$, the flow rate $\\dot{M}_r$ would decrease inward and the temperature approaches a constant value higher than $\\sim (1/3)\\beta T_0$ toward the cluster center.  This may be the case when cooling in the core is not yet very significant or when some heating works against local radiative cooling so that $\\tilde{t}_{cool}$ is effectively longer than given by Eq. (\\ref{t_cool}).  The gas gains momentum during inflow.  In the quasi-static model this is included into Eq. (\\ref{dP/rho}), which determines the hydrostatic structure in cooperation with EOS.  This is valid as long as the momentum flux is small compared to the thermal pressure.  In practice, however, particularly in the region where the temperature approaches a constant value (see Fig. 1), the inflow gas may partly overshoot because of a fluctuation about $\\dot{M}_r \\sim \\tilde{M}_r/\\tilde{t}_{cool}$.  This may lead to instability and result in a break in the spherically symmetric structure.  The resultant dense part cools more rapidly, but the gas remains quasi-hydrostatic on the whole if the momentum flux is still small on average.  Within the context of quasi-hydrostatic cooling, the gas inflow must be mild at every radius. However, if the flow evolved so that the momentum flux would become considerable, though still subsonic, the quasi-hydrostatic balance would fail unless some force works against the inertia.  In such a case, for the gas to be quasi-hydrostatic, some momentum flux not in the form of the thermal pressure by heat may be needed against the inflow.  Observations by coming missions can examine the radial profiles predicted by the quasi-hydrostatic cooling model for several quantities, such as local cooling time $\\tilde{t}_{cool}$, mass flow rate $\\dot{M}_r$ or local luminosity $\\tilde{L}_r$, and possibly inflow velocity as well as the density and temperature. The observed temperature profiles can be understood invoking quasi-hydrostatic cooling, though  some momentum/heat source may be working additionally.  The simple quasi-hydrostatic model is expected to give a basis for understanding the thermal properties of the intracluster gas under gravity and for advanced study using hydrodynamics codes."
    },
    "0404/astro-ph0404484_arXiv.txt": {
        "abstract": "To determine the nature of dark energy from observational data, it is important that we use model-independent and optimal methods. We should probe dark energy using its density (allowed to be a free function of cosmic time) instead of its equation of state. We should minimize gravitational lensing effect on supernovae by flux-averaging. We need to include complementary data (for example, from the Cosmic Microwave Background [CMB] and large scale structure [LSS]) in a consistent manner to help break the degeneracy between the dark energy density and the matter density fraction. We should push for ambitious future supernova surveys that can observe a large number of supernovae at the highest possible redshifts. I discuss these and other issues that will be important in our quest to unravel the mystery of the nature of dark energy.  Current supernova, CMB, and LSS data already rule out dark energy models with dark energy densities that vary greatly with time; with the cosmological constant model providing an excellent fit to the data. A precise measurement of dark energy density as a free function of cosmic time will have a fundamental impact on particle physics and cosmology. ",
        "introduction": "\\label{sec.intro}  The data of type Ia supernovae (SNe Ia) reveal that the expansion of our universe is accelerating \\cite{Riess98,Perl99}. This implies the existence of dark energy. Solving the mystery of the nature of dark energy will have revolutionary implications in particle physics and cosmology.  The first most fundamental question about dark energy is whether the dark energy density depends on time. The time dependence of the dark energy density can be used to distinguish the various models. \\cite{Freese}  A powerful probe of dark energy is SNe Ia, which can be calibrated to be good cosmological standard candles \\cite{Phillips,Riess95,Branch98}, and used to measure how distance changes with redshift. For a flat universe, the luminosity distance is \\ba \\label{eq:dL(z)} d_L(z) &=& cH_0^{-1} (1+z) \\int_0^{z} \\frac{dz'}{E(z')},\\nonumber\\\\ E(z) &= &\\left[ \\Omega_m (1+z)^3 + (1-\\Omega_m) \\rho_X(z)/\\rho_X(0)\\right]^{1/2}, \\ea where $\\rho_X(z)$ is the dark energy density. Note that dark nergy only appears in the form of the function $E(z)$; $\\rho_X(z)$ is on the same footing as the matter density fraction $\\Omega_m$ in the context of observational tests. The dark energy equation of state is \\cite{Wang01a} \\be \\label{eq:wXz} w_X(z)= \\frac{p_X(z)}{\\rho_X(z)}= \\frac{(1+z)}{3} \\frac{\\rho_X'(z)} {\\rho_X(z)}-1. \\ee ",
        "conclusions": ""
    },
    "0404/astro-ph0404167_arXiv.txt": {
        "abstract": "We present three new binary pulsars discovered during a search for pulsations in 56 unidentified mid-latitude EGRET $\\gamma$-ray error boxes with the Parkes multibeam receiver. Timing observations of these sources is on-going with both the Parkes and the Green Bank telescopes. We discuss the place of these new systems in the population of binary pulsars and suggest that they are all somewhat atypical systems. ",
        "introduction": "We have recently completed a search of 56 unidentified mid-latitude EGRET $\\gamma$-ray sources for pulsations using the Parkes telescope and the multibeam receiver. This survey covered $\\sim$140 square degrees of sky at Galactic latitudes $> 5^{\\circ}$ and revealed six new pulsars, three of which are in binary systems. Here we present initial timing results for the three binaries. ",
        "conclusions": ""
    },
    "0404/astro-ph0404398_arXiv.txt": {
        "abstract": "{The pulsational properties of the Cepheid models along the evolutionary tracks from the Padova group (Girardi \\etal), as calculated with our turbulent convective pulsation code, are in good agreement with the resonance constraints imposed by the observational OGLE-2 data of the Small and Large Magellanic Clouds.  Our study suggests that the $P_4/P_1=1/2$ resonance for the overtone Cepheids occurs for periods clustering around 4.2\\th d, in disagreement with the suggestion of Antonello \\& Poretti based on the observations of light curves, but in agreement with Kienzle \\etal\\ and Feuchtinger \\etal.  For the fundamental Cepheids the lowest order Fourier decomposition coefficients from the light curves, \\viz $R_{21}$ and $\\phi_{21}$ can be used to locate the resonance region, but not so for the first overtone Cepheids.  Here, the radial velocity curves can be used to locate the overtone resonance region, or in their absence, one needs to resort to numerical hydrodynamic modelling. ",
        "introduction": "Cepheids are perhaps the best observed variable stars.  They are very interesting heat engines from a physical point of view, and it is hardly necessary to stress their importance as cosmological distance indicators. Obviously, a good theoretical understanding of these important stars is essential, and they have received the most theoretical attention among the variable stars, going as far back as a half-century.  It was first noted by Hertzsprung (1926) that the shape of the light curves (LC) of the Galactic fundamental (F) mode Cepheids exhibit a progression with pulsation period in the vicinity of $10$\\th days.  Subsequently Payne-Gaposhkin (1947) found concomitant sharp variations in the Fourier decomposition coefficients of the LCs.  Later, Simon \\& Schmidt (1976) noted a correlation of the location of the bump on the LC with the location of a resonance between the excited fundamental pulsation mode and the second overtone ($P_2/P_0 = 1/2$).  Intrigued by these features Buchler \\& Goupil (1984) developed the mathematical 'amplitude equation' formalism that is necessary to understand how the occurrence of an internal resonance can produce, both qualitatively and quantitatively, a progression of the Fourier coefficients of the LCs and of the radial velocity ($V_{r}$) curves (see also Buchler 1993 and Buchler \\& Kov\\'acs, 1986).  The amplitude equations for a given star yield the modal amplitudes as a function of the stellar parameters, such as the period or \\Teff\\ for example, and these in turn are related to the Fourier coefficients of the LC.  Klapp, Buchler \\& Goupil (1985) and Kov\\'acs \\& Buchler (1989) applied this formalism to the specific case of the F Cepheid progression.  A comparison of a number of sequences of full amplitude hydrodynamic Cepheid models in Figs.~3, 5 and 7 of Buchler, Moskalik \\& Kov\\'acs (1990) shows very clearly that the structure of the (magnitude) Fourier coefficients $\\phi^m_{21}= \\phi^m_2-2\\phi^m_1$ and $R^m_{12}= A_2/A_1$ correlates very well with the period ratio $P_r=P_2/P_0$, which is indicative of the resonance, rather that with the period $P_0$ itself.  (We note that it is because of the finite width of the IS that stars with different masses and luminosities, and thus different periods can have the same $P_r$ as Fig.~\\ref{ber} below will show.  A corollary is that there is no real Cepheid resonance center, but rather a resonance region).  For the observed Cepheids we know the periods, but we do not have any direct information on the period ratios $P_{r}$.  We therefore try to take advantage of the behavior of the Fourier coefficients to localize the resonance region.   The Fourier coefficients $\\phi^{\\v}_{21}$ and $R^{\\v}_{21}$ of the observed $V_{r}$ curves of Galactic F Cepheids also show structure in the vicinity of a 10\\th d period (\\eg Kov\\'acs, Kisvarsanyi \\& Buchler 1990), but the structure is different from that of the LCs.  This is to be expected because the luminosity $L\\sim R^2 T^4$ involves not just the radial displacement eigenvectors but also the temperature eigenvectors.  The coefficients in the transformation from the modal amplitudes to LCs and to $V_{r}$ curves are therefore different as explained in Kov\\'acs \\& Buchler (1989).  Thus, for example, in the F Cepheids the sharp feature in the magnitude phases $\\phi^{m}_{21}$ occurs close to 10\\th d, and $R^{m}_{21}$ has a minimum there \\th~(\\cf \\eg Fig.~6 of Moskalik, Buchler \\& Marom 1992).  The corresponding $V_{r}$ phase $\\phi^{\\v}_{21}$ has a rapid, but not sharp variation from 5 to 15\\th d, while the minimum for $R^{\\v}_{21}$ is less pronounced and appears at a somewhat larger period of $\\sim$ 12\\th d.  Either one of the two coefficients, $\\phi^{m}_{21}$ or $R^{m}_{21}$, can therefore serve as a criterion for locating the resonance, but $\\phi^{\\v}_{21}$ and $R^{\\v}_{21}$ are less useful discriminants.     \\begin{figure*} \\centerline{\\resizebox{8.9cm}{!}{\\includegraphics{plot_bertelli_SF.ps}} \\hskip 4pt \\resizebox{8.9cm}{!}{\\includegraphics{plot_bertelli_LF.ps}}} \\vskip 3pt \\centerline{\\resizebox{8.9cm}{!}{\\includegraphics{plot_bertelli_SO.ps}} \\hskip 4pt \\resizebox{8.9cm}{!}{\\includegraphics{plot_bertelli_LO.ps}}} \\parbox[b]{18cm}{ \\caption[]{\\small {\\bf F and O1 Cepheid models in SMC and LMC)}: Sections of the 3.5, 4.0, 4.5, 5.0, 6.0 and 7.0 \\Mo\\ Padova ($Z$=0.004 and 0.008) tracks corresponding to the third crossing of the IS; the periods [d] are noted above the tracks, the period ratios $P_r$ below; the {\\sl thickened sections} of the tracks represent the linear IS (linearly unstable models);  the pairs of thick {\\sl vertical curves} delimit the resonance regions, $P_r = 0.5 \\pm 0.015$.} \\label{ber} } \\end{figure*}    We now turn to the O1 Cepheids.  Antonello \\& Poretti (1986) suggested evidence for an internal resonance of the excited first overtone with the fourth overtone $P_{r}=P_4/P_1\\sim 1/2$ in the vicinity of $P_1\\sim 3.2$\\th d, on the basis of the sharp variations of $\\phi^{m}_{21}$ of the Galactic LC Fourier coefficients.  Linear models of O1 Cepheids indeed display such a resonance in that general vicinity (Antonello 1994, Buchler \\etal, 1996), the exact location depending however on the $M$ -- $L$ relation that is used.  The first results of nonlinear radiative model did not correctly reproduce the variation of the Fourier coefficients (Antonello \\& Aikawa 1995).  In contrast, on the basis of the $V_{r}$ Fourier coefficients Kienzle, Moskalik, Bersier \\& Pont (1999), instead, put the resonance near 4.6\\th d.  The reason for this discrepant conclusion is that the Fourier coefficients, especially $\\phi^m_{21}$ have a substantially different behavior for the LCs and the $V_{r}$ curves.  See \\eg Feuchtinger, Buchler \\& Koll\\'ath (2000, hereafter FBK) for a juxtaposition.  In order to resolve this discrepancy it was necessary to make a detailed numerical hydrodynamical survey of full amplitude pulsations (FBK).  The code contains a time-dependent mixing length model for convection because purely radiative models had failed to reproduce the observed features of the Fourier coefficients, in particular near 3.2\\th d.  This survey clearly put the resonance  near 4.2\\th d.  But it also showed that only the broad maximum of $\\phi^{\\v}_{21}$ correlates with the resonance.  The conclusion is that while it is true that sharp features in the Fourier coefficients generally indicate the presence of internal resonances, there is {\\sl no simple} \\apriori criterion for locating precisely the resonance region from either the LC or from the $V_{r}$ data.  Numerical hydrodynamic modelling of full amplitude pulsations is needed to determine the criterion to be used and the location of the resonance.  The Magellanic Cloud F Cepheid LCs appear very similar to those of the Galactic ones in terms of their Fourier coefficients (\\eg Beaulieu et al., 1995, Beaulieu \\& Sasselov 1997).  What matters in particular for this paper is that the resonant features are very similar.  The observed LC features of the LMC and SMC F Cepheids occur in the same place as for the Galactic ones.  A systematic and comparative hydrodynamic full amplitude survey of Galactic and MC F Cepheids is still lacking, and in its absence we make the reasonable assumption that the Fourier structure as a function of $P_r$ is the same.   There are some small differences though for the overtones.  Despite a fair amount of scatter, it would appear that the prominent sharp features, namely the minimum in $R^{m}_{12}$ occurs near 2.7 and 2.4\\th d for the LMC and SMC O1 Cepheids respectively (Beaulieu et al., 1995, Beaulieu \\& Sasselov 1997, and the larger OGLE sample by Udalski \\etal 1999a, 1999b), whereas they occur closer to 3.2\\th d for the Galactic O1 Cepheids Antonello \\& Poretti (1986). There is a similar small shift in the sharp drop in $\\phi^m_{21}$ to lower periods.  In the absence of a numerical hydrodynamic survey similar to that of FBK, we will therefore assume on the basis of these shifts in $R^m_{21}$ that there are probably corresponding shifts of --0.5 and --0.7\\th d of the resonance period for the LMC and SMC O1 Cepheids.  In a previous attempt to use the resonances to obtain constraints on the $M$ and $L$ of the Cepheid models Buchler, Koll\\'ath, Beaulieu \\& Goupil (1996) (see also Simon \\& Kanbur 1994, Aikawa \\& Antonello 2000) made the assumption that the resonance boundaries, for example for F Cepheids, were defined for the resonance center, \\viz $P_2/P_0$= 1/2, at $P_0$ = 11\\th d at the blue edge, and 9\\th d at the red edge of the IS.  This led to their suggestion that the derived masses for the SMC and LMC were too small.  The current paper reexamines this assumption.  Fig.\\ref{ber} and the discussion below show that instead one should define the resonance boundaries to be $P_r=P_2/P_0=0.5+\\Delta$ at the blue edge and $P_r=0.5-\\Delta$ at the red edge, with $\\Delta = 0.015$ for example, both for $P_0$ = 10\\th d and $P_1=4.2$\\th d for O1 Cepheids).  If one repeats their procedure with these new conditions one finds that the sensitivity to the exact values of the resonance boundaries ($P_r$) is too great to constrain $M$ and $L$.  Hence we are led here to an alternate comparison, still involving the resonances, that is based on a combination of stellar evolution tracks and pulsation properties.  Other studies (Baraffe et al., 1998, Alibert et al., 1999, Bono et al., 2000a, Bono et al., 2000b, Bono et al., 2001) usually concentrate on the comparison of the models with the observations in the period -- luminosity planes in different bands and ignore the strong constraints on dynamics coming from the resonances.  The main purpose of this paper is to check how well the models satisfy these resonance constraints when confronted with the observational OGLE-2 Magellanic Cloud (MC) data.    \\subsection{Models on the Tracks of Girardi \\etal}   In Fig.~\\ref{ber} we display sections of the evolutionary tracks of \\cite{Girardi} for $Z=0.004$ and for $Z=0.008$ corresponding to the third crossing of the instability strip (IS).  They are meant to be approximately representative of the average properties of the SMC and LMC Cepheids.  We do not show the second crossings of the IS here because they will give similar results.  We note the well known fact that the blueward extents of the low mass evolutionary tracks, \\eg for 3.5\\Mo\\, are too short, and for the LMC the 3.0\\Mo\\ track does not even penetrate the IS.  This problem is not specific to the Padova tracks, but is encountered by all recent evolution calculations. See Cordier, Goupil \\& Lebreton (2003) for a recent discussion of this problem with evolutionary tracks and a possible explanation.  However the problem occurs outside both the F and the O1 Cepheid resonance regions that we are concerned with here.  As far as pulsation modelling is concerned the Padova evolutionary models provide us with a mass $M$, a luminosity $L$, and effective temperature \\Teff. For a given composition (specified by $X$ and $Z$) these three quantities {\\sl uniquely specify the stellar envelope} \\th \\ie without a need to know the properties of the stellar interior.  (One could equivalently, but much less conveniently characterize the Cepheid envelope instead by the luminosity, temperature and pressure at the core radius, for example).  We note that the envelope composition is uniform having recently undergone a fully convective stage.  The problems associated with core convection and overshooting (\\eg Cordier \\etal\\ 2003) thus need not concern us here because pulsation is limited to the part of the star (the envelope) that is located above the burning shells (\\ie $T<$ a few million K).       \\begin{figure*} \\centerline{\\resizebox{16cm}{!}{\\includegraphics{mc_evol_osc_phiFB.ps}}} \\vskip -1pt \\parbox[b]{17cm}{ \\caption[]{\\small {\\bf MC F Cepheids -- $\\phi^m_{21}$ resonance criterion:} {\\sl dots (open circles)} represent linearly stable (unstable) OGLE models; {\\sl squares}: $\\phi^m_{21} > 5.45$ (and $P_0<20$), {\\sl tilted squares}: $0 < \\phi^m_{21} < 2$. The resonant Cepheid models lie between the vertical curves and the envelopes of the thickened tracks (instability strip). } \\label{phiFB} } \\end{figure*}   \\begin{figure*} \\centerline{\\resizebox{16cm}{!}{\\includegraphics{mc_evol_osc_r21FB.ps}}} \\vskip -1pt \\parbox[b]{17cm}{ \\caption[]{\\small  {\\bf MC F Cepheids -- $R^m_{21}$ resonance criterion:} {\\sl dots (open circles)} represent linearly stable (unstable) OGLE models; {\\sl squares}: $R^m_{21} < 0.15$ \\th (and $P_0 < 16$), {\\sl tilted squares}: $0.15 < R^m_{21} < 0.2$ \\th (and $P_0<16$), {\\sl triangles}:  $0.2 < R^m_{21} < 0.25$ \\th (and $6<P_0<16$). The resonant Cepheid models lie between the vertical curves and the envelopes of the thickened tracks (instability strip).}  \\label{r21FB} } \\end{figure*}    The equilibrium Cepheid models and their {\\sl linear} vibrational stability properties are computed with our turbulent convective, TC code (\\eg Yecko, Koll\\'ath \\& Buchler 1998, Koll\\'ath, Buchler, Szab\\'o \\& Csubry 2002) for the model sequence along the Padova tracks.  One might think that there is a remaining inconsistency because Girardi \\etal~(2000) use standard (time-independent) mixing length in the envelope, whereas we have additional ($\\alpha$) parameters associated with our time-dependent mixing length model, which also contains a turbulent flux and a turbulent pressure.  The structure of the envelope turns out to be relatively insensitive to these differences, as are the pulsation periods of the lowest modes.  In contrast, as one may expect, the modal stability is more sensitive to the $\\alpha$'s.  We can therefore safely adopt the $M$, $L$ and \\Teff\\ given by Girardi \\etal\\ and expect our envelopes to be very close to their evolutionary ones.  The top plots in Fig.~\\ref{ber} show the F Cepheid models and the bottom plots the O1 Cepheid models.  The numbers above the tracks represent the linear pulsation periods, $P_0$ and $P_1$ [d], respectively. \\th (We recall here that the nonlinear and linear pulsation periods are known to differ only by a few tenths of a percent.)  The numbers below the tracks refer to the (linear) period ratios $P_{r}=P_2/P_0$ and $P_{r}=P_4/P_1$, respectively, in days.  The pairs of thick vertical curves delineate the boundaries of the resonance regions which we define as $0.485<P_{r}<0.515$.  The thickened portions of the tracks in Fig.~\\ref{ber} represent the linear IS, \\ie the region where the F or O1 Cepheid model, resp., are linearly unstable. The actual IS is a little narrower because nonlinear effects are known to shift the F blue edge a little lower in temperature, 0 -- 200\\th K, depending on $M$, and the O1 red edge a few hundred K to higher temperatures (\\eg Fig.~2 of Buchler 2000).  In addition, it is well known that there is a region where either F or O1 pulsation is possible depending on the direction of the evolutionary track (\\eg Koll\\'ath \\etal\\ 2002).    \\begin{figure*} \\centerline{\\resizebox{16cm}{!}{\\includegraphics{mc_evol_osc_r21OB.ps}}} \\vskip -1pt \\parbox[b]{17cm}{ \\caption[]{\\small  {\\bf MC O1 Cepheids -- $R^m_{21}$ resonance criterion:} {\\sl squares}: $R^m_{21} < 0.05$ \\th (and $P_1<4.5$), {\\sl tilted squares}: $0.05 < R^m_{21} < 0.07$ \\th   (and $P_1<4.5$). The resonant Cepheid models lie between the vertical curves and the envelopes of the thickened tracks (instability strip). The apparent lack of agreement is discussed in the text. } \\label{r21OB} } \\end{figure*}   \\begin{figure*} \\centerline{\\resizebox{16cm}{!}{\\includegraphics{mc_evol_osc_per1B.ps}}} \\vskip -1pt \\parbox[b]{17cm}{ \\caption[]{\\small {\\bf SMC O1 Cepheids -- period  'resonance' criterion:} {\\sl dots (open circles)} represent linearly stable (unstable) OGLE models; {\\sl circles}:        $4.45 < P_1 < 4.95$, {\\sl tilted squares}: $3.95 < P_1 < 4.45$, {\\sl squares}:        $3.45 < P_1 < 3.95$, {\\sl hexagons}:       $2.95 < P_1 < 3.45$, {\\sl triangles}:      $2.45 < P_1 < 2.95$. The resonant Cepheid models lie between the vertical curves and the envelopes of the thickened tracks (instability strip). } \\label{perOB} } \\end{figure*}     In Fig.~\\ref{ber} the resonant F Cepheids therefore should fall in the {\\sl resonance regions} defined by the resonance boundaries, $P_2/P_0 = 0.485$ and $P_2/P_0 = 0.515$ (vertical curves) and the boundaries of the IS. \\th (The pulsationally unstable models marked by the thickened sections of the tracks). These are the approximate values for which the (cos) Fourier coefficients $R^m_{21}$ and $\\phi^m_{21}$ have an egregious behavior, \\viz the $R^m_{21}$ have a dip, and the phases $\\phi^m_{21} \\approxgt 5.5$ or $\\phi^m_{21} \\approxlt 2.0$.  We see that for the resonant Cepheids, independently of $Z$ and of the pulsation mode, the \"NW\" corner on the {\\sl blue edge} of the IS, with the highest $L$ and $M$ has the highest $P_r$.  For example for the SMC F Cepheids one finds $M\\sim 7.3$\\Mo, with a period ($P_0\\sim 16.5$d).  Conversely, the \"SE\" corner of the {\\sl red edge} of the resonant IS, with the lowest $L$ and $M$ has the lowest $P_r$.  For the SMC F Cepheids $M\\sim 4$\\Mo with $P_0\\sim 5$.  The resonant F Cepheid models along the $Z=0.004$ Padova tracks therefore have periods ranging from $\\sim 5$ to 16.5\\th d.  A similar topography is seen to obtain for the LMC Cepheids.  From the periods that are indicated in Fig.~\\ref{ber} one can readily infer that constant period curves run NW to SE, albeit with a slope that is shallower than that of the constant $P_{r}$ curves.  If instead the resonance regions were defined by the edges of the IS and the edges of the constant period curves, they would be reasonably close to the resonance regions that we have just defined on the basis of $P_r$, as one would expect from the correlation of $R^m_{21}$ with the period in the OGLE data, for example.  A similar situation occurs for the overtone Cepheids where the resonant O1 Cepheid models along the $Z=0.004$ Padova tracks have periods ranging from $\\sim 2.0$ to 6.5\\th d.  The constant $P_{r}$ curves that demarcate the resonant O1 models show a pinch near 6\\Mo.  Interestingly, this is not a numerical artifact, but can be traced to the nonmonotone behavior of the the period $P_4$ of the fourth overtone, because of the occurrence of a strange mode (Buchler \\& Koll\\'ath 2001, \\eg Fig.~2; Buchler, Yecko \\& Koll\\'ath, 1997).  For the $Z=0.008$ Padova tracks (third crossing of the IS) that are approximately representative of the LMC Cepheids, we obtain a very similar overall picture.  In particular, we find approximately the same period ranges for the resonant F Cepheid models ($P_0 \\sim$ 5.5 to 16\\th d) and for the resonant O1 Cepheid models ($P_1 \\sim$ 2.2 to 7.2\\th d).  Note that a similar pinch in the resonance curves also occurs for the $Z=0.008$ tracks.   \\subsection{Comparison with the OGLE-2 Cepheids}  \\subsubsection{Instability Strip: Model Stability}  Beaulieu, Buchler \\& Koll\\'ath (2001, hereafter BBK) analyzed the SMC and LMC OGLE Cepheids (OGLE-2 web database) with the intent of extracting $M$ -- $L$ relations from the observational data.  They first converted the observed magnitudes and colors into $L$ and \\Teff\\ with assumptions about reddening and distance moduli.  We shall refer to these observationally derived $L$, \\Teff~ and periods as defining the {\\sl OGLE stars}.  From these OGLE star data BBK then derived stellar masses $M$ as well as linear stability properties with the help of a linear pulsation code.  In the following we shall refer to these as the {\\sl OGLE models} when there is a need to distinguish between them.  We recall that BBK had concluded that observational luminosity and reddening uncertainties cause a small spread in \\Teff\\ that is difficult to disentangle from the spread in \\Teff\\ due to the width of the IS.  The derived $M$, $L$ and \\Teff\\ values for the OGLE models thus contain small errors.  The two consequences are that (1) in an $M$ -- $L$ diagram one obtains a swarm rather than a narrow strip as one would expect from a relatively homogeneous group of Cepheids, and that (2) in an \\Teff\\ -- $L$ plot some stars necessarily fall outside the actual IS.  Figs.~\\ref{phiFB} and \\ref{r21FB} show theoretical HR diagrams (Log $L$ \\vs Log \\Teff) in which we redisplay the OGLE F Cepheid stars for the BBK's (preferred) choice B of distance modulus, namely 18.55$\\pm$0.10 for the LMC and 19.97$\\pm$0.15 for the SMC, and mean reddenings of E(B-V) = 0.1 and 0.08, respectively.  These figures are identical except for the resonance criteria to be addressed below.  In these figures we now represent the linearly stable/unstable OGLE models with open circles/dots.  Despite the above mentioned observational uncertainties in the derived $M$, $L$ and \\Teff\\ we see that the majority of the OGLE Cepheid models are unstable, as they should be. There appears to be some small systematic discrepancy, however, for the low luminosity F Cepheids on the red side, both in the LMC and SMC that goes beyond the uncertainties in the model parameters.  Considering that the linear growth rates are the least certain of the calculated pulsation quantities, because they depend on the $\\alpha$ parameters that are used in the time-dependent mixing length equations we are not too concerned because a small adjustment of the $\\alpha$ parameters in the convective terms could fix this problem.  We stress again that, in contrast, the periods are largely independent of the $\\alpha$\\th s.  In Figs.~\\ref{r21OB} and \\ref{perOB} we reproduce the OGLE O1 Cepheids.  One notes that the vast majority of the OGLE models are linearly unstable, and that the region occupied by the unstable evolutionary models coincides well with that of the unstable OGLE models.  Overall, the figures thus indicate good agreement between the tracks, the pulsation calculations and the observations as far as the stability of the models is concerned.  \\subsubsection{Resonance Conditions}  We now go on to examine how well the resonance information along the evolutionary tracks agrees with that of the Fourier coefficients of the OGLE LCs.  We start with the F Cepheids, both in SMC and LMC.  \\vskip 5pt  \\underbar{\\sl Fundamental Mode Cepheids:}  \\vskip 5pt  As we have already discussed in the Introduction, for the observed F Cepheids we know the periods, but we do not have any direct information on the {\\sl period ratios} $P_{r}$.  However, near $P_r\\sim 0.50$ the period ratio correlates with the dip of the LC's Fourier amplitude ratio $R^m_{21}$ .  It also correlates well with the egregious values of the phases $\\phi^m_{21}$ in the resonance region, \\viz $\\approxlt 2.0$ and $\\phi^m_{21} \\approxgt 5.3$. The phases are normalized to [0, 2$\\pi$ ], mod 2$\\pi$.  In Figs.~\\ref{phiFB} and \\ref{r21FB} we display \\Teff\\ -- $L$ plots for the F Cepheids.  The OGLE F Cepheids are represented by dots for the linearly unstable models, and by open circles for the stable ones.  Superposed are the Padova tracks where again the thickened sections represent the linearly unstable models (as in Fig.~\\ref{ber}), and the vertical curves denote the edges of the resonance regions, defined by $0.485 < P_r < 0.515$.  In Fig.~\\ref{phiFB} we use the Fourier phase ($\\phi^m_{21}$) resonance criterion.  We have surrounded by {\\sl squares} the stars for which $\\phi^m_{21} > 5.45$ (and $P_0<20$), and by {\\sl tilted squares} those for which $\\phi^m_{21} < 2$.  The position of these resonant OGLE stars is seen to agree rather well with the resonance region that we have defined with the help of the models along the Padova tracks, despite the large scatter in the OGLE Fourier coefficients that manifests itself in our figures.  In Fig.~\\ref{r21FB} we use instead the resonance criterion based on the $R^m_{21}$ coefficients.  We surround the stars for which $R^m_{21} < 0.15$ and $P_0<16$ by square boxes, those for which $0.15 < R^m_{21} < 0.20 $ and $P_0<16$ by tilted squares, and those for which $0.20 < R^m_{21} < 0.25 $ and $6<P_0<16$ by triangles.  We have chosen these numerical values on the basis of the plots of $R^m_{21}$ \\vs period of the OGLE Cepheids (Udalski \\etal 1999a, 1999b).  One notes that this second criterion gives results that are very similar to those derived with the $\\phi^m_{21}$ criterion.  Again there is excellent agreement between the F Cepheid models along the Padova tracks and the OGLE F Cepheids.   \\vskip 5pt  \\underbar{\\sl First Overtone Cepheids:}  \\vskip 5pt  We have already noted in the Introduction that for the O1 Cepheids neither the phase $\\phi^m_{21}$ nor the amplitude ratio $R^m_{21}$ are useful discriminants for the location of the resonance (FBK).  Fig.~\\ref{r21OB} shows that if we persist in using $R^m_{21}$ as a resonance criterion, where the squares are stars with $R^m_{21} < 0.05$ and $P_1<4.5$, and the tilted squares those with $0.05 < R^m_{21} < 0.07 $ and $P_1<4.5$, we see that these 'resonant' stars fall substantially below the resonance region defined by the evolutionary tracks and pulsation theory.  The reason, as we have already pointed out, is that the behavior of $R^m_{21}$ through the resonance is different (as a function of $P_r$) for the O1 Cepheids and for the F Cepheids.  This point was argued in FBK where it was shown that nonlinear Galactic overtone Cepheid models give excellent agreement with the observational Fourier decomposition data, but that the resonance occurs around 4.2\\th d, rather than 3.2\\th d. Kienzle \\etal~ (1999) had also suggested a resonance at 4.6\\th d on the basis of the Fourier decomposition of the $V_{r}$ curves of the Galactic overtone Cepheids.  Since neither the $\\phi^m_{21}$ nor the $R^m_{21}$ resonance criteria are very simple or useful in the case of O1 Cepheids, we turn to the period $P_1$ as an alternate resonance criterion in Fig.~\\ref{perOB}, even though it is less restrictive as pointed out in the Introduction.  The different period ranges are indicated with different symbols, \\viz circles: $4.45 < P_1 < 4.95$, tilted squares: $3.95 < P_1 < 4.45$, squares: $3.45 < P_1 < 3.95$, hexagons: $2.95 < P_1 < 3.45$, triangles: $2.45 < P_1 < 2.95$.  In the Introduction we suggested on the basis of the comparative location of the minima of $R^m_{21}$ that in going from the Galaxy to the LMC and SMC one might expect similar shifts of --0.5 and --0.7\\th d in the locations of the resonance regions.  The latter would thus be centered on $P_1\\sim$ 3.7\\th d for the LMC (in the region of the squares in Fig.~\\ref{perOB}) and on $\\sim$ 3.5\\th d for the SMC (in the region straddling the squares and the hexagons). Fig.~\\ref{perOB} is certainly compatible with these conclusions.    \\vskip 5pt  In this paper we have shown the properties of the OGLE Cepheids obtained with BBK's choice B of distance modulus and reddening.  For completeness we have repeated the same calculations with their choice A.  It is noteworthy that choice B, which BBK labelled as preferred on other grounds, also gives better agreement between the resonance properties of the observations and the models than choice A. ",
        "conclusions": "We have used the properties of the internal 2:1 resonance between fundamental and second overtone modes for the F Cepheids (originally known as the Hertzsprung progression), as well as the 2:1 resonance between the first and fourth overtone modes for the O1 Cepheids to compare the observational, stellar evolutionary and pulsational properties.  We find a very good agreement between the evolutionary tracks of Girardi \\etal\\ (2000), our turbulent convective pulsation code and the OGLE observational constraints.  The only major, and well known disagreement is the shortness of the blueward evolutionary tracks for low masses where they do not even penetrate into the instability strip, but this occurs outside the resonance regions.  Another, but smaller disagreement occurs in the precise location of the ISs and their widths.  It has its origin partially in the errors in the $L$ and \\Teff\\ which themselves come from reddening and magnitude uncertainties in the observational OGLE data.  From the modelling side uncertainties arise from the values of the $\\alpha$ parameters that enter the pulsation code through time-dependent mixing length theory.  In addition, nonlinear effects that we have not considered here shift the F blue edge to lower \\Teff\\ and the O1 blue edge to higher \\Teff.   In the past the sharp features in the Fourier coefficients, that are indicative of the presence of internal resonances, have been used to localize internal resonances.  This has worked reasonably well for the Galactic F Cepheids, but as we have discussed here, in general, there exists {\\sl no simple} \\apriori criterion for localizing precisely the resonance region from either the LC or from the $V_{r}$ data.  One needs to resort to full amplitude numerical hydrodynamic modelling to determine the criterion to be used to localize the resonance center.  We find that we get much better agreement between the resonant models and the OGLE data if we put the O1 Cepheid resonance center ($P_4/P_1$=1/2) at $\\sim$ 3.7\\th d for the LMC and at $\\sim$ 3.5\\th d for the SMC, rather than where the $\\phi_{21}$ has its sharp drop.  The same suggestion was reached already by FBK in their detailed study of Galactic Cepheids where this resonance appear $\\sim$4.2\\th d."
    },
    "0404/astro-ph0404351_arXiv.txt": {
        "abstract": "{We use a grid of 30 line-blanketed unified stellar photosphere and wind models for O-type stars; computed with the code {\\sc cmfgen} in order to evaluate its potential in the near-infrared spectral domain. The grid includes dwarfs, giants and supergiants. We analyse the equivalent width behaviour of the 20 strongest lines of hydrogen and helium in spectral windows that can be observed using ground-based instrumentation and compare the results with observations. Our main findings are that: {\\em i)} \\hei/\\heii\\ line ratios in the J, H and K bands correlate well with the optical ratio employed in spectral classification, and can therefore be used to determine the spectral type; {\\em ii)} in supergiant stars the transition from the stellar photosphere to the wind follows a shallower density gradient than the standard approach followed in our models, which can be mimicked by adopting a lower gravity in our prescription of the density stratification. {\\em iii)} the \\brg\\ line poses a number of peculiar problems which partly might be related to wind clumping, and {\\em iv)} the \\bra\\ line is an excellent mass-loss indicator. For the first and last item we provide quantitative calibrations. ",
        "introduction": "A large fraction of the galactic population of massive stars lies hidden behind tens of magnitudes of visual extinction. The reasons for this are that massive stars are so rare that their typical distances are a sizable fraction of the Galactic scale and that they are concentrated in the Galactic disk. Therefore, these stars suffer from obscuration by intervening molecular gas and dust clouds in the line-of-sight.  Moreover, due to the short lifetimes of high-mass stars, they are located in star forming environments. As they typically form in the densest parts of these giant molecular clouds, they spend a significant fraction of their life embedded in this natal environment, before either moving out or breaking out.  Over the last decade several tens of massive stars were discovered in the near-infrared spectral window, where extinction by intervening dust is strongly reduced compared to the optical and ultraviolet window (Hanson et al. 2002, Kaper et al. 2002, Kendall et al. 2003). Studying the physical properties of these stars from their near-infrared radiation alone is essential (for instance in relation to their formation mechanism), but not an easy task. The J, H, K, and L spectral windows contain relatively few lines, mostly of hydrogen and helium. These lines are difficult to model as, for O-type stars, most near-infrared lines are formed in the transition region from the stellar photosphere (where the optical absorption spectrum originates) to the super-sonic stellar wind (see e.g. Kudritzki \\& Puls 2000). This makes a treatment of the stellar wind an integral aspect of quantitative spectroscopic studies of O-type stars in the near-infrared, perhaps excluding late-type O\\,V stars of which the stellar winds are relatively weak ($\\mdot \\leq 10^{-7}$ \\msunyr). A fundamental problem is that so far we have only a poor knowledge of the way in which the density structure in the transition region and lower part of the wind (up to a few times the sonic velocity) behaves. Though the basic driving mechanism of stellar winds has been identified (e.g. Castor, Abbott, \\& Klein 1975, Abbott 1982, Pauldrach et al. 1986) a fully self-consistent numerical implementation of radiative line-driving is, at present, not feasible for models which have been constructed for the objective of atmospheric analyses. Moreover, the inevitable assumptions made in describing this theory are anticipated to have severe effects on the physics of the transition region. As an example, the neglect of line source-function gradients when using the Sobolev approximation might lead to erroneous values for the radiative acceleration just in this transition region, at least in the case of thin winds (cf. Owocki \\& Puls 1999).  One may identify two essentially complementary approaches for making progress in the development of near-infrared diagnostics that allow for a characterisation of the basic stellar and wind properties of O-type stars. The first is to establish the near-infrared spectroscopic characteristics of O stars with known properties from studies at other wavelengths. One may then try to correlate the behaviour of these lines with their basic properties and see if one can retrieve the same information. This requires high-quality near-infrared spectra of a large sample of MK standard stars. So far, a good coverage is available only for the 2.0-2.2 \\mum\\ range (see Hanson et al. 1996, 2003). The second approach is to model the near-infrared lines using state of the art techniques, and then to study the dependence between spectral and basic properties.  So far modelling of the near-infrared spectral region has been done for extreme early-type stars, i.e., Luminous Blue Variables (Najarro et al. 1997), Of/WN stars (Crowther et al. 1995) and galactic centre objects (Najarro et al. 1994). This study aims at an improved treatment of normal O stars. With the use of sophisticated models that include a detailed treatment of stellar winds and that properly describe the spectrum, we study the predictions for the near-infrared regime. We investigate to what extent the near-infrared lines can be used to determine the spectral type, luminosity class, and mass loss of O-type stars, focussing on lines that are observable using ground-based instrumentation.  This paper is organized as follows: in Sect.~\\ref{sec_grid} we introduce our grid of models. Predicted equivalent widths (EW) of near-infrared lines are presented and their dependence on model parameters are discussed in Sect.~\\ref{sec_trend}. Model prediction are compared with observations gathered from the literature in Sect.~\\ref{sec_obs}. In Sect.~\\ref{sec_param}, we present near-infrared spectral classification schemes for O-type stars and a means to determine wind properties. We end with conclusions. ",
        "conclusions": "\\label{sec_conclu}  We have confronted observations of near-infrared spectra with predictions using the current state-of-the-art model atmosphere code {\\sc cmfgen} of Hillier \\& Miller (1998). This study was prompted by the fact that an increasing number of spectra of embedded O-type stars is presently becoming available, allowing the study of stars hidden beyond tens of magnitudes of circumstellar or foreground visual extinction. Though first attempts to make an inventory of the potential and predictive power of models in this wavelength range have been undertaken (Najarro et al. 1999), a systematic approach using a large grid of models was so far missing.  We summarize the main conclusions of this study.  \\begin{enumerate}  \\item The general trends and strengths of the lines are fairly compatible with observed properties. Near-infrared lines are typically formed in the transition region from photosphere to stellar wind, or even in the lower part of the wind (up to a few times the sonic velocity). The largest discrepancies are found for helium lines in supergiant spectra, which are systematically underpredicted. As low gravity models appear to reproduce these lines much better, it is expected that for these stars the connection of the photosphere to the wind is much smoother, i.e. the density changes more gradually, than is adopted in our models. (It does not imply that masses for supergiants are systematically underestimated).  \\item The models show a good correlation between optical and near-infrared helium line ratios in the J, H and K band. Notably, the \\hei~\\lam~1.2788 / \\heii~\\lam~1.1676, \\hei~\\lam~1.7007 / \\heii~\\lam~1.6921, and \\hei~\\lam~2.1136 / \\heii~\\lam~2.1885 \\mum\\ ratios appear good candidates. Application of these near-infrared ratios requires good quality data and are only applicable for stars with spectral type O4 to O8. Cooler stars do not show \\heii\\ lines. This is to be expected as the near-infrared continuum is formed further out in the atmosphere, relative to where the optical spectrum originates.  \\item The \\bra, \\brg, and \\heii\\,\\lam 3.0917 lines in giants and supergiants are predominantly formed in the stellar wind. Comparison with observations shows that the \\brg\\ lines are systematically underpredicted, while in most cases \\bra\\ is reproduced well. We have argued that \\brg\\ might be affected by an inhomogeneous density distribution or ``clumping'', while the stronger \\bra\\ line suffers much less from this effect, which would suggests a distance dependent clumping. Additional arguments concerning line formation depth and profile shape of \\brg, however, point also to presently unknown processes which have to be identified in future investigations.  \\item We find that the \\bra\\ line is a fairly good diagnostic for the wind density and we provide a means to correlate the measured equivalent width of this line to the wind density invariant $Q = \\mdot / R^{3/2} \\vinf \\teff^{2}$. The relation, however, depends on the \\heii\\ abundance in the wind, as the \\bra\\ line is contaminated by a potentially strong \\heii\\ line and may strongly affect mass loss predictions for Iaf stars.   \\end{enumerate}"
    },
    "0404/nucl-ex0404022_arXiv.txt": {
        "abstract": "{The Coulomb dissociation of $^8B$, measured with high precision by the GSI group, is in excellent agreement with the astrophysical cross section factor ($S_{17}$) measured by the Weizmann group with a $^7Be$ target. The GSI and Weizmann data are in good agreement with the Seattle data at high energies, but at low energies we observe a slight systematic (up to 2$\\sigma$) deviation, yet the Seattle group repeatedly rejects the CD method. We show that when plotting the slopes, they mis plotted one CD data point and exclude measured slopes that contradict their claim. Indeed the measured slope is essential to elucidate the d-wave correction to $S_{17}(0)$ that could be as large as 15\\%, and is the last open question that needs to be resolved before $S_{17}(0)$ can be quoted with an accuracy of 5\\% or better. We show that this goal has not been achieved (in spite of the strong claim of the Seattle group), since  currently there is no agreement among experiments as well as among theoretical models on the value of the slope. In addition, currently there is no theoretical framework within which (for example the Seattle) data can be analyzed and $S_{17}(0)$ extrapolated with the claimed high precision of 2.5\\%. This (last) issue of the slope and the d-wave correction must be resolved (by future measurements) so as to allow quoting $S_{17}(0)$ with an accuracy of 5\\% or better.} \\keyword{Coulomb dissociation, astrophysical cross section factor, solar $^8B$.} \\PACS{25.20.Dc, 25.70.De, 95.30-K, 26.30.+K, 26.65.+t}  \\begin{document} ",
        "introduction": "The astrophysical cross section factor of the $^7Be(p,\\gamma)^8B$ reaction, $S_{17}(0)$ as defined in Ref. \\cite{Adel}, is essential for predicting the $^8B$ solar neutrino flux \\cite{SSM}, which is now measured with 7.3\\% accuracy \\cite{SNO}, and extracted from a global fit of solar and reactor neutrino experiments \\cite{Bah03} with an accuracy of 4\\%. Hence it is essential to know $S_{17}(0)$ with a comparable accuracy. While recent claims suggest \\cite{Seatt} $S_{17}(0)$ measured with high accuracy, we demonstrate that this accuracy has not yet been achieved. ",
        "conclusions": ""
    },
    "0404/astro-ph0404437_arXiv.txt": {
        "abstract": "Monochromatic opacities from the Opacity Project (OP) (Seaton et al.) have been augmented by hitherto missing inner-shell contributions (Badnell \\& Seaton). OP  Rosseland-mean opacities, $\\kr$, are compared with results from OPAL (Iglesias \\& Rogers) for the elements H, He, C, O, S and Fe. The OPAL data are obtained from the website www-phys.llnl.gov/Research/OPAL/index.html.  Agreement for H is close everywhere except for the region of $\\log(T)\\simeq 6$ and $\\log(R)\\simeq -1$ ($R=\\rho/T_6^3$ where $\\rho$ is mass-density in g cm$^{-3}$ and $T_6=10^{-6}\\times T$ with $T$ in K). In that region $\\kr$(OPAL) is larger than $\\kr$(OP) by up to 13\\%. The differences are due to different equations of state (EOS). In the region concerned OP has the H ground state undergoing dissolution, leading to a small H-neutral ionization fraction, while OPAL has larger values for that fraction. A similar difference occurs for He at $\\log(R)\\simeq-1$ and $\\log(T)\\simeq 6.4$, where OP has the He$^+$ ground-state undergoing dissolution.  The OPAL website does not provide single-element Rosseland means for elements other than H and He. Comparisons between OP and OPAL are made for mixtures with $X=0.9$, $Z=0.1$ and $Z$ containing pure C, O or S. There are some differences: at the lower temperatures, say $\\log(T)\\leq 5.5$, due to differences in atomic data, with the OP R-matrix data probably being the more accurate; and at higher temperatures  mainly due to differences in level populations resulting from the use of different EOS theories.  In the original OP work, R-matrix data for iron were supplemented by data obtained using the configuration-interaction (CI) code {\\sc superstructure}. The experiment is made of replacing much of the original iron data with new data from the CI code {\\sc autostructure}. Inclusion of intercombination lines gives an increase in $\\kr$ of up to 18\\%.  The OPAL website does not allow for $Z$ containing pure iron. Comparisons are made for an iron-rich mixture, $X=0.9$, $Z=0.1$ and $Z$ containing C and Fe with C:Fe=2:1 by number fraction. There are some differences between OP and OPAL for that  case: the OP `$Z$-bump' in $\\kr$ is shifted to slightly higher temperatures, compared to OPAL.  Overall, there is good agreement between OP and OPAL Rosseland-mean opacities for the 6-elements, but there are some differences. Recent work (Bahcall et al.) has shown that helioseismology measurements give a very accurate value for the depth of the solar convection zone, $R_{\\rm CZ}$, and that solar models give agreement with that value only if opacities at $R_{\\rm CZ}$ are about 7\\% larger than OPAL values. For the 6-element mix at $R_{\\rm CZ}$ we obtain $\\kr$(OP) to be larger than $\\kr$(OPAL) by 5\\%. ",
        "introduction": "Rosseland-mean opacities from the Opacity Project (OP), as originally presented in \\cite{symp}, were in good agreement with those from the OPAL project \\cite{opal} over much of the temperature--density plane, but were smaller than those from OPAL at high temperatures and densities. Iglesias and Rogers \\cite{IR96} offered the explanation that OP was missing some data for inner-shell transitions and that was confirmed in a recent paper by the present authors \\cite{bs03}. The results of \\cite{IR96} and \\cite{bs03} were for a mixture of 6 elements, H, He, C, O, S and Fe with abundances, by number-fractions, given in Table 1: we refer to that as the 6-element mix; mass-fractions are $X=0.7$ for H, $Y=0.28$ for He, $Z=0.02$ for `metals'.  Fig. 1 shows the level of agreement between OP and OPAL, for that mix, as obtained in \\cite{bs03}: values of $\\log(\\kappa_{\\rm R})$, where $\\kappa_{\\rm R}$ is Rosseland-mean opacity in cm$^2$ g$^{-1}$, are plotted against $\\log(T)$ for 5 different values of $\\log(R)$ where $R=\\rho/T_6^3$, $\\rho=$  density in g cm$^{-3}$ and $T_6=10^6\\times T$ with $T$ in K (fig. 1 of \\cite{symp} shows the behaviour of $\\log(R)$ for a few typical stellar models). The OPAL data are obtained from the OPAL website \\cite{web} and the OP data are with inclusion of the inner-shell contributions discussed in \\cite{bs03}. It is seen that the agreement between OP and OPAL is fairly good in all cases but that there are some differences. The purpose of the present paper is to make more detailed comparisons of OP and OPAL for the six elements H, He, C, O, S and Fe. Work on the inclusion of OP inner-shell data for other elements is in progress. ",
        "conclusions": "Our main conclusion is that there is good general agreement between OP and OPAL opacities for the 6 elements H, He, C, O, S and Fe. We expect to obtain equally good agreement for other elements once work is complete on the calculation of further inner-shell data.  There are some indications that the OPAL interpolation procedures described in \\cite{corr} may introduce some errors near the edges of the domains of validity which are claimed.  At lower temperatures, $\\log(T)\\lesssim 5.5$, there are some modest differences between OP and OPAL which may be a consequence of the greater sophistication of the OP atomic-data work. At higher temperatures and higher densities, $\\log(R)\\gtrsim -2$, the OP opacities tend to be larger than those from OPAL in consequence of the use of different equations of state. There is some helioseismology evidence that the OP results are the more accurate."
    },
    "0404/astro-ph0404092_arXiv.txt": {
        "abstract": "{We report on an analysis of the polarization of single pulses of PSR B0329+54 at 328~MHz. We find that the distribution of polarization orientations in the central component diverges strongly from the standard picture of orthogonal polarization modes (OPMs), making a remarkable partial annulus on the Poincar\\'{e} sphere.  A second, tightly clustered region of density appears in the opposite hemisphere, at a point antipodal to the centre of the annulus.  We argue that this can be understood in terms of birefringent alterations in the relative phase of two elliptically polarized propagation modes in the pulsar magnetosphere (i.e. generalised Faraday rotation). The ellipticity of the modes implies a significant charge density in the plasma, while the presence of both senses of circular polarization, and the fact that only one mode shows the effect, supports the view that refracted ordinary-mode rays are involved in the production of the annulus.  At other pulse longitudes the polarization (including the circular component) is broadly consistent with an origin in elliptical OPMs, shown here quantitatively for the first time, however considerable non-orthogonal contributions serve to broaden the orientation distribution in an isotropic manner. ",
        "introduction": "\\label{sec:intro} Whenever sufficient sensitivity is available, radio pulsar emissions are seen to be rich in phenomenology. Their polarization is no exception to this rule. The dependence of linear polarization position angle on pulse longitude (i.e. rotational phase) can, for some pulsars, be explained as arising in the vicinity of the magnetic pole, polarized linearly at the position angle of the sky projection of magnetic field lines \\citep{rc69a}. For other pulsars this is not the case, and for some of these the distribution of position angles (PAs) in individual pulses has been shown to be bimodal about two values separated by 90$\\degr$ --- so-called orthogonal polarization modes (OPMs; e.g. \\citealt{mth75,brc76,crb78, br80,scr+84,scw+84}). Many pulsars also show $90\\degr$ ``jumps'' in their position angle profiles, a fact that received explanation in the discovery of OPMs through a longitude dependence of the relative intensities of the modes, which themselves tend to have position angle swings consistent with the magnetic pole model of \\citet{rc69a}. These studies of OPMs also found evidence for deviations from orthogonality in the fact that PA distributions were broader than expected and/or their peaks were not separated by $90\\degr$.  Attempts have been made to explain these deviations from orthogonality by means of superposition of two modes that are not orthogonal due to origins on different field lines and subsequent birefringent refraction (e.g. \\citealt{scr+84,mck03a}), by means of the superposition of a range of modal orientations arising from a distribution of field lines that are visible due to their finite beam width \\citep{gl95}, or {\\changed due to the presence of two instantaneously orthogonal modes, the orientation of which varies with time due to coherent wave coupling effects} (e.g. \\citealt{cr79,lp99}). Of these, only the first and third scenarios have been quantitatively tested. \\citet{mck03a} showed that the distribution of PAs, and the shape of the average PA curve of PSR B2016+28 are consistent with the superposition of two non-orthogonal modes.  On the other hand, \\citet{pet03} showed that the PA and circular polarization of the average pulse profiles of PSR B0355+54 and PSR B0628$-$28 are consistent with the predictions of \\citet{lp99} for alterations in the polarization due to coherent effects. However, neither work examines whether the data are also consistent with other models, and indeed the statistics considered in each case are not well suited to distinguishing between models.   The question of the origin of circular polarization in pulsars is also not addressed by the magnetic pole model, in which the polarization is expected to be linear due to the high magnetic field strength.  The fact that orthogonal linear polarization states have been found to be associated with opposite signs of circular polarization leads naturally to the suggestion that the modes are in fact elliptically polarized orthogonal modes \\citep{crb78}, although to our knowledge, to date no observational tests of the expected proportionality between linear and circular polarization under this hypothesis have been made. In explaining the origin of elliptical OPMs, most authors point to the so-called polarization-limiting region (PLR), where birefringent propagation effects no longer significantly alter the relative phase of the modes, as the determinant of the observed polarization \\citep{cr79}.  It has been proposed that either the propagation modes themselves are elliptical at this point, requiring a net charge density in the plasma \\citep{cr79,am82, hlk98}, or that weak birefringence in the vicinity of the PLR causes initially linearly polarized rays to suffer changes in their polarization if their position angle deviates from that of the local linear modes, for example due to rotational aberration or refraction \\citep{cr79,lp99,pl00,pet01,pet03}.  In this work we make a detailed study of the polarization of single pulses from PSR B0329+54, employing new techniques in the hope of placing greater constraints on the nature and origin of ellipticity and non-orthogonality of pulsar polarization. ",
        "conclusions": "We have performed a detailed analysis of the polarization of individual pulses from PSR B0329+54 at 328~MHz. The use of statistics involving the ellipticity angle instead of Stokes $V$ or $V/I$ enabled us to reveal for the first time the rich structure of the circular polarization distribution, which is correlated with the position angle distribution.  Displaying the distribution of polarization orientations in an equal-area projection, a remarkable structure was uncovered for pulse longitudes near the centre of the average pulse. The annulus-like form of the distribution in one mode is, in our view, indicative of Generalised Faraday Rotation (GFR) in the pulsar magnetosphere, while the fact that it apparently only affects one polarization mode, and is capable of producing either sense of circular polarization, is taken as an indication that refraction of the ordinary mode is taking place. The apparent ellipticity of the polarization state upon which the annulus centres and orthogonal to the typical polarization of states dominated by the other mode are taken to indicate that the GFR occurs in a region with elliptically polarized plasma propagation modes, indicating a net charge density in the plasma.  Through the analysis of the covariance of the Stokes parameters, by means of eigenvector decomposition, we have shown  quantitatively for the first time that  the circular polarization is consistent with an origin in elliptically polarized orthogonal polarization modes. Moreover, we find that a significant apparently randomly polarized component dilutes the purity of OPM-driven variations in polarization state, as also found recently by \\citet{mck04} for  PSR B1929+10 and PSR B2020+28 at 1404~MHz.  In this work we have shown how to detect and characterise deviations from OPM in more powerful ways than previously available. Future application of these techniques to other pulsars over a broad frequency range should allow renewed progress in the resolution of decades-long debate about the origin of pulsar polarization."
    },
    "0404/astro-ph0404571_arXiv.txt": {
        "abstract": "{We present the results of VLBA observations of the radio galaxy 3C\\,317, associated with the cD galaxy UGC 09799 at the centre of the cooling core cluster of galaxies A2052.\\\\ These observations were carried out at 1.7 GHz, 4.9 GHz and 8.3 GHz, in polarimetric mode, and allowed us to  image the parsec scale region of the source. Our analysis suggests that the nucleus of 3C\\,317 hosts a very young radio source. The estimated radiative age for the radio structure within the inner 10 pc is $\\sim$ 170 yr. Given the existence of extended radio emission on the arcsecond scale, we suggest that 3C\\,317 is a restarted radio galaxy. The implications of this result in the light of the interaction between radio plasma and thermal hot gas in clusters of galaxies are briefly discussed.\\\\ ",
        "introduction": " ",
        "conclusions": "Our polarimetric VLBA multifrequency observations of 3C\\,317, the central radio galaxy in the cooling core cluster A\\,2052, revealed a compact double morphology over a scale of few mas. The observational properties are consistent with the idea that the nucleus of 3C\\,317 is a young CSO. In particular, the peak in the radio spectrum at $\\sim$ 2.6 GHz, the lack of polarisation, the orientation almost in the plane of the sky, and the values of the intrinsic physical parameters are common to CSS and to GPS sources. The radiative age of the inner source structure, estimated from our observations, is only $\\sim$ 170 yr.\\\\ Assuming that the radiative age and the dynamic age of the source are similar, we predict that the parsec--scale lobes in 3C\\,317 are moving away from each other at a speed $v_{sep}\\sim 0.14 ~hc$.\\\\  On the arcsecond scale, the nucleus of 3C\\,317 is surrounded by an extended amorphous ``halo'' of low brightness radio emission, which spatially anticorrelates with the cluster X--ray emission. In particular, the radio plasma is found to fill gaps of X--ray brightness, supporting the recent suggestion that the radio plasma sweeps the thermal intracluster gas and compresses it into bright shells. An intriguing possibility is that the extended amorphous halo is the result of reacceleration processes due to MHD turbulence in the cooling flow.\\\\ The young age estimated for the radio nucleus of 3C\\,317, coupled with the existence of large scale radio emission, leads us to suggest that the radio galaxy is characterised by intermittent radio emission in the nucleus, and that a new phase of activity has just started. This result is relevant not only in the light of our understanding of radio source birth and evolution, but it is also an important piece of information for our knowlegde of the interaction between the non thermal radio emission and the hot thermal gas in clusters of galaxies."
    },
    "0404/hep-ph0404272_arXiv.txt": {
        "abstract": "Neutrinos from far-away sources annihilating at the Z resonance on relic neutrinos  may give origin to the extreme-energy cosmic rays (EECR). If ``Z-bursts'' are responsible for the EECR events, then we show that the non-observation of cosmic ray events at energies above 2 $\\times 10^{20}$ eV by the AGASA Collaboration implies a lower bound $\\sim 0.3$~eV on the relic neutrino mass. Since this mass exceeds the mass-squared differences inferred from oscillation physics, the bound in fact applies to all three neutrino masses. Together with the upper bound provided by comparisons of the CMB anisotropy with large-scale structure, this bound leaves only a small interval for neutrino masses around 0.3 eV, if Z-bursts are to explain the existing EECR events. ",
        "introduction": "The existence of extreme-energy cosmic rays (EECR) with energies above the Greisen-Zatsepin-Kuzmin (GZK) cutoff~\\cite{gzk}  of about $5\\times 10^{19}$ eV, presents an outstanding problem~\\cite{data}.  Nucleons and photons with those energies have short attenuation lengths and could only come from distances of 100 Mpc or less~\\cite{50Mpc,40Mpc}, while plausible astrophysical sources for those energetic particles are much farther away.  Data from the HiRes experiment \\cite{HiRes} brought the violation of the GZK cutoff into question. Yet the EE events in even the HiRes data set remain unexplained since the local Universe ($\\sim 100$~Mpc) is devoid of strong candidate sources. This controversy will be solved conclusively by the Pierre Auger hybrid observatory~\\cite{Auger}, perhaps as soon as summer of 2005. Here we assume that the published AGASA spectrum is correct. Interestingly, it is the absence of events above $2\\times 10^{20}$~eV in these data, rather than the presence of events above $E_{\\rm GZK}\\sim 5\\times 10^{19}$~eV, that motivates this work.  Among the solutions proposed for the origin of the highest energy events observed by AGASA, an elegant and economical solution  to this problem is the ``Z-burst'' mechanism: annihilation at the $Z$-resonance of extreme-energy neutrinos ($\\nu_{\\rm EE}$) coming from remote sources, and relic background neutrinos within $\\sim 50$~Mpc of Earth, produces  the nucleon and photon EECRs~\\cite{Zbursts}. The observed EECRs from the Z-bursts are the emission analogues of the absorption features (Z-dips) predicted long ago~\\cite{Zdips}. One of the most appealing features of the $Z$-burst mechanism is that the energy scale of $\\gsim 10^{20}$~eV at which the unexpected events have been detected, is generated naturally. The $Z$-resonance occurs when the energy of the incoming $\\nu_{\\rm EE}$ is \\begin{equation} \\Eres = \\frac{M_Z^2}{2~m_\\nu} = 4\\times 10^{21}{\\rm eV}\\,\\left(\\frac{\\rm eV}{m_\\nu}\\right)\\,, \\label{Eres} \\end{equation} where $m_\\nu$ is any of the three masses of the relic neutrinos. Given the lower limit $\\sim 0.04$~eV deduced from atmospheric oscillations for the heaviest neutrino mass, at least one $\\Eres$ is below $10^{23}$~eV.  Since the individual energies of the nucleons and photons emerging from the Z-burst cannot exceed the total energy of the burst, $\\Eres$ is the new end-point of the EECR energy-spectrum in this mechanism. Partitioning this burst energy among the $\\langle N\\rangle \\sim 40$ final state particles, one arrives at precisely the primary energies needed to produce events observed above the GZK energy.  Combining their recent  measurements of the anisotropies of the Cosmic Microwave Background (CMB) radiation with data on the large scale structure of the Universe, WMAP~\\cite{WMAP} produced a strong limit on neutrino masses, $\\Sigma_i m_{\\nu_i} < 0.69$~eV. Since a single neutrino mass above $\\sim 0.04$~eV implies near mass-degeneracy for all three active neutrinos (given the mass-squared splittings from neutrino oscillation data, described below) one has $m_\\nu < 0.23$~eV at the 95\\% CL. However, objections~\\cite{objections} to the priors assumed, or to the data sets included, have led to a relaxed bound $\\Sigma_i m_{\\nu_i} < 1$~eV, or $m_\\nu < 0.33$~eV. A subsequent analysis~\\cite{BMT} which includes WMAP, 2dF, and SDSS data, and another which includes these and galaxy cluster data~\\cite{ASB}, have arrived at the bound, $\\Sigma_i m_{\\nu_i} \\lsim 0.7$~eV. Still another analysis~\\cite{CLP} with CMB and LSS data finds $\\Sigma_i m_{\\nu_i} \\lsim 1$~eV, but finds $\\Sigma_i m_{\\nu_i} \\lsim 0.6$~eV when priors from supernova data and the Hubble Key Project are included. These newer bounds are very similar to the original WMAP bound.  In this work we focus on a particular feature of the AGASA spectrum, namely the end-point energy.  Requiring that the Z-burst mechanism not overproduce events beyond this AGASA end-point, we derive a lower bound on the neutrino mass. The significance of an end-point energy for constraining model fluxes has been noted much earlier in the context of topological defect decay~\\cite{PJ-PS}. A more ambitious project would be to actually fit the neutrino mass to the full AGASA spectrum. This has been done, by Fodor, Katz and Ringwald (FKR)~\\cite{FKR}. However, the allowed range of the neutrino mass that results is not tight. Rather, it depends sensitively on how one parameterizes the transition from non-burst spectrum to burst spectrum near and above the ankle. FKR find $\\pm 2\\sigma$ fitted mass ranges of 0.1 to 1.3~eV if all EE events are assumed to originate with Z-bursts, 0.02 to 0.8~eV if an additional extra-galactic source of EE protons is allowed, and 1 to 7~eV if an additional Galactic-halo source of EE protons is allowed.  Super-Kamiokande has provided a strong evidence for the oscillation in atmospheric showers of two neutrino species with mass-squared splitting $\\delta m^2 $ = $m_3^2-m_2^2$ = $(1-3) \\times 10^{-3}$~eV, consisting of nearly equal amounts of $\\nu_\\mu$ and $\\nu_\\tau$, and little or no $\\nu_e$~\\cite{SK}. If neutrino masses are hierarchical, like the other leptons and quarks, then $\\sqrt{\\delta m^2} \\simeq 0.04\\,{\\rm eV} \\equiv m_{\\rm SK}$ \\underline{\\sl is} the mass of the heavier of the two oscillating neutrinos (call it $\\nu_{\\rm SK}$). Some previous works~\\cite{gk1,gv} on Z-burst production of the EECRs have made this assumption. However we will show in this work that $m_\\nu\\sim 0.04$~eV is not compatible with the AGASA end-point spectrum if Z-bursts produce the EECR. The reason is that with $m_\\nu = 0.04$~eV, the new EECR cutoff $\\Eres \\simeq 10^{23}$~eV predicts too many super GZK EECR events beyond the AGASA end-point.  As the relic neutrino mass increases, $\\Eres\\propto 1/m_\\nu$ decreases, and the features in the EE spectrum move progressively to lower energies. In particular, the number of events predicted at high energy decreases, and the number beyond the AGASA end-point becomes compatible with zero for $m_\\nu \\gsim0.3$~eV. Note that our result is much sharper than, but compatible with, the allowed mass ranges of the FKR models with no additional EE proton source, or with an additional extra-galactic source of EE protons.  In the next section we present our simulations and the resulting spectrum of EECR in detail. ",
        "conclusions": "The non-observation  of cosmic ray events at energies higher than $\\sim 2\\times 10^{20}$ eV by AGASA provides a lower bound of about 0.3~eV on relic neutrino masses, if Z-bursts are responsible for the EECRs. This bound together with comparisons of the microwave background anisotropy to today's large-scale structure, leave a small interval for the relic neutrino mass, centered around 0.3 eV, if Z-bursts explain the EECRs. The $\\sim 0.3$~eV mass provides a very accessible signal for neutrinoless double beta-decay experiments (assuming neutrinos are Majorana particles), and is in fact compatible with a reported positive signal~\\cite{Klapdor}. This $\\sim 0.3$~eV mass interval may also be probed by the future KATRIN tritium end-point experiment~\\cite{katrin}, and by expected improvements to the astrophysical bound from weak lensing of distant galaxies by intervening matter distributions~\\cite{AD02}, and from mining future Planck measurements of the CMB~\\cite{PlanckCMB}. We note that our small allowed range for $m_\\nu$ is incompatible with an upper bound of 0.15~eV that arises from invoking the see-saw mass spectrum as a source of successful early-Universe baryogenesis~\\cite{leptogenesis}.   We have not included in this work any possible effects of extragalactic magnetic fields. For the analysis of this paper to hold as is, these fields should be sufficiently small, $\\lsim 10^{-9}~G$. A survey of the literature shows that extragalactic fields much larger than $10^{-9}~G$ are not expected, except in regions around galaxy clusters. In any event,  large extragalactic magnetic fields which would reduce or eliminate the EE photon component of the Z-bursts (thus lowering the neutrino mass we obtain) would also lead to larger required  EE neutrino fluxes at the Z-resonance energy. This possibility is rejected by the GLUE and FORTE bound.  We have shown that $Z$-bursts  may account  not only for the super GZK events, but for the ``ankle\" and all EECR events above it. In doing so, the $Z$-burst mechanism meets the ``naturalness\" requirements of isotropy and spectral smoothness. In our simulation we found the ``ankle\"  close to $E_{ankle} = 10^{19.0}$~eV as measured by AGASA.\\ We found that most EECR primaries above the ankle should be nucleons up to about $10^{20.0}$~eV and nucleons and photons  at higher energies. We also found that primary nucleons do accumulate at the GZK cutoff energy, which could  account for the slight accumulation seen in the data (see, however, \\cite{depression}).  With $m_\\nu\\sim 0.3$~eV needed to accommodate the AGASA end-point, the Z-burst hypothesis predicts a new absolute energy cutoff for EECRs, at $\\Eres \\simeq 1.3\\times 10^{22}~eV$. However, due to the large multiplicity in Z-decay, the vast majority of incident primaries are expected to carry only a few percent of this cutoff energy.  In summary, we believe that the $Z$-burst mechanism provides a plausible explanation to the puzzle of extreme-energy cosmic rays, not only for the super GZK events, but for the ``ankle'' and all events above it. But the hypothesis will stand or fall on the required nearness of $m_\\nu$ to 0.3~eV, and/or on the proximity of the required neutrino flux at $\\Eres=1.3\\times 10^{22}$~eV to existing flux limits.  If Z-bursts are invalidated as the source of the observed EECRs, then we would encourage future, larger experiments (e.g., Auger, EUSO) to continue the search for Z-burst signatures, for they provide the only known window to discovery of the relic neutrino background. The Z-burst mechanism is entirely Standard Model particle physics coupled with Standard Hot Big-Bang Cosmology.  Consequently, the mechanism itself cannot be ruled out. However, the Z-burst rate depends linearly on the resonant-energy neutrino flux which Nature provides.  Nature may not be generous.  \\vspace{5 mm}   This work was supported in part by NASA grant NAG5-13399, the US Department of Energy grants DE-FG03-91ER40662, Task C, and DE-FG05-85ER40226, and Vanderbilt University's Discovery Grant and sabbatical programs. G.V. was also supported by an award from the Research Corporation. We thank A. Kusenko, S. Nussinov and D. Semikoz for many valuable discussions, P. Biermann and E. Roulet for comments and suggestions, and the Aspen Center for Physics for a supportive environment during part of this work."
    },
    "0404/hep-ph0404012_arXiv.txt": {
        "abstract": "We re-examine the model of Natural inflation, in which the inflaton potential is flat due to shift symmetries.  The original version of the model, where the inflaton is a pseudo Nambu-Goldstone boson with potential of the form $ V(\\phi) = \\Lambda^4 [1 \\pm \\cos(\\phi/f)]$, is studied in light of recent data. We find that the model is alive and well. Successful inflation as well as data from the Wilkinson Microwave Anisotropy Probe require $f > 0.6 m_{\\rm Pl}$ (where $m_{\\rm Pl} = 1.22 \\times 10^{19}$ GeV) and $\\Lambda \\sim m_{GUT}$, scales which can be accommodated in particle physics models. The detectability of tensor modes from natural inflation in upcoming microwave background experiments is discussed. We find that natural inflation predicts a tensor/scalar ratio within reach of future observations. ",
        "introduction": "The inflationary universe model was proposed \\cite{guth} to solve several cosmological puzzles, the horizon, flatness, and monopole problems, via an early period of accelerated expansion. To satisfy a combination of constraints on inflationary models, in particular, sufficient inflation and microwave background anisotropy measurements \\cite{WMAP} of density fluctuations, the potential for the inflaton field must be very flat.  For a general class of inflation models involving a single slowly-rolling field (including new \\cite{newinf}, chaotic \\cite{chaotic}, and double field inflation \\cite{doublefield}, the ratio of the height to the (${\\rm width})^4$ of the potential must satisfy \\cite{afg} \\begin{equation} \\label{eq:ratio} \\chi \\equiv \\Delta V/(\\Delta \\phi)^4 \\le {\\cal O}(10^{-6} - 10^{-8}) \\, , \\end{equation} where $\\Delta V$ is the change in the potential $V(\\phi)$ and $\\Delta \\phi$ is the change in the field $\\phi$ during the slowly rolling portion of the inflationary epoch.  Thus, the inflaton must be extremely weakly self-coupled, with effective quartic self-coupling constant $\\lambda_{\\phi} < {\\cal O}(\\chi)$ (in realistic models, $\\lambda_{\\phi} < 10^{-12}$).  The small ratio of mass scales required by Eq. (\\ref{eq:ratio}) quantifies how flat the inflaton potential must be and is known as the ``fine-tuning'' problem in inflation.  Three approaches have been taken toward this required flat potential characterized by a small ratio of mass scales.  First, some simply say that there are many as yet unexplained hierarchies in physics, and inflation requires another one.  The hope is that all these hierarchies will someday be explained.  In these cases, the tiny coupling $\\lambda_{\\phi}$ is simply postulated {\\it ad hoc} at tree level, and then must be fine-tuned to remain small in the presence of radiative corrections.  But this merely replaces a cosmological naturalness problem with unnatural particle physics. Second, models have been attempted where the smallness of $\\lambda_{\\phi}$ is protected by a symmetry, {\\it e.g.,} supersymmetry. In these cases (e.g., \\cite{hrr}), $\\lambda_{\\phi}$ may arise from a small ratio of mass scales; however, the required mass hierarchy, while stable, is itself unexplained.  In addition, existing models have limitations.  It would be preferable if such a hierarchy, and thus inflation itself, arose dynamically in particle physics models.  Hence, in 1990 we proposed a third approach, Natural Inflation \\cite{natural}, in which the inflaton potential is flat due to shift symmetries.  Nambu-Goldstone bosons (NGB) arise whenever a global symmetry is spontaneously broken.  Their potential is exactly flat due to a shift symmetry under $\\phi \\rightarrow \\phi + {\\rm constant}$. As long as the shift symmetry is exact, the inflaton cannot roll and drive inflation, and hence there must be additional explicit symmetry breaking.  Then these particles become pseudo-Nambu Goldstone bosons (PNGBs), with ``nearly'' flat potentials, exactly as required by inflation.  The small ratio of mass scales required by Eq. (\\ref{eq:ratio}) can easily be accommodated. For example, in the case of the QCD axion, this ratio is of order $10^{-64}$.  While inflation clearly requires different mass scales than the axion, the point is that the physics of PNGBs can easily accommodate the required small numbers\\footnote{For example, in `invisible' axion models \\cite{invisible} with Peccei-Quinn scale $f_{PQ} \\sim 10^{15}$ GeV, the axion self-coupling is $\\lambda_a \\sim (\\Lambda_{QCD}/f_{PQ})^4 \\sim 10^{-64}$. (This simply reflects the hierarchy between the QCD and GUT scales, which arises from the slow logarithmic running of $\\alpha_{QCD}$.)  Due to the nonlinearly realized global symmetry, the potential for PNGBs is exactly flat at tree level.  The symmetry may be explicitly broken by loop corrections, as in schizon \\cite{hillross} and axion \\cite{axion} models.  In the case of axions, for example, the PNGB mass arises from non-perturbative gauge-field configurations (instantons) through the chiral anomaly.  When the associated gauge group becomes strong at a mass scale $\\Lambda$, instanton effects give rise to a periodic potential of height $\\sim \\Lambda^4$ for the PNGB field \\cite{instanton}.  Since the nonlinearly realized symmetry is restored as $\\Lambda \\rightarrow 0$, the flatness of the PNGB potential is natural in the sense of 't Hooft \\cite{hooft}.}.  We first proposed this model and performed a simple analysis in \\cite{natural}. Then, in 1993, we followed with a second paper which provides a much more detailed study \\cite{natural2}.  The results of Section III of the second paper, which presents a careful analysis of the dynamics of the natural inflaton, are of particular relevance here.  Many types of candidates have subsequently been explored for natural inflation.  For example \\cite{kyy} used shift symmetries in Kahler potentials to obtain a flat potential and drive natural chaotic inflation in supergravity.  Additionally, \\cite{hccr} examined natural inflation in the context of extra dimensions and \\cite{kapwei} used PNGBs from little Higgs models to drive hybrid inflation.  Also, \\cite{firtye} and \\cite{hsukallosh} use the natural inflation idea of PNGBs in the context of braneworld scenarios to drive inflation. Freese \\cite{term} suggested using a PNGB as the rolling field in double field inflation \\cite{doublefield} (in which the inflaton is a tunneling field whose nucleation rate is controlled by its coupling to a rolling field).  We will focus in this paper on the original version of natural inflation, in which there is a single rolling field; we will comment further on other variants of natural inflation in Section 6.  In the current paper we show that the original proposal of natural inflation is live and well, contrary to recent criticisms (which we address in Section 3).  In particular, the single-field version of the model is successful for $f > 0.6 m_{Pl}$ (and does not require $f >> m_{Pl}$, contrary to the claims of \\cite{hccr}). A second focus of the current paper is to discuss tests of natural inflation from existing and upcoming data from microwave background experiments.  Recent results from the Wilkinson Microwave Anisotropy Probe \\cite{WMAP} are used to constrain our model, and predictions are made for upcoming experiments such as the PLANCK satellite which will begin taking data in 2007.  We begin in Section 2 by reviewing the basic idea of natural inflation.  In Section 3 we present results of the evolution of the scalar field driving inflation, including explicit numerical calculation of the evolution of the scalar field in its potential.  In Section 4 we discuss density fluctuations, and find the constraint on the potential due to comparison with WMAP data. In Section 5 we compute the tensor modes from natural inflation, and discuss their detectability in upcoming microwave background experiments.  In Section 6, we conclude with a discussion of the pros and cons of having a model in which the width of the potential is of order the Planck scale. ",
        "conclusions": "In conclusion, natural inflation is alive and well.  Recent WMAP data constrain the width of the potential to be $f >0.6 m_{\\rm Pl}$, and our predictions show that upcoming CMB observations such as the PLANCK satellite may be able to see the tensor modes.  In this section, we discuss the pros and cons of $f \\sim m_{\\rm Pl}$, as well as comment on some of the literature of PNGB models using shift symmetries.  Although it is {\\it not} true that the original model of natural inflation requires $f>>m_{\\rm Pl}$ for the width of the potential, it does require $f$ to be of order $m_{\\rm Pl} \\sim 10^{19}$ GeV.  In fact, virtually all 4D inflationary models require $f\\sim m_{\\rm Pl}$ and the height of the potential $\\sim m_{GUT}$, in order to satisfy the simultaneous requirements of sufficient inflation and the right amplitude of density perturbations; this fact is emphasized by the conclusions of Ref. \\cite{afg}. However, the height of the potential is generically of order the GUT scale, far enough below the Planck scale that we can safely ignore quantum gravitational effects on the background evolution. However, energy density is not the only issue.  In \\cite{kammarch, holman}, it was argued that Planck-scale physics results in the violation of all global symmetries, including the Peccei Quinn symmetry of the axion and the underlying symmetry from which we derive the PNGB inflaton.  Wormholes are suggested as one mechanism for this violation, and black holes another (as a consequence of black hole no-hair theorems, the global charge of a black hole is not defined). The authors argue that, as a consequence, one is required to add all higher dimension operators (suppressed by powers of $m_{\\rm Pl}$) consistent with the symmetries of the full theory, which then does not respect global symmetries.  One should include terms of the form: \\begin{equation} {\\cal L} = {1 \\over 2} m^2 \\phi^2 + \\lambda \\phi^4 + \\sum_{n = 6}^{\\infty}{\\lambda_n \\left(\\phi^n \\over m_{\\rm Pl}^{n-4}\\right)}. \\end{equation} Without a complete theory of quantum gravity, the validity of these arguments is not clear.  If true, then the idea of using a PNGB directly as the inflaton would fail; however, the axion also could not exist and we would have no theory at all to escape the strong CP problem in QCD.  In \\cite{lyth}, Lyth discussed the failure of effective field theory if the width of an inflationary potential approaches $m_{\\rm Pl}$. Again, if inflation is to be formulated as an effective low-energy field theory, he argues that we expect additional nonrenormalizable operators in the Lagrangian to be suppressed by inverse powers of $m_{\\rm Pl}$ as above.  Then if observational constraints require the field to travel a distance $\\Delta \\phi \\sim m_{\\rm Pl}$, the effective field theory will begin to break down due to radiative corrections from the nonrenormalizable operators. Such a theory rapidly becomes inconsistent as $\\Delta \\phi \\gg m_{\\rm Pl}$.  Motivated by the desire to evade these issues, in 1995 Kinney and Mahantappa \\cite{km} constructed natural inflation models in which symmetries suppress the mass terms and the potential is of the form $V \\simeq 1 - \\sin^4{(\\phi / f)} \\sim 1 - \\phi^4$.  Then one automatically obtains an effective width of the potential $f << m_{\\rm Pl}$ \\cite{knox92}. One does so, however, at the expense of an unobservably small tensor component in the CMB.  Natural inflation has been implemented in the context of extra dimensions with varying degrees of success.  Recently, Arkani-Hamed {\\it et al.} \\cite{hccr} examined natural inflation in the context of extra dimensions, and also found models for which the mass terms were suppressed by a symmetry with $f << m_{\\rm Pl}$, similar to the work of \\cite{km}.  Focusing on a Wilson line in a fifth dimension, Arkani-Hamed {\\it et al.} alternatively suggested that one might obtain large $f >> m_{\\rm Pl}$ if the inflaton is the extra component of a gauge field propagating in the bulk.  However, Banks {\\it et al.} examined the general question of whether it is ever possible to obtain large values of $f>> m_{\\rm Pl}$ in string theory \\cite{Banks:2003sx}.  While their study was not exhaustive, it strongly suggests that it is not possible.  Generically there is a T-dual description in which small radii become large and the value of $f$ is small ($f<< m_{\\rm Pl}$); hence the model of Arkani-Hamed {\\it et al.} does not succeed in providing large $f$ inflation.  In addition, in a variety of regions in moduli space Banks {\\it et al.} find that there are low action instantons which give rise to rapidly varying contributions to the potential that effectively rescale the value of $f$ to the Planck scale.  Although they do not provide an exhaustive proof, Banks {\\it et al.} suggest the very strong statement that natural inflation cannot work in the context of string theory for $f>>m_{\\rm Pl}$.  However, natural inflation with $f \\sim m_{\\rm Pl}$, as discussed throughout this paper, is fine.  Natural inflation has been implemented in brane-world scenarios as well.  Shift symmetries have been studied in brane inflation by Firouzjahi and Tye \\cite{firtye} and in the work of Hsu and Kallosh \\cite{hsukallosh}. The four dimensional effective field theory description of some braneworld scenarios is likely to be described by the physics in this paper.  The shift symmetries of natural inflation have also been used in multiple field models.  Freese \\cite{term} suggested using a PNGB as the rolling field in double field inflation \\cite{doublefield} (in which the inflaton is a tunneling field whose nucleation rate is controlled by its coupling to a rolling field).  Kaplan and Weiner have examined natural inflation-like models in the context of ``little'' fields \\cite{kapwei}.  Kawasaki {\\it et al.} proposed a supergravity inflation model in which the inflaton potential is flat due to a shift symmetry, again utilizing the basic idea of natural inflation.  Choi {\\it et al.} have discussed thermal inflation in the context of radiatively generated axion scale in supersymmetric axion models \\cite{choi97}.  While arguments based on theoretical prejudice are useful guidelines for model building, the ultimate test is an observational one. It is a remarkable coincidence that the borderline between consistent and inconsistent effective field theory models for inflation is roughly the same as the borderline between whether or not tensor modes are, in a practical sense, observable \\cite{kinney2003}. In order for the tensor/scalar ratio to be large enough to be measured by foreseeable future experiments, the width of the potential must be of order $f \\sim m_{\\rm Pl}$ or larger. Natural inflation is still very much in the running as a candidate model for the early universe.  These models are also especially attractive from a particle physics perspective because they possess a hierarchy of scales which is stable against radiative corrections. Such a hierarchy is necessary for the suppression of density perturbations, but, in other models, is typically left to an unexplained fine tuning of the inflaton self-coupling to order $\\lambda \\sim 10^{-14}$. In Natural Inflation this heirarchy of scales arises naturally. Finally, the simplest models of natural inflation predict a large enough gravitational wave component that a detection by advanced CMB measurements will be possible."
    },
    "0404/astro-ph0404182_arXiv.txt": {
        "abstract": "We analysed 246 scanned ACO clusters and we do not find any dependence of galaxy cluster ellipticity on redshift. ",
        "introduction": "The recent substantial evolution in the ellipticity of rich galaxy clusters was noted by Mellot, Chambers \\& Miller (2001). This result has been supported by Plionis (2002). In particular, he argues that the  ellipticity of  ACO clusters is correlated with their redshifts. Using, among others, 309 ACO clusters taken from APM, for which about 180 have determined redshifts, he claimed the existence of a significant ellipticity - redshift dependence. ",
        "conclusions": "\\label{sec:concl}  It is clear that we do not observe any relation between ellipticities and redshifts. This differs greatly from the Plionis (2002) result. However, the claimed dependence of ellipticity on redshift reported by him is very weak. Our result is affected by such factors as e.g. the adopted cluster centre, the way of smoothing the distribution and the applied method of ellipticity determination. These factors are being carefully studied now. The current result is consistent with Mulchaey \\& Zabludoff (1998) and Zabludoff \\& Mulchaey (1998) observation reporting the lack of recent significant dynamical evolution for poor groups and cD clusters. Moreover, the lack of the eccentricity evolution is in agreement with the latest hydrodynamic simulations of clusters (Floor et al. 2003a, 2003b).  \\begin{table} \\caption{The result of correlation analysis of the investigated samples} $$\\begin{tabular}{cccccc} \\hline \\hline \\emph{N} &\\emph{R} & \\emph{A} & \\emph{B} & \\emph{r} \\\\ \\hline \\hline 1 & 0.50 & 0.298 & 0.218 & 0.05 \\\\ 2 & 0.75 & 0.215 & 0.118 & 0.05 \\\\ 3 & 1.00 & 0.207 & -0.097 & 0.05 \\\\ 4 & 1.25 & 0.186 & -0.115 & 0.06 \\\\ 5 & 1.50 & 0.176 & -0.061 & 0.04 \\\\ \\hline \\hline 6 & Mean & 0.216 & 0.013 & 0.01 \\\\ \\hline \\hline \\end{tabular}$$ \\end{table}  \\bigskip \\bigskip \\bigskip"
    },
    "0404/astro-ph0404527_arXiv.txt": {
        "abstract": "{We investigate how galaxy-galaxy lensing measurements depend on the knowledge of redshifts for lens and source galaxies. Galaxy-galaxy lensing allows one to study dark matter halos of galaxies statistically using weak gravitational lensing. Redshift information is required to reliably distinguish foreground lens galaxies from background source galaxies and to convert the measured shear into constraints on the lens model. Without spectroscopy or multi-colour information, redshifts can be drawn from independently estimated probability distributions. The COMBO-17 survey provides redshifts for both lens and source galaxies. It thus offers the unique possibility to do this investigation with observational data. We find that it is of great importance to know the redshifts of individual lens galaxies in order to constrain the properties of their dark matter halos. Whether the redshifts are derived from $UBVRI$ or the larger number of filters available in COMBO-17 is not very important. In contrast, knowledge of individual source redshifts improves the measurements only very little over the use of statistical source redshift distributions as long as the source redshift distribution is known accurately. ",
        "introduction": "\\label{sect: introduction}  Weak gravitational lensing of galaxies by galaxies (galaxy-galaxy lensing) provides a unique tool to study the dark matter distribution in lens galaxies. A foreground galaxy distorts the images of background galaxies such that they are, on average, tangentially aligned with respect to the lens. This alignment can be measured statistically, see e.g.\\ \\citet{mellier1999,bartelmann2001} for reviews.  In the beginning, galaxy-galaxy lensing was measured from imaging data alone without any direct redshift information \\citep[e.g.][]{brainerd1996,dellantonio1996,fischer2000}. However, multi-colour or spectroscopic data are desirable for two reasons: \\begin{enumerate} \\item The strength of the gravitational shear and thus the image distortion depends on the angular diameter distances between observer, lens and source. In turn, redshift estimates for the lens and source galaxies are therefore needed to translate the measured distortion reliably into constraints on the lens galaxies. \\item The need for averaging over at least hundreds of lens galaxies complicates the interpretation of the results severely unless physically similar subsets of galaxies can be identified. As the luminous parts of galaxies cover a large range of properties, it is clear that galaxies also differ in the properties of their dark matter halos \\citep[e.g.][]{mckay2002,prada2003}. To learn something about different galaxy types one needs to be able to classify the potential lens galaxies at least by their luminosities and rest-frame colours, which requires redshift estimates. Note, that for this appliciation the redshift precision required is only a few percent, not the $10^{-4}$ of spectroscopic surveys. \\end{enumerate} Redshifts and/or classification for the lens galaxies have been used by several authors now \\citep[e.g.][]{smith2001,wilson2001,mckay2001,guzik2002,hoekstra2003}, while redshifts for the source galaxies have so far only been used by \\citet{hudson1998} who measured galaxy-galaxy lensing in the Hubble Deep Field North, and more recently by \\citet{sheldon2004} using the Sloan Digital Sky Survey (SDSS). All other studies used redshift probability distributions at a given apparent magnitude to assign redshifts to the source and -- where necessary -- to the lens galaxies. These probability distributions have been derived from galaxy redshift surveys like the Canada-France Redshift Survey \\citep{crampton1995} or the Caltech Faint Galaxy Redshift Survey \\citep{cohen2000}. However, for the deeper data sets, the measured redshift probability distributions had to be extrapolated to the fainter magnitudes of the source galaxies.  In this paper we address the question of how measurements of dark matter halos with galaxy-galaxy lensing are affected by incomplete redshift information. We use observational data with accurate photometric redshifts for both lens and source galaxies. We perform the measurements with (a) the full information for lens and source galaxies, (b) full information for the lens galaxies but drawing the redshifts of the sources from probability distributions, and (c) using probability distributions to estimate redshifts for lenses and sources. Although in case (a) we use the full redshift information, this measurement is not an optimal measurement of galaxy-galaxy lensing from COMBO-17. In this paper, the lens and source selection is always based on apparent magnitudes and angular separations which allows us to compare the results from the different cases. However, for retrieving the tightest and most meaningful constraints on the dark matter halos of galaxies, one should clearly use redshifts also for the lens and source selection. This will be done in a companion paper \\citep{kleinheinrich2004}.  In Sect.\\ \\ref{sect:data} we describe the data set followed by an overview of our method of measuring galaxy-galaxy lensing in Sect.\\ \\ref{sect:method}. Section \\ref{sect:z_estimates} explains how we estimate redshifts and luminosities. In Sect.\\ \\ref{sect:measurements} we present our measurements and results. We close with a summary in Sect.\\ \\ref{sect:outlook}. ",
        "conclusions": "\\label{sect:outlook}  We used the COMBO-17 survey to investigate the importance of individual redshift measurements on galaxy-galaxy lensing studies. We find that redshift information for the lens galaxies plays a crucial role in constraining their dark matter halos while for the source galaxies it is sufficient to have roughly correct redshift probability distributions available. We have also seen that redshifts obtained from just $UBVRI$ ($\\langle \\sigma_z/(1+z)\\rangle =0.03$ for $R=18-22$) instead of the full 17 filter set available in COMBO-17 ($\\langle \\sigma_z/(1+z)\\rangle =0.015$ for $R=18-22$) are sufficient to constrain the redshifts of the lenses and their rest-frame luminosities.  However, some benefits from the 17-filter classification remain. Most importantly, the 17-filter classification was used to derive the redshift probability distribution of the source population. In deep weak lensing studies, sources are typically fainter than the magnitude limit of current galaxy redshift surveys. Therefore, any redshift probability distribution can only be obtained from extrapolation to faint magnitudes. In Sect.\\ \\ref{sect:sources_sub4} we have shown that this extrapolation can introduce a bias into the measurement that is of the same order or even larger than the errors from large surveys like the RCS or SDSS. In principle, the redshift distribution of the sources could also be derived from just 5 filters. However, from 5 filters, the decrease in redshift accuracy is larger for the fainter sources than for the brighter lenses. Furthermore, fewer sources get reliable redshift estimates from 5 than from 17 filters. Therefore, the source redshift distribution derived from 5 filters will be less accurate than from 17 filters.  The gain from multi-colour or spectroscopic data for the lenses will even be larger than shown from the comparisons in this paper. Knowledge of individual redshifts will allow one to select lenses and sources not from less reliable magnitude cuts but from the redshifts themselves. Redshifts will also allow one to include fainter lenses or, when redshifts are available for the sources as well, brighter sources and will thus improve the statistics.  Any information on the nature of the lens galaxies will additionally allow one to investigate if and how the properties of the dark matter halos depend on galaxy classes which can be defined according to e.g.\\ luminosity, colour, stellar mass or environment. Such investigations are necessary to test and improve our understanding of galaxy formation and evolution where dark matter is supposed to play a major role."
    },
    "0404/astro-ph0404077_arXiv.txt": {
        "abstract": " ",
        "introduction": "Photodissociation regions (PDRs) form on the surfaces of molecular clouds whenever these clouds are struck by far-ultraviolet radiation from hot young stars.  These regions are characterized by the transition from hot, ionized gas to cold, molecular gas as the far-ultraviolet field is attenuated farther from the ionization front.  % The temperature profile of the PDR varies depending on the density and strength of the FUV field.  Derivation of this profile must take into account the local heating and cooling, the chemical equilibrium, and radiative coupling to other layers within the structure. The Orion Bar is a dense molecular structure at the southeast boundary of the Orion Nebula. Early theoretical models of the Orion Bar by Tielens and Hollenbach \\cite{tielens} (for a density of 2.3 $\\times$ 10$^4$ and UV field strength (G$_0$) of 10$^5$) predict temperatures of $\\sim$1000 K at A$_{\\mathrm{V}}$=0-2, dropping to less than 100 K by A$_{\\mathrm{V}}$=4.  Ground state rotational lines of H$_2$ are good temperature probes of moderately hot (200-1000 K) gas.  The low A-values of these lines result in low critical densities while ensuring that the lines are optically thin.  ISO observations of H$_2$ rotational lines in PDRs reveal large quantities of warm gas that are difficult to explain via current models\\cite{draine}, but the spatial resolution of ISO does not resolve the temperature structure of the warm gas.  We present and discuss high spatial resolution observations of H$_2$ rotational line emission from the Orion Bar. ",
        "conclusions": "PDR models predict that H$_2$ emission arises at a depth in the PDR with steep temperature gradients \\cite{tielens,draine}. We do not see evidence for such spatial gradients even though we have resolved the PDR.  The agreement of temperatures derived from the 0-0 S(1)/0-0 S(2) and 1-0 S(1)/0-0 S(1) line intensity ratio indicates that the emitting gas is isothermal. We re-interpret the 400-700 K temperatures found for galactic PDRs in the large ISO beam \\cite{wright} as emerging from isothermal gas, and not the beam-average of a temperature gradient. Future PDR modeling efforts must be able to explain the ubiquity of 600 K H$_2$ in PDRs."
    },
    "0404/astro-ph0404594_arXiv.txt": {
        "abstract": "The Reuven Ramaty High Energy Solar Spectroscopic Imager (RHESSI) has been observing gamma-ray lines from the Sun and the Galaxy since its launch in February 2002.  Here I summarize the status of RHESSI observations of solar lines (nuclear de-excitation, neutron capture, and positron annihilation), the lines of $^{26}$Al and $^{60}$Fe from the inner Galaxy, and the search for positron annihilation in novae. ",
        "introduction": "The \\it Reuven Ramaty High Energy Solar Spectroscopic Imager (RHESSI)\\rm, which was launched on February~5, 2002, is part of NASA's Small Explorer series of satellites \\citep{lin}.  \\it RHESSI \\rm was designed to perform imaging and spectroscopy of solar flares in the hard x-ray and gamma-ray range, with both spatial and spectral resolution far superior to previous missions.  It uses Rotating Modulation Collimators (RMCs) for solar imaging down to 2.3''\\citep{hur}.  Its nine coaxial germanium detectors cover the range from 3~keV to 17~MeV with extremely good energy resolution ($\\rm{E}/\\Delta\\rm{E}$ around 400 from 1--2~MeV) \\citep{sm1}.  The detectors are in the rear of the spacecraft; the spacecraft bus is wrapped around the central tube that contains the optical axis.  Each detector is segmented into a thin front segment, which records the copious hard x-ray photons from a solar flare, and a thick rear segment built to observe solar gamma-rays at higher energies that penetrate the front segment.  Only {\\it RHESSI's} coarsest grids are thick enough to effectively modulate gamma rays above 1~MeV.  Imaging is therefore limited to 35'' at the 2.223~MeV line from neutron capture in solar flares, and to 3' for the 4 and 6 MeV lines from carbon and oxygen de-excitation. Solar spectroscopy can be performed across the entire energy range of the detectors, subject to the limitations of photon statistics at gamma-ray energies.  Because these detectors have no shielding, and because the spacecraft is very light, the {\\it RHESSI} detectors can also see emission at gamma-ray energies from any direction in the sky.  For photons on the order of 1~MeV or higher, in fact, the instrument's effective area (about 29 cm$^2$ at 1809 keV) has little dependence on the direction of incidence, with the exception of a $\\sim$30\\%  depression at near-solar angles due to the effects of the grids, absorption in the front segments, and opacity of parts of the spacecraft bus. Since {\\it RHESSI} is in low-Earth orbit, the spectrum of emission from any given point in space can be derived by using the Earth as an occulter, subtracting blocked from unblocked data to determine background, if the source in question is the dominant source in the sky at the wavelength (and during the time period) under consideration. ",
        "conclusions": ""
    },
    "0404/astro-ph0404288_arXiv.txt": {
        "abstract": "Krsti\\'c has carried out the first quantum mechanical calculations near threshold for the charge transfer (CT) process ${\\rm H}^+ + {\\rm H}_2(X\\ ^1\\Sigma_g^+,\\ \\nu=0,\\ J=0) \\to {\\rm H}(1s) + {\\rm H}_2^+$.  These results are relevant for models of primordial galaxy and first star formation that require reliable atomic and molecular data for obtaining the early universe hydrogen chemistry.  Using the results of Krsti\\'c, we calculate the relevant CT rate coefficient for temperatures between 100 and 30,000~K.  We also present a simple fit which can be readily implemented into early universe chemical models.  Additionally, we explore how the range of previously published data for this reaction translates into uncertainties in the predicted gas temperature and H$_2$ relative abundance in a collapsing primordial gas cloud.  Our new data significantly reduce these cosmological uncertainties that are due to the uncertainties in the previously published CT rate coefficients. ",
        "introduction": "\\label{sec:Intro}  Molecular hydrogen is an important coolant during the epoch of primordial galaxy and first star formation.  Model calculations indicate that radiatively-induced cooling due to collisions of H$_2$, primarily with H, dominate the cooling of collapsing primordial gas clouds from temperatures beginning at $\\sim 10^4$~K and going down to $\\sim 5\\times 10^2$~K \\citep[e.g.,][]{Sas67a,Abe00a,Nak02a}.  The importance of H$_2$ in the early universe is supported by the recent {\\it Wilkinson Microwave Anisotropy Probe} ({\\it WMAP}) measurements which suggest (a) that reionization occured at high redshift and (b) a large Thompson scattering optical depth.  These results require star formation at this redshift to produce the ionizing radiation, which implies that H$_2$ cooling is important for early star formation \\citep{Spe03a,Hai03a}.  Given the pivotal role that H$_2$ is predicted to play, an accurate understanding of its formation and destruction in the early universe is crucial for understanding the formation of hierarchical structure (Haiman, Thoul, \\& Loeb 1996; Abel et al.\\ 1997; Abel, Bryan, \\& Norman 2002; Galli \\& Palla 1998; Flower 2002; Lepp, Stancil, \\& Dalgarno 2002; Nakamura \\& Umemura 2002; Ricotti, Gnedin, \\& Shull 2002).  The early universe chemistry of H$_2$ during the epoch of primordial galaxy and first star formation has been reviewed recently by a number of groups \\citep[e.g.,][]{Abe97a,Gal98a,Lep02a,Oh02a}.  These authors have found that the relative H$_2$ abundances are determined by only a handful of processes.  Of particular interest, \\citet{Oh02a} point out that the charge transfer (CT) process \\begin{equation} \\label{eq:CTH2} {\\rm H}^+ + {\\rm H}_2(X\\ ^1\\Sigma_g^+,\\ \\nu=0,\\ J=0)  \\to {\\rm H}(1s) + {\\rm H_2}^+ \\end{equation} is the dominant destruction mechanism of H$_2$ during the formation at this epoch of halos with viral temperatures $T_{vir}>10^4$~K.  This process dominates until a collapsing cloud of primordial gas begins internally to emit ionizing radiation or until the background extragalactic UV radiation field has reached significant levels.  Then the dominant destruction mechanism becomes photodissociation through the Lyman-Werver bands via the two-step Solomon process (Haiman, Rees, \\& Loeb 1997).  Because of the fundamental importance of reaction~(\\ref{eq:CTH2}), reliable rate coefficients for this process are needed for our understanding of the formation of structure in the early universe. However, as discussed below, the various rate coefficients adopted by the astrophysics community for this process differ from one another by orders of magnitude.  This translates into orders of magnitude uncertainties in the predicted relative abundance of H$_2$ at key epochs. Furthermore, this uncertainty directly affects the predicted properties of collapsing gas clouds during primordial galaxy and first star formation.  To help address the need for accurate rate coefficients for reaction~(\\ref{eq:CTH2}), \\citet{Krs02a} has recently carried out the first quantum mechanical calculations for this process at collision energies relevant for the cosmological formation of structure.  His results are appropriate for current models of early universe chemistry, which treat all molecules as being in their ground rovibrational state.  These models have not yet evolved to the point where they take into account the electronic and rovibrational distribution of molecules \\citep[e.g.,][]{Oh02a,Lep02a}.  Throughout the rest of this paper only ground electronic configurations are considered.  In \\S~\\ref{sec:Review}, we give a brief review of the previous experimental and theoretical work for reaction~(\\ref{eq:CTH2}).  A short description of the theoretical calculations by \\citet{Krs02a} is given in \\S~\\ref{sec:Theoretical}.  In \\S~\\ref{sec:Results} we present our new rate coefficient, provide a simple fitting formula that can readily be incorporated into early universe chemistry models, and comment on our results.  Lastly, in \\S~\\ref{sec:Cosmological} we briefly discuss some of the cosmological implications of our results. ",
        "conclusions": "\\label{sec:Results}  We have calculated the rate coefficient for reaction~(\\ref{eq:CTH2}) for temperatures of $10^2-3\\times 10^4$~K.  This covers the range of temperatures for which reaction~(\\ref{eq:CTH2}) is cosmologically important.  We present our results in Figure~\\ref{fig:h2hplusct}.  The calculated CT rate coefficient has been fitted using \\begin{equation} \\alpha(T) = \\exp(-a/T)\\sum_{i=0}^{7} b_i (\\ln T)^i. \\label{eq:fit} \\end{equation} The best fit values are listed in Table~\\ref{tab:fitparameters}.  The fit is accurate to better than 4\\% for $10^2 \\le T \\le 3\\times 10^4$~K.  Our calculated results represent the most sophisticated theoretical treatment of reaction~(\\ref{eq:CTH2}) at the present time.  Previous rate coefficient calculations, based on detailed balance or the TSH method, are unreliable for the reasons discussed in \\S~\\ref{sec:Review} and cannot be used to verify the validity of our results.  Similarly, \\citet{Lin95a} raise questions about the reliability of the laboratory measurements by \\citet{Hol71a}.  Hence, in our opinion, rate coefficients derived from the cross section data of either \\citet{Hol71a} or \\citet{Lin95a} cannot be used to verify the reliability of our calculated results.  In short, new laboratory measurements are needed to benchmark these new calculations."
    },
    "0404/astro-ph0404241_arXiv.txt": {
        "abstract": "Most direct source detections beyond $z\\sim7$ are likely to arise from wide-field narrowband surveys of \\Lya\\ emission in the J band. For this to be true, the \\Lya\\ emission must somehow escape from compact star-forming regions (CSR) presumably associated with massive haloes. Since the \\Lya\\ escape fraction is $\\la10$\\% from an emitting region of size $\\sim 1$ kpc, these objects will be difficult to find and hard to detect, requiring $\\sim 30$ -- 100 hours at each telescope pointing on 8 -- 10 m telescopes.  For CSR sources, existing large-format IR arrays are close to ideal in terms of their noise characteristics for conducting {\\it wide-field} narrowband surveys where pixel sizes are 0.1\\arcsec\\ or larger.  However, we stress that \\Lya\\ can also arise from external large-scale shocks (ELS) due to starburst winds, powered by CSRs, ploughing into gas actively accreting onto the dark halo.  The winds effectively carry energy from the dense, dusty environment of a starburst into lower density regions, where the escape probability for \\Lya\\ photons is greater.  ELS emission is expected to be considerably more clumpy ($\\la 100$ pc) than CSR emission.  For ELS sources, IR arrays will need $1-2$ orders of magnitude improvement in dark current in order to detect dispersed clumpy emission within the environments of massive haloes. These sources require an IR camera with small pixels ($\\sim 0.01$\\arcsec) and adaptive optics correction (\\eg\\ GSAOI on Gemini), and will therefore require a {\\it targetted} observation of a Dark Age source identified by a wide-field survey.  For either targetted or wide-field surveys, the deepest observations will be those where the pixel sampling is well matched to the size of the emitting regions.  We show that there are only 3 -- 4 J-band windows [$z = 7.72, (8.22), 8.79, 9.37$] suitable for observing the Dark Ages in \\Lya; we summarize their cosmological properties in Table 1. ",
        "introduction": "In recent years, we have come to recognize that the `ionization epoch' may lie just beyond our current observational horizon \\citep[\\cf][]{fwd00,fss03,kes04}.  Structure formation models in CDM cosmologies appear to show that this epoch took place at $z=7$ -- 12 \\citep{go97,hl98}. Recent wide-angle polarization measurements with WMAP \\citep{ksb03} suggest the ionization epoch could have been well under way by $z\\sim 17$. However, this may conflict with the inferred high column densities in $z\\sim 6$ quasars \\citep[\\eg][]{wl04}.  Since cosmic ionization requires only a tiny fraction of the primordial gas to be converted into stars or black holes \\citep{lb01}, it is possible to construct a wide range of scenarios \\citep{bcl99,bcl02,nu01,abn00,abn02,bl04,bl00, wl03,htl96,fwh01}.  Clearly, the first observations of the Dark Ages will have a dramatic impact on our understanding of this new frontier.  In conventional CDM simulations, CSR sources are associated with the cores of massive haloes. Since the distribution of dense peaks collapsing out of an evolving Gaussian density field is well constrained \\citep{me03,p99}, bright CSR sources are expected to be rare. These will be difficult to find even with well optimized wide-field surveys which exploit large pixels ($\\sim$0.1\\arcsec). However, the number of detectable sources above a given flux level at a fixed epoch is more uncertain. This requires a detailed understanding of how and when \\Lya\\ emission is produced, and how it manages to escape its immediate environment \\citep{n91,hs99}.  Even the most optimistic calculations \\citep{bw87} show that \\Lya\\ detections will be a major observational challenge on 8 -- 10m telescopes. The possible discovery of a lensed candidate galaxy at $z = 10$ by \\citet{psr04} demonstrates the power of lensing to extend our observational reach.  However, the total number of sources accessible this way is limited by the small total volume of the universe that is strongly magnified by foreground lenses.   During the Dark Ages, gas accretion onto protogalactic cores must have been well under way.  Galactic nuclei at the highest redshifts observed to date ($z\\sim5$) exhibit solar metallicities, and therefore appear to have undergone many cycles of star formation \\citep{hf99}. Most galaxy cores early in their evolution must have experienced starburst-driven winds. Early protogalactic winds will have carried large amounts of energy, with relatively low dust/metal content, away from the complex circumnuclear environment \\citep{dnh04}.     In order to understand physical processes associated with the first sources, we will need to resolve the \\Lya\\ structures. But, as we now show, detecting emission powered by outflows is a different experimental challenge from that posed by wide-field searches.   \\bigskip \\bigskip ",
        "conclusions": "It is clear that something quite extraordinary took place over a time span of a few hundred million years which led to the complete ionization of the intergalactic medium. The epoch of (re)ionization is likely to attract a great deal of attention over the coming decade.  We have shown that detecting sources beyond $z=7$ is going to be very challenging on 8 -- 10m telescopes. But since 30 -- 100m telescopes will not be in operation until the next decade, narrowband imaging is likely to dominate studies of the Dark Ages for years to come. We have outlined a strategy based on wide-field surveys to identify star formation in massive haloes, followed by detailed studies with an adaptive optics imager. The optimal technologies for both kinds of study will be available in the next few years.  It is often observed that major discoveries are made within five years of a new technology \\citep[\\eg][]{h81}.  What is less well known is that this statement normally applies to experiments which achieve the systematic noise limit of the apparatus.  \\citet{gb01} argue that the systematic limit possible with nod \\& shuffle spectroscopic observations argues for total integration times measured in months or even years.  Since large telescopes are general user facilities, it is relatively rare that hundreds of hours are devoted to a single target field. But in order to reach back to the Dark Ages, observing programmes of one or two weeks at a time will be essential."
    },
    "0404/astro-ph0404131_arXiv.txt": {
        "abstract": "{We present the first results obtained from our pilot ultra-deep near-IR survey with ISAAC/VLT, aimed at the detection of z$>7$ sources using lensing clusters as natural Gravitational Telescopes. Evolutionary synthesis models of PopIII and extremely metal-poor starbursts have been used to derive 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 us 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 magnification in the core of lensing clusters improves the search efficiency and subsequent spectroscopic follow up. The research efficiency will be significantly improved by the future near-IR multi-object facilities such as EMIR/GTC and KMOS/VLT-2, and later the JWST. }    \\addkeyword{Cosmology: Early Universe} \\addkeyword{Galaxies: High Redshift} \\addkeyword{Galaxies: Evolution} \\addkeyword{Infrared: Galaxies}   \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404307_arXiv.txt": {
        "abstract": "The contribution of Source Cosmic Rays (SCRs), accelerated and still confined in Supernova Remnants (SNRs), to the diffuse high energy \\gr emission above 1 GeV from the Galactic disk is studied. The \\grs produced by the SCRs have a much harder spectrum than those generated by the Galactic Cosmic Rays (GCRs). Extending a previous paper, a simple SNR population synthesis is considered and the Inverse Compton emission from the SCR electrons is evaluated in greater detail. Then the combined spectrum of \\gr emission from the Galactic Supernova Remnant population is calculated and this emission at low Galactic latitudes is compared with the diffuse \\gr emission observed by the EGRET and ground based instruments.  The average contribution of SCRs is comparable to the GCR contribution already at GeV energies, due to Supernovae of types II and Ib exploding into the wind bubbles of quite massive progenitor stars, and becomes dominant at \\gr energies above $100$~GeV. At TeV energies the dominant contribution is from SCRs in SNRs that expand into a uniform interstellar medium. In fact, the sum of hadronic and Inverse Compton \\grs would exceed the limits given by the existing experimental data, unless the confinement time $T_{SN}$, i.e. the time until which SNRs confine the main fraction of accelerated SCRs, is as small as $T_{SN}\\sim 10^4$~yr and the typical magnetic field strength in SNRs as large as 30~$\\mu$G.  Both situations are however possible as a result of field amplification through CR backreaction in the acceleration process. It is pointed out that accurate measurements of the low-latitude diffuse Galactic \\gr spectrum at TeV-energies can serve as a unique consistency test for CR origin from the Supernova Remnant population as a whole. ",
        "introduction": "The diffuse Galactic \\gr emission was measured with the Energetic Gamma Ray Experiment Telescope (EGRET) on the Compton Gamma Ray Observatory up to about 20 GeV. To first approximation it can be described by a suitable model for the diffuse interstellar gas, the Galactic Cosmic Ray (GCR) distribution, and the diffuse photon fields \\citep{hunteretal97a,hunteretal97b}. The resulting GCR distribution in the Milky Way essentially corresponds to that measured in situ in the local neighborhood of the Solar System. This is also our starting point.  However above about 1~GeV the observed average diffuse \\gr intensity in the inner Galaxy, $300^{\\circ}<l<60^{\\circ}$, $|b|\\leq 10^{\\circ}$, exceeds the model prediction significantly, and this excess appears to increase monotonically with \\gr energy. Spatially, the diffuse \\gr intensity is strongly increasing towards the Galactic midplane.  In a previous paper \\citet[hereafter referred to as Paper I]{bva00}, we have calculated the contribution of an unresolved distribution of CR sources to the diffuse \\gr flux in the disk, assuming the GCR distribution within its large confinement volume to be the same as that measured in situ. The CR sources contribute through the \\gr emission of the accelerated particles that are confined in their interior -- the Source Cosmic Rays (SCRs). It was assumed that Supernova Remnants (SNRs) are the dominant sources of the GCRs. On this premise it was found that the SCRs give an important contribution to the diffuse high-energy $\\gamma$-ray emission at low Galactic latitudes. Since the CR energy spectrum inside SNRs is much harder than on average in the Galaxy --- the average spectrum being softened by rigidity-dependent escape from the Galaxy in the diffusion/convection region above the disk --- the relative SNR contribution increases with energy and was found to become in fact dominant at \\gr energies $\\epsilon_{\\gamma} \\gsim 100$~GeV. This led to a substantial increase of the \"diffuse\" TeV $\\gamma$-ray emission from the Galactic disk so as to constitute a significant observational background. This hard-spectrum background must also be taken into account in the search for spatially extended Galactic CR sources above the GeV region.  Extending Paper I we take into account here that a significant fraction of the SNR population, the type II and Ib Supernovae (SNe) with progenitor masses in excess of $15M_{\\odot}$, expand into a nonuniform circumstellar medium strongly modified by the wind from the progenitor star. In such SNRs the majority of CRs should be produced in a thin dense shell of swept-up ISM. As a consequence of the high density of this shell the expected $\\pi^0$-decay $\\gamma$-ray emission with energies $\\epsilon_{\\gamma}\\lsim 100$~GeV is shown to be an order of magnitude larger than for an identical explosion into a uniform ISM. This gives a natural explanation for the observed diffuse $\\gamma$-ray excess for energies $\\gsim 1$~GeV. We also consider in detail the contribution of the high-energy SCR electrons to the high energy \\gr background due to inverse Compton (IC) scattering on the background photon field. Since the energy spectrum of the SCR electrons is strongly influenced by synchrotron losses, the IC \\gr emission depends on the effective magnetic field strength in the SNR interior. This field may in turn be amplified relative to the external field by the acceleration process itself \\citep{belll}. We take such a field amplification into account as an average property of SNRs relative to the contribution from the GCRs. As demonstrated below, this results in an order of magnitude increase of the Galactic \\gr emission at TeV-energies. In a preliminary form these results have been presented in \\citet{bv03}.  We therefore conclude that a detailed measurement of the low-latitude diffuse Galactic \\gr spectrum within the energy interval from 1 to $10^4$~GeV will constitute an important indirect test for the GCR origin from the {\\it population} of Galactic SNRs as a whole. ",
        "conclusions": "Our above consideration is based upon a picture in which the population of SNRs, considered here as the main GCR sources, are distributed across the Galactic disk similar to the Interstellar gas. In this case the calculated ratio $R(\\epsilon_{\\gamma})$ of the SCR to the GCR contribution does not depend upon the observing direction in the disk. Since the number of these sources $N_{SN}=\\nu_{SN}T_{SN}$ is very limited, their mean number within the field of view $\\sim {\\mathrm few}$~ degrees of a stereoscopic system of imaging atmospheric Cherenkov telescopes for \\gr energies $\\gsim 100$~GeV is so low, that one should expect large fluctuations of the actual value of $R(\\epsilon_{\\gamma})$, especially for lines-of-sight other than those directed towards the inner Galaxy.  The actual SNR distribution within the Galactic disk most probably is not uniform. According to \\citet{case} the SN explosion rate as a function of Galactocentric radius $r$ has a peak at $r\\approx 5$~kpc and drops exponentially with a scale length of $\\sim 7$~kpc. Within $5<r< 20$~kpc this agrees fairly well with the radial distributions of {\\it Supernovae} in a sample of 36 {\\it external galaxies} of Hubble type Sb--Sbc, statistically considered equivalent to our Galaxy \\citep{drag}.  From our results, for \\gr energies larger than 10 GeV, the ratio $R$, corrected for the nonuniform SNR distribution, is then given by the ratio $R_c=\\eta(r) R$, where the parameter $\\eta(r)$ represents the ratio of the actual SNR number along the given line-of-sight to that SNR number which would corresponds to their uniform distribution. For lower \\gr energies, $100< \\epsilon_{\\gamma}< 10^4$~MeV, where the truly diffuse emission from the GCRs dominates, it is known that the \\gr emissivity gradient is significantly shallower than the gradient of $\\eta(r)$, as shown by the EGRET instrument on board of the CGRO satellite \\citep{strm} in basic agreement with the results from the COS-B satellite \\citep{stretal}. This much shallower, indeed truly diffuse Galactocentric \\gr gradient can be understood as a nonlinear propagation effect from the disk into the Galactic Wind where the latter is driven by the GCRs themselves \\citep{breit}. In contrast, the radial galactocentric variation of the direct radiation from the sources is not smoothed by any propagation effect. Therefore its amplitude should spatially vary like the distribution of the sources, integrated along the line-of-sight, and this variation should become directly measurable in the $> 100$~GeV range. Given several such line-of-sight integrals, one can infer the radial galactocentric \\gr emissivity gradient. The excellent angular resolution of atmospheric Cherenkov telescopes and the high sensitivity of instruments like the CANGAROO III and H.E.S.S. arrays in the Southern Hemisphere, which at TeV energies have an order of magnitude higher sensitivity compared with previous generation instruments like HEGRA \\citep{aha01}, should allow this measurement in particular in the inner Galaxy.  The assumed average value $B=30~\\mu$G of the magnetic field inside SNRs is considerably smaller than the value $120~\\mu$G inferred for e.g. SN~1006 in the very early Sedov phase \\citep{bkv02,bkv03}. However the amplification must be a monotonically increasing function of the shock strength. In particular, \\citet{belll} argue for a dependence $B \\propto V_s$. If so, then the effective amplified magnetic field in SNRs varies from $B\\sim 100$~$\\mu$G at the beginning of the Sedov phase to the typical interstellar value $B\\approx 10$~$\\mu$G at the end of its active period $t=T_{SN}$. Since the value of the magnetic field in SNRs has a dominant influence on the IC \\gr spectrum, and since all evolutionary phases of a SNR contribute to this spectrum in a roughly equal way, the adopted value $B=30$~$\\mu$G appears to be quite realistic, averaging over the entire active period.  One should also note that the SNR magnetic field influences the maximum energy of SCRs. Since the most energetic SCRs are created at the beginning of the Sedov phase, their actual maximum energy $\\epsilon_{max}\\propto B(t_0)$ is expected to be larger than considered here, because the amplified magnetic field $B(t_0)$ in this phase is by a factor of several larger than on average. Compared with the Fig.\\ref{fig2} this effect will produce more a extended and smooth high energy tail of the \\gr spectrum up to the energy $\\epsilon_{\\gamma}\\approx 10^5$~GeV.  The actual process of SN shock interaction with the dense shell could be much more complicated compared with the ideal picture considered here. If due to some physical factors (inhomogeneities in the surrounding ISM, instabilities ...) the shell is {\\it always} strongly distorted, then one would expect that the SN shock will penetrate through the shell more rapidly. This will lead to a decrease of the propagation time $t_f$ and/or to the decrease of the effective gas number density $N_g^{SCR}$, that in turn will decrease the amount of \\grs produced in such types of SNRs. In other words, the calculated \\gr spectrum at $\\epsilon_{\\gamma}\\lsim 100$~GeV (see Fig.\\ref{fig1}, \\ref{fig2}), which is mainly due to the wind SNe, has to be considered as an upper limit, which corresponds to the case of stable unmodified bubbles."
    },
    "0404/astro-ph0404313_arXiv.txt": {
        "abstract": "The structure of the universe is believed to have formed by clustering hierarchically from small to large scales. Much of this evolution occurs very slowly but at a few special times more, rapid, violent activity may occur as major subunits collide at high velocities.  The Abell 2125 complex (z=0.247) appears to be undergoing such an event as shown by modeling of the optical velocity field and by the detection with the VLA of an unusually large number of associated radio active galaxies. We present an 80 ksec {\\sl Chandra} imaging of Abell 2125, together with extensive complementary multi-wavelength data. We show direct evidence for galaxy transformation and destruction during the cluster formation. The {\\sl Chandra} data unambiguously separate the X-ray contributions from discrete sources and large-scale diffuse gas in the Abell 2125 complex, which consists of various merging clusters/groups of galaxies and low-surface brightness emission. This enables us to study processes affecting galaxy evolution during this special time from scales of Mpc down to a few kpc. The overall level of activity plus the special time for the cluster-cluster merger suggests that an important phase of galaxy evolution can take place during such events. ",
        "introduction": "We have identified a large-scale hierarchical complex (Fig.\\ 1  left panel) that is well-suited for investigating the structure formation process and the galaxy/environment interactions. Revealed in a survey of 10 Butcher \\& Oemler clusters observed with the \\rosat\\ PSPC, this complex contains various X-ray-emitting features, which are associated with concentrations of optical and radio galaxies (Fig.\\ 1; Wang, Connolly, \\& Brunner 1997; Wang, Owen, \\& Ledlow 2004). The overall optical galaxy concentration of the region was classified previously as a cluster Abell 2125 (richness 4). In addition to its large blue galaxy fraction ($\\sim 20\\%$), Abell 2125 also contains a unusually high number of radio sources, a factor of a few higher than those in more-or-less relaxed clusters at about the same redshifts. The \\rosat\\ image and follow-up optical observations have further shown that the complex contains three well-defined X-ray bright clusters (Wang et al. 1997). In addition, substantial amounts of unresolved low-surface brightness X-ray emission (LSBXE) are also present. The overall size of the complex is at least $\\sim 5$~Mpc (assuming the now standard $\\Lambda$CDM cosmology). The complex thus represents an X-ray-bright hierarchical filamentary superstructure, as predicted by numerical simulations of the structure formation (e.g.,  Cen \\&  Ostriker 1996).   \\begin{figure} \\centerline{ } \\caption{\\protect\\footnotesize \\rosat\\  PSPC X-ray image of the Abell 2125 complex in the 0.5-2~keV band (left panel; Wang et al. 1997). Point-like X-ray sources detected in the image have been excised.  Spectroscopically confirmed, radio detected members of the Abell 2125 complex (Owen et al. 2004a) are marked by {\\sl pluses}. The box outlines the field covered by our \\chandra\\ ACIS-I observation shown in the right panel, where the intensity image is in the 0.5-2 keV band. {\\sl Squares} mark the positions of complex member galaxies.} \\label{fig1} \\end{figure}  We have obtained an 80 ksec \\chandra\\ observation that covers part of the Abell 2125 complex (Fig. 1, right panel) to characterize its detailed X-ray properties and to study the interplay between galaxies with their environments. The high spatial resolution of \\chandra\\ enables us to detect point-like sources down to a limit of $\\sim 1 \\times 10^{-15}$ ${\\rm~erg~cm^{-2}~s^{-1}}$ in the 0.5-8 keV band and to examine diffuse X-ray structures down to a scale of $\\sim 3.8$~kpc ($\\sim 1^{\\prime\\prime}$). Here we summarize results from the analysis of the data and from the comparison with extensive multi-wavelength observations of the complex. ",
        "conclusions": ""
    },
    "0404/astro-ph0404125_arXiv.txt": {
        "abstract": "{We present an $RI$ photometric survey covering an area of 430 arcmin$^2$ around the multiple star $\\sigma$ Orionis. The observations were conducted with the 0.8 m IAC-80 Telescope at the Teide Observatory. The survey limiting $R$ and $I$ magnitudes are 22.5 and 21, and completeness magnitudes 21 and 20, respectively. We have selected 53 candidates from the $I$ vs.\\ $R$--$I$ colour-magnitude diagram ($I$ = 14--20) that follow the previously known photometric sequence of the cluster. Adopting an age of 2--4 Myr for the cluster, we find that these objects span a mass range from 0.35 $M_{\\odot}$ to 0.015 $M_{\\odot}$. We have performed $J$-band photometry of 52 candidates and $K_{\\rm s}$ photometry for 12 of them, with the result that 50 follow the expected infrared sequence for the cluster, thus confirming with great confidence that the majority of the candidates are bona fide members. $JHK_{\\rm s}$ photometry from the Two Micron All Sky Survey (2MASS) is available for 50 of the candidates and are in good agreement with our data. Out of 48 candidates, which have photometric accuracies better than 0.1 mag in all bands, only three appear to show near-infrared excesses. ",
        "introduction": "The Orion complex, because of its relative closeness and youth, is one of the most suitable sites for understanding low-mass star formation processes. Recently, {\\itshape ROSAT} pointed observations within this complex has led to the discovery of the very young stellar cluster $\\sigma$ Orionis, around the multiple star of the same name (Walter et al. \\cite{walter97}; Wolk \\& Walter \\cite{wolk98}). Follow-up photometric and spectroscopic studies have revealed a sequence of objects in the colour--magnitude diagram that extends well below the substellar limit (B\\'ejar et al. \\cite{bejar99}, hereafter BZOR; B\\'ejar et al. \\cite{bejar01}, hereafter BMZO). Studies of the depletion of lithium in the atmosphere of K6--M8.5 type low-mass members of the cluster impose an upper limit of 8 Myr on the age and suggest a most likely cluster age in the interval 2--4 Myr (Zapatero Osorio et al. \\cite{osorio02}). {\\itshape Hipparcos} provides a distance modulus of $m-M = 7.7 \\pm 0.7$ (Perryman et al. \\cite{perryman97}) for the central star, $\\sigma$ Orionis. This star is affected by a low extinction value of $E(B-V) = 0.05$ (Lee \\cite{lee68}), and the associated cluster also seems to exhibit very little reddening (see BMZO and Oliveira et al. \\cite{oliveira02}). These combined characteristics of youth, proximity and low extinction make $\\sigma$ Orionis one of the most interesting clusters for studying young substellar objects and the substellar mass function.    In this paper we present an extension of the $RI$ survey conducted by  BZOR, with the aim of detecting new low-mass star and brown dwarf candidates in the $\\sigma$ Orionis cluster. We also present near-infrared data for these objects. Details of the observations are indicated in Section 2. In sections 3.1 and 3.2 we explain the selection of the candidates and discuss their membership and the presence of near-infrared excesses. Our conclusions are given in section 4. ",
        "conclusions": "In this paper we present an $RI$ survey of the $\\sigma$ Orionis cluster, covering an area of 430 arcmin$^2$. We have selected 53 candidates in the magnitude range $14<I<20$ that follow the optical photometric sequence of the cluster corresponding to masses in the range 0.35--0.015 $M_{\\odot}$. All but two of the candidates follow the cluster sequence in the infrared and are considered to be likely cluster members: around 31 are stars and 20 are brown dwarfs. The available spectroscopy for some of these objects confirms them as bona fide members. Using near-infrared photometry from 2MASS, we conclude that only three of the 48 (6 \\%) candidates show near-infrared excesses possibly related to the presence of discs."
    },
    "0404/astro-ph0404549_arXiv.txt": {
        "abstract": "A new set of software applications and libraries for use in the archival and analysis of pulsar astronomical data is introduced. Known collectively as the PSRCHIVE scheme, the code was developed in parallel with a new data storage format called PSRFITS, which is based on the Flexible Image Transport System (FITS). Both of these projects utilise a modular, object-oriented design philosophy.  PSRCHIVE is an open source development environment that incorporates an extensive range of C++ object classes and pre-built command line and graphical utilities.  These deal transparently and simultaneously with multiple data storage formats, thereby enhancing data portability and facilitating the adoption of the PSRFITS file format.  Here, data are stored in a series of modular header-data units that provide flexibility and scope for future expansion.  As it is based on FITS, various standard libraries and applications may be used for data input/output and visualisation.  Both PSRCHIVE and PSRFITS are made publicly available to the academic community in the hope that this will promote their widespread use and acceptance. ",
        "introduction": "\\subsection{Collaborative scientific software development}  Modern, highly specialised experimental systems often require extensive original software development. This is true for all tasks from direct hardware control through to data reduction. Whilst individual research groups often approach such software development from an isolated perspective, the proliferation of digital hardware and wide-area networking makes global cooperative software development far more attractive, provided suitable common ground exists. Both cooperative software development and the adoption of standard packages provide a number of distinct advantages to the research community. For instance, as less effort is wasted unnecessarily duplicating the work of others, cooperative development can lead to more efficient allocation of resources. In addition, supporting the requirements of a larger user community promotes the development of basic, general purpose routines that may be used in a wider variety of situations. These influences result in more modular and extensible software.\\\\  However, it should be noted that a greater level of care and cooperation is required in collaborative software development, especially the open source approach advocated in this paper.  For example, in contrast to most commercial software, ``black box'' modularity is undesirable in scientific analysis, especially when the application of certain algorithms requires experienced judgment. Open, well-documented code provides researchers with an accurate understanding of third-party analytical tools.  Therefore, contributing developers must be willing to put their code in the public domain, making it freely available for non-commercial use by any other academic organisation.  Although this facilitates the exchange of ideas, it also raises the issue of potential loss of intellectual property, which might discourage some authors.\\\\  It is also the case that collaborative development tends to become de-centralised, especially when multiple developers have the ability to commit fundamental changes to the code.  Effective communication between the core developers becomes essential to the smooth running of the project, necessitating greater attention to version control, the maintenance of stable releases, and the development of extensive and concise documentation.  Also, when a wider user community is affected by modifications to the software, exhaustive methods must be employed to ensure the validity of changes and the integrity of the system as a whole.  Although each of these issues tend to increase the workload of the collaborative developer, a much larger body of users will benefit from the effort.  \\subsection{Software development in the pulsar community}  The global pulsar community is ideally suited to adopt a collaborative approach to software development.  It consists of a relatively small number of locally centralised groups that deal with different telescopes and instruments, leading to several parallel but incompatible software development paths.  As each path tends to be built around a highly specific data storage format, cross examination of data and algorithms is problematic.  In addition, because such software is generally designed for a limited purpose, it is often difficult to extend its functionality without introducing obfuscated code.  This is especially true when the program develops in an experimental fashion, as is often the case with scientific applications.  In order to avoid future inflexibility, sufficient time and care must be invested during the planning stage, calling on input from both experienced software developers and pulsar astronomers.   \\subsection{Radio pulsar data}  Radio pulsars are broadband point sources of highly polarised emission that exhibit rapid pulsations with a characteristic period anywhere between one millisecond and ten seconds.  They are thought to be rotating neutron stars with a strong dipolar magnetic field whose axis is not aligned with the rotation axis of the star (Gold 1968). Intense beams of emission originate at the magnetic poles, which sweep across the sky with each rotation of the star and produce the pulsed radio signal observed at Earth.\\\\  The characteristic \\textit{signature} of any radio pulsar is its integrated polarimetric pulse profile, given by the observed Stokes parameters averaged (folded) as a function of pulse longitude over several thousand individual pulses (Helfand, Manchester, \\& Taylor 1975). Under the influence of electrons in the interstellar medium (Taylor \\& Cordes 1993), this pulsed signal is broadened by dispersive frequency smearing, which must be corrected in order to infer the shape of the characteristic profile at the source.  This is normally done by dividing the observed bandwidth into narrow frequency channels, which are appropriately delayed relative to each other before summing the detected flux densities in each channel.  However, as the dispersion measure may vary with time or may not be known with sufficient accuracy at the time of the observation, it is often necessary to store the individual pulse profiles observed in each frequency channel.\\\\  Additionally, it is possible to create a mean pulse profile only if a suitably accurate model of the pulsar's spin frequency and phase is available.  The apparent pulse period is affected by a number of phenomena, including the spin-down, timing noise, and/or glitches intrinsic to the pulsar, variations in the interstellar dispersion, and Doppler effects introduced by the relative motions of the Earth and pulsar.  Inaccuracies in the model that describes these effects introduce phase errors that accumulate with time and cause the integrated profile to become smeared.  Therefore, it is often beneficial to store multiple, shorter integrations of the mean pulse profile instead of a single, long integration.  Furthermore, when a pulsar is bright enough, a great deal of additional information about the characteristics of the pulsar emission can be obtained by recording and analysing each individual pulse.  Therefore, a useful pulsar data format must be able to represent pulse profiles observed over multiple epochs of arbitrary length.\\\\  In summary, pulsar observations generally consist of a four-dimensional array of data indexed by polarization component, pulse phase, frequency, and epoch.  Software support for sensible groupings in other dimensions, such as orbital phase, is also highly desirable.  In addition, data from a number of telescopes can be combined to increase sensitivity and contribute to the eventual detection of new phenomena, such as the cosmic background of stochastic gravitational radiation (e.g. Hellings \\& Downs 1983; Stinebring et al. 1990).  Therefore, the data storage format should have a flexible structure that provides efficient access to key parameters, removed from any considerations of individual instruments or signal processing schemes.  \\subsection{Processing radio pulsar data} \\label{sec:processing}  Pulsars are observed for a variety of reasons, from studying the nature of their structure and emission mechanism (Dodson, McCulloch, \\& Lewis 2002) to utilising them as highly stable clocks and astrophysical probes (Taylor \\& Weisberg 1982). Consequently, the same pulsar observation can be used in a number of different contexts, one focusing on the variation of polarization with frequency, another measuring general relativistic effects on pulse times of arrival, etc. Nevertheless, our experience has shown that there exist many common tasks associated with pulsar data analysis that can be standardised within a modern open source development environment.\\\\  As a demonstration of the types of operations performed on pulsar data, consider the specific example of pulse time of arrival calculation.  In order to increase the signal-to-noise ratio (S/N) of each observation, data are often integrated (scrunched) by several factors in one or more of the available dimensions.  Each resultant profile is then cross-correlated with a high S/N standard profile known as a template, yielding an estimate of the longitudinal offset between the two.  This offset is added to the reference epoch associated with a fiducial point in the observed pulse profile, yielding an arrival time in the reference frame of the observatory, which is later converted into a barycentric arrival time using a Solar System ephemeris.  This data reduction operation involves a number of typical tasks, including loading the arrays of numbers that represent the folded profiles and computing sums, products, rotations, weighted averages, and correlations of these arrays; sometimes in the Fourier domain.  Most of these various operations must be performed in a manner consistent with the observational parameters, taking into account dispersive delays, observation time stamps and relative weightings of different frequency channels, for example.  At each step, the software must also ensure that all parameters are updated accordingly.  \\subsection{Scope and design of PSRCHIVE and PSRFITS}  It should be noted that the pulsar data under consideration represents a point near the end of the typical pulsar data reduction chain. The software presented in this paper is not intended for the direct handling of radio data such as that recorded by baseband systems, or for the purposes of performing computationally expensive offline searching, although some support for the storage of such data is provided in PSRFITS. The code is also not designed to perform any phase-coherent dispersion removal or formation of filter-bank data; these techniques are treated as separate computational tasks. Code for such data reduction is also available from the repository at the Swinburne Centre for Astrophysics and Supercomputing under the umbrella name of BASEBAND DSP\\footnote{http://astronomy.swin.edu.au/pulsar/software/libraries/dsp}, a general library for use in digital signal processing.\\\\  The PSRCHIVE and PSRFITS schemes were designed from the beginning to form an object-oriented framework into which existing algorithms and data structures could be transplanted.  By introducing layers of abstraction between the various levels of responsibility, the design remains both flexible and extensible.  For example, different telescopes and instruments require the storage of different types of information, including configuration parameters, observatory and instrumental status information, and other site-specific data. Because there is no way of knowing exactly what future systems might include, both PSRCHIVE and PSRFITS implement a generalised scheme for incorporating arbitrarily complex data extensions, as described in Sections~\\ref{sec:abstraction} and~\\ref{sec:psrfits_defn}.\\\\  In addition, a basic framework of crucial parameters common to all pulsar observations and a wide variety of fundamental data reduction algorithms, such as those described in Section~\\ref{sec:processing}, have been implemented.  Each of these basic data structures and reduction operations may be used in the composition of more complex scientific processing algorithms.  By virtue of continued development amongst the authors, the PSRCHIVE library includes an extensive array of such high-level algorithms for use in the calibration, visualisation, and analysis of pulsar data; these can be used immediately on any of the supported file formats.\\\\  PSRCHIVE and PSRFITS were developed in parallel and are presented in the hope that they will promote increased data portability.  The PSRFITS file format also serves as an example of how to incorporate other, pre-existing file formats into the new scheme, as described in Section~\\ref{sec:plugin}.  After two years of development, the code is now ready for formal release to the wider pulsar community.  In the following sections, we describe the implementation of the new schemes and outline the specific advantages that they offer. ",
        "conclusions": "The task of organising astronomical data into a logical format lends itself surprisingly well to the object-oriented programming paradigm. The combination of PSRCHIVE and PSRFITS provides a powerful, ready-to-use pulsar data archival and reduction system that can be rapidly adapted to new instruments.  We hope that the ready availability of an open source data reduction framework will facilitate large scale collaborative projects, such as an extended pulsar timing array (Foster and Backer 1990).  Therefore, we encourage both scientists and engineers involved with pulsar data acquisition and reduction to consider taking advantage of these packages."
    },
    "0404/astro-ph0404580_arXiv.txt": {
        "abstract": "{The Earth is a planet with a dipolar magnetic field which is agitated by a magnetized plasma wind streaming from the Sun. The magnetic field shields the Earth's surface from penetrating high energy solar wind particles, as well as interstellar cosmic rays. The magnetic dipole has reversed sign some hundreds of times over the last 400 million years. These  polarity reversals correspond to drastic breakdowns of the dynamo action. The question arises what the consequences for the Earth's atmosphere, climate, and, in particular, biosphere are. It is shown by kinematic estimates and three-dimensional plasma-neutral gas simulations that the solar wind can induce very fast a magnetic field in the previously completely unmagnetized Earth's ionosphere that is strong enough to protect Earth from cosmic radiations comparable to the case of an intact magnetic dynamo. ",
        "introduction": "Paleomagnetic records show that the magnetism of Earth has reversed itself hundreds of times over the last 400 million years (Valet \\& Meynardier 1993; Juarez et al. 1998; Gee et al. 2000; Selkin \\& Tauxe 2000; Valet 2003). In fact, geomagnetic polarity reversals represent the most dynamic feature of the Earth's magnetic field.  The polarity reversals do not occur instantaneously. Rather, transition periods, that span some thousand years and are characterized by unstable varying magnetic fields with no clear shape, lay between the stable dipole field states. During the transition periods the magnetic field strength can drop well below 10$\\%$  of the average value (Juarez et al. 1998; Guyodo \\& Valet 1999; Selkin \\& Tauxe 2000) which signifies a drastic breakdown of the Earth's dynamo. In the present time the magnetic south pole has wandered over more than 1100 km during the last 200 years. The strength of the Earth's field has decreased by $5\\%$ per century. This decrease is by far the fastest that has been verified since the last total field reversal, the end of the so-called Matuyama chron, 730,000 years ago. Also, by statistical estimates the Earth's dynamo is overdue for a reversal. Thus, we have to expect a transition period characterized by a very small Earth's magnetic field in the near future. Since the Earth's dipole field provides us with a shield against cosmic rays and solar high-energy radiation one may wonder about the consequences for life on Earth. Also, having in mind that Mars lost the atmosphere almost completely after the final breakdown of the magnetic field \\cite{luh2}, one may speculate that stripping by the solar wind could alter the Earth's atmosphere. Severe climate changes could result. Interesting enough, during the last Brunhes-Matuyama reversal, no major changes in plant and animal life have been detected. This may be partly due to the fact that the atmospheric layers block some fraction of the cosmic radiation by scattering. On the other hand, the cosmogenic radionuclide production varies at least over the last 200,000 years as function of short-term variations of the magnetic field \\cite{frank00}.  In this contribution. we consider the interaction of the solar wind with a completely unmagnetized Earth. When the solar wind encounters unmagnetized objects, such as Venus \\cite{luh95} and comets \\cite{konz}, magnetic barriers and ionopauses develop. Although the interaction of a fully ionized and a weakly ionized gas is very complex, an important characteristic can be identified - the generation of magnetic fields caused by relative plasma- neutral gas shear flows. It has been shown \\cite{hub93} that this process operates in the Venus' ionosphere and is responsible for the non-dipole magnetic field measured there. The same process has been studied in detail for the generation of seed magnetic fields in emerging galaxies (Wiechen et al. 1998; Birk et al. 2002) and in circumstellar disks \\cite{birk03}. A kinematic estimate indicates that relatively strong magnetic fields are generated in the Earth's ionosphere. So far, the only way to study the dynamical non-linear interaction of the magnetized fully ionized solar wind plasma and the partially ionized Earth's ionosphere are three-dimensional plasma-neutral gas simulations. Our numerical studies show the draping of the magnetic field of the solar wind and the self-generation of a relatively strong magnetic field in the Earth's ionosphere. ",
        "conclusions": "We studied the interaction of the magnetized fully ionized solar wind plasma with the unmagnetized partially ionized Earth's ionosphere. When the solar wind hits the Earth the magnetic field lines carried along with it are draped around the planet. What is more important, the relative motion between the solar wind plasma and the ionosphere results in the self-generation of magnetic fields in the ionospheric layer. The strengths of these fields depend on the shear length of the relative flows, which, in contrast to the other relevant physical parameters, is not well known. For a reasonable shear length of $L=10$ km the maximum strength of the newly generated magnetic field is comparable to the one of the present dipole field. Consequently, even in the case of a complete breakdown of the Earth's dynamo, the biosphere is still shielded against cosmic rays, in particular coming from the sun, by the magnetic field induced by the solar wind.   \\begin{figure} \\centering{ \\includegraphics[width=9cm,keepaspectratio]{Gb091_f3.eps}} \\caption{The strength of the induced magnetic field in the Earth's ionosphere at t=500 s.} \\label{fig3} \\end{figure}"
    },
    "0404/astro-ph0404063_arXiv.txt": {
        "abstract": "{ The objective of this series of two papers is to investigate small-scale molecular structures in the Magellanic Clouds (hereafter MCs). We report on the {\\small FUSE} detections of the HD and CO molecules {\\bf on the lines of sight towards three Large Magellanic stars}: Sk\\,$-$67D05, Sk\\,$-$68D135, and Sk\\,$-$69D246. HD is also detected for the first time {\\bf on the lines of sight towards two Small Magellanic Cloud stars}: AV\\,95 and Sk\\,159. While the HD and CO abundances are expected to be lower in the Large Magellanic Cloud where molecular fractions are a third of the Galactic value and where the photodissociation flux is up to thousands times larger, we report an average HD/H$_2$ ratio of 1.4$\\pm$0.5 ppm and CO/H$_2$ ratio ranging from 0.8 to 2.7 ppm similar to the Galactic ones. We tentatively identify a deuterium reservoir (hereafter D--reservoir) towards the Small Magellanic Cloud, along the light path to AV\\,95. We derive a D/H ratio ranging from 1. 10$^{-6}$ to 1.1 10$^{-5}$. Combining {\\small FUSE} and {\\small HST}/{\\small STIS} data we also analyzed the H$_2$, Cl\\,{\\small I}, Cl\\,{\\small II}, Fe\\,{\\small II}, S\\,{\\small II}, C\\,{\\small I}, C\\,{\\small I}*, and C\\,{\\small I}** content, when available. High resolution {\\small VLT} observations of Na\\,{\\small I}, Ca\\,{\\small II}, and Ca\\,{\\small I} were obtained in support of the lower resolution {\\small FUSE} and {\\small STIS} data for three targets in order to unravel the sightline velocity structures. These observations form the only such set of detections in the Magellanic Clouds to date and allow us to investigate in {\\bf detail} some of the physical properties of the intervening molecular gas. Our observation of the HD and CO molecules in the Magellanic Clouds is an argument for dense ({\\bf $n_{\\rm H}$} $>$ 100 cm$^{-3}$) components. Furthermore, we demonstrate that these components are probably extremely small molecular clumps (possibly as small as 10$^2$ pc) or filaments similar to the tiny--scale atomic structures (TSAS) recently observed in the halo of our Galaxy by {\\bf Richter et al. (2003)}. Interestingly enough, for these five sightlines, we also detected molecular hydrogen originating in the Galactic disk. From these observations, we conclude that tiny--scale molecular filamentary structures are present in the disk of the Galaxy as well. ",
        "introduction": "For many years, the study of molecular clouds through the CO emission lines has greatly improved our knowledge of the physical properties within the Magellanic Clouds themselves and helped constrain the theoretical models of molecular clouds formation and chemistry in different Galactic environments (Cohen et al. 1988). Over the past decade, in particular, the Swedish--ESO--Submillimeter--Telescope (SEST) has been used to map the radio--emission of molecular complexes within the LMC at an $8.\\arcsec0$  angular resolution which is an order of magnitude higher compared to previous surveys. Such resolution corresponds to a linear extent of about 10 pc at the LMC distance of 50 kpc. The latter survey led to the clear detection of small CO clumps with sizes possibly smaller than 10 pc (Lequeux et al. 1994; Garay et al. 2002), typical densities of a few thousand molecules per cubic centimeter and virial masses around 10$^{4}$ solar masses. However, smaller clumps and AU-scale structures, such as those tentatively identified in absorption in the Milky Way (Frail et al. 1994; Lauroesch, Meyer \\& Blades 2000; Richter, Sembach \\& Howk 2003) are still out of reach.  An alternative method relies on detailed spectroscopic analyses of selected sightlines that might actually help to unravel such small-scale structures in the Magellanic InterStellar Medium (MC ISM). The high-resolution spectroscopic abundance studies of Welty et al. (1999) based on the observations of SN 1987A by Vidal-Madjar et al. (1987) showed that the MC ISM exhibits properties similar to those of the warm, low-density Galactic disk clouds and can contain complex velocity structures with more than 20 components. Some of the absorbers observed toward SN 1987A also showed signatures of relatively cool gas {\\bf ($T<1500$ K)}.  The lack of information concerning the molecular content precluded to further characterize this cool phase, however.  With the {\\small FUSE}  data it is now possible to investigate the molecular content of this cool phase as well via H$_2$, CO, and HD. Although the HD molecule is the most abundant molecule in the interstellar medium after H$_2$, relatively few observations have been reported to date in the Galaxy and few comparisons with theoretical models of translucent clouds have been done (Le Petit et al. 2002). Similarly to H$_2$, the {\\bf presence of} HD molecule is best determined through its ground state transitions which only occur in the FUV domain. The detection is however difficult due to a low molecular abundance (HD/H$_2$ $\\approx$ 10$^{-6}$). The ground-state transitions of HD were first observed with {\\it Copernicus} in the local ISM towards bright, nearby stars (Spitzer \\& Cochran 1973; Wright \\& Morton 1979). Only recently did the {\\small FUSE} mission (Moos et al. 2000) re--open the FUV window with sufficiently high sensitivity to allow the detection of HD further away in the Galactic disk (Ferlet et al. 2000; Rachford et al. 2002; Lacour et al. 2004 in prep.). As to extragalactic search for HD, only two detections were found in the literature, one in the LMC by Bluhm \\& de Boer (2001) with the {\\small FUSE} satellite, the other in a Damped Lyman Alpha system towards PKS1232$+$0.82 (Varshalovich et al. 2001). The conclusions from the work of Bluhm \\& de Boer (2001) are consistent with the existence of a cool phase in the LMC with molecular abundances (H$_2$, HD, CO) typical of Galactic translucent sightlines (Rachford et al. 2002).  The project of this series of two papers is to provide a more extensive knowledge of the ``dense'', small-scale molecular structures within the Magellanic Clouds. We wish to point out that we specifically use the term ``dense'' to designate gas with proton density, $n_{\\rm H}$, above 100 cm$^{-3}$ in this work. In the present paper, emphasis is put on the analysis of H$_2$, HD, CO, Cl\\,{\\small I}, Cl\\,{\\small II}, Na\\,{\\small I}, Ca\\,{\\small II}, S\\,{\\small II}, Fe\\,{\\small II}, C\\,{\\small I}, C\\,{\\small I*}, and C\\,{\\small I**}, when available, towards five bright Magellanic Cloud stars.  Sk\\,$-$67D05, Sk\\,$-$68D135 and Sk\\,$-$69D246 are located in the LMC. AV\\,95, and Sk\\,159 are in the SMC. In $\\S$ 2 we describe the reduction of the {\\small FUSE}, {\\small VLT} and {\\small STIS} datasets. The analysis methods used to derive the species column densities are described in $\\S$ 3. Individual conclusions for each target are presented in $\\S$ 4. A preliminary discussion of these results is given in $\\S$ 5 and $\\S$ 6 summarizes our study. ",
        "conclusions": "\\begin{enumerate}  \\item We have studied the sightlines towards three LMC bright stars: Sk\\,$-$67D05, Sk\\,$-$68D135, and Sk\\,$-$69D246. The combination of ground--based observations ({\\small VLT}) along with archival data from UV satellites ({\\small FUSE}, {\\small STIS}) makes it possible to investigate H$_2$, CO, HD, Cl\\,{\\small I}, Cl\\,{\\small II}, C\\,{\\small I}, C\\,{\\small I*}, C\\,{\\small I**}, Na\\,{\\small I}, Ca\\,{\\small I}, Ca\\,{\\small II}, Fe\\,{\\small II}, and S\\,{\\small II}. We also report on the investigation of two SMC sightlines with the help of {\\small FUSE} and {\\small STIS} data. With the exception of Na\\,{\\small I}, Ca\\,{\\small I}, and Ca\\,{\\small II}, we have obtained measurements or upper limits on all the species observed towards the LMC targets.  \\item In the LMC we find $\\langle$HD/H$_2$$\\rangle$ $=$ 1.4 ppm, \\mbox{$\\langle$CO/H$_2$ $\\rangle$ $=$ 1.4 ppm} (0.8 ppm when excluding the large value obtained toward Sk\\,$-$67D05), and $\\langle$Cl\\,{\\small I}/H$_2$$\\rangle$ $=$ 1.3 ppm. After correcting for the MW reddening we also derive $\\langle$$N$(H$_2$)/$E_{B-V}$$\\rangle$ $=$ 3.5 $\\pm$ 1.8 $\\times$ 10$^{20}$ mag$^{-1}$ cm$^{-2}$. The analysis of the rotational levels of the H$_2$ molecules gives the kinetic temperature of the crossed molecular gas about $\\langle$$T_{01}$$\\rangle$ $\\approx$ 70 K. All these values are similar to values obtained in the MW disk and seem a bit surprising in the LMC where the radiation field is 5 to 3,000 times more intense than in the MW disk and the metallicity down by a factor 2. We suggest that these observations are consistent with the existence of cold, dense molecular structures embedded in the diffuse gas.  \\item In the SMC, we find $\\langle$$N$(H$_2$)/$E_{B-V}$$\\rangle$ $=$ 4.7 and 1.7 $\\times$ 10$^{20}$ mag$^{-1}$ cm$^{-2}$ for AV\\,95 and Sk\\,159 respectively. The sightline towards AV\\,95 in particular shows unexpectedly large HD/H$_2$ (21 ppm) ratio. We tentatively conclude that a D--reservoir is detected that can give access to a direct measure of the D/H in the SMC. However, more observations are needed, with better S/N and higher resolution, in order to confirm our analysis.  \\item The detailed analysis of the excited states of the carbon atoms gives an interesting diagnostic of the H number density along the sightlines towards Sk\\,$-$69D246 and AV\\,95. For both stars we find that the density is greater than 10$^2$ cm$^{-3}$ in the LMC molecular components. This leads to clump sizes (or filament) of the order of 0.1 pc. Indeed, we note that there are hints for an even denser medium and sizes of the order of 10$^2$ AU ($\\approx$ 0.001 pc). We reached the same conclusions for the Galactic molecular component sampled by the Sk\\,$-$68D135 and the Sk\\,$-$69D246 sightlines. We argue that this component probably corresponds to a sheet-like structure such as the ones recently inferred in the Galactic halo (Richter et al. 2003). More quantitative conclusions are to be given in the following paper based on the chemical model of Le Petit et al. (2002).  \\end{enumerate}"
    },
    "0404/hep-ex0404025_arXiv.txt": {
        "abstract": "We present the results of indirect searches for Weakly Interacting Massive Particles (WIMPs), with 1679.6 live days of data from the Super-Kamiokande detector using   neutrino-induced upward through-going muons. The search is performed by looking for an excess  of high energy muon neutrinos from WIMP annihilations in  the Sun, the core of the Earth, and the Galactic Center, as compared to the number expected from the atmospheric neutrino background. No statistically significant excess was seen. We calculate the flux limits in various angular cones around each of the above celestial objects.  We obtain conservative  model-independent upper limits on the WIMP-nucleon cross-section as a function of WIMP mass, and compare these results with the corresponding results from direct dark matter detection experiments. ",
        "introduction": "There is growing  evidence which indicates that the universe contains, at all length-scales,  ``dark matter''  which does not emit or absorb electromagnetic radiation at any known wavelength, but manifests itself only through gravity~\\cite{Trimble87}. Recent data from WMAP and other CMB experiments, large-scale structure surveys, Lyman-$\\alpha$ forest,  Type 1a supernovae  etc.~\\cite{WMAP,Tegmark,Perlmutter,Riess,Tonry} point with increasing accuracy towards  a standard cosmological model~\\cite{Turner02,Silk02}, in which the universe is flat, with about 5\\% baryons, 25\\% non-baryonic dark matter,   and about 70\\% dark energy. In this model,  the non-baryonic dark matter consists of ``Cold Dark Matter'' particles, which decoupled from rest of the matter and radiation while moving at non-relativistic velocities~\\cite{Peebles,Primack}.     Weakly Interacting Massive Particles (WIMPs) are one of the plausible  cold dark matter candidates~\\cite{Jungman96,Bergstrom00}.  WIMPs are stable particles which are predicted to occur  in extensions of the standard model, such as  supersymmetry.  WIMPs undergo only weak-scale interactions with matter, and could have masses in the GeV-TeV range.  If WIMPs exist, their relic abundance (which is governed by electroweak scale interactions) is remarkably close to the inferred density of dark matter in the universe~\\cite{Weinberg77}.  The lightest supersymmetric particle (LSP) of supersymmetric theories is the most theoretically well developed WIMP candidate~\\cite{Goldberg83,Ellis84}.  Most SUSY theories contain a multiplicatively conserved quantum number called R-parity, where R-parity is equal to~1 for ordinary standard model particles, and is equal to~$-1$ for their supersymmetric partners~\\cite{Jungman96}. If R-parity is conserved, superpartners can form and decay only in pairs. Thus, the LSP is stable and hence should be present in the Universe as a cosmological relic from the Big Bang.  The most likely candidate for this LSP is the neutralino~\\cite{Ellis84}.  The neutralino $\\tilde{\\chi}$ is a linear combination of the supersymmetric particles that mix after electroweak-symmetry breaking: \\begin{equation} \\tilde{\\chi}= a_{1}\\tilde{\\gamma} + a_{2}\\tilde{Z} + a_{3}\\tilde{H}_{1} + a_{4}\\tilde{H}_{2}, \\end{equation} \\noindent where $\\tilde{\\gamma}$ and $\\tilde{Z}$ are the supersymmetric partners of the photon and the $Z$ boson, and $\\tilde{H}_{1}$ and $\\tilde{H}_{2}$ are the supersymmetric partners of the Higgs bosons (Higgsino). The current lower limit on neutralino mass, without assuming anything about the SUSY breaking mechanism, and taking into account the results from LEP2 data,  is about 18~GeV~\\cite{Hooper02}. The upper limit on neutralino mass depends on the model assumed for SUSY breaking, and ranges from   500 GeV in some models~\\cite{Ellis03} and up to 10~TeV in other models~\\cite{Gondolo97}.  In this paper, we perform an indirect search for WIMP dark matter  using the Super--Kamiokande detector, by looking for a high energy neutrino signal resulting from WIMP annihilation in the Earth, the Sun, and the Galactic Center.  The signature would be an excess of neutrino-induced events coming from the direction of these objects over the background expected from atmospheric neutrinos.  Such a search is complementary  to direct detection experiments, which look for direct interaction of WIMPs with a nucleus in a low background detector.  However, both direct and indirect detection experiments probe the coupling of WIMPs to nuclei.  We then compare our results with those of direct detection of dark matter experiments. ",
        "conclusions": ""
    },
    "0404/astro-ph0404533_arXiv.txt": {
        "abstract": "We present late-time near-infrared (NIR) and optical observations of the type~IIn SN~1998S. The NIR photometry spans 333--1242~days after explosion, while the NIR and optical spectra cover 333--1191~days and 305--1093~days respectively. The NIR photometry extends to the $M'$-band (4.7~$\\mu$m), making SN~1998S only the second ever supernova for which such a long IR wavelength has been detected. The shape and evolution of the H$\\alpha$ and He\\,{\\sc i} 1.083~$\\mu$m line profiles indicate a powerful interaction with a progenitor wind, as well as providing evidence of dust condensation within the ejecta.  The latest optical spectrum suggests that the wind had been flowing for at least 430~years.  The intensity and rise of the $HK$ continuum towards longer wavelengths together with the relatively bright $L'$ and $M'$ magnitudes shows that the NIR emission was due to hot dust newly-formed in the ejecta and/or pre-existing dust in the progenitor circumstellar medium (CSM).  The NIR spectral energy distribution (SED) at about 1~year is well-described by a single-temperature blackbody spectrum at about 1200~K. The temperature declines over subsequent epochs. After $\\sim$2~years the blackbody matches are less successful, probably indicating an increasing range of temperatures in the emission regions.  Fits to the SEDs achieved with blackbodies weighted with $\\lambda^{-1}$ or $\\lambda^{-2}$ emissivity are almost always less successful. Possible origins for the NIR emission are considered. Significant radioactive heating of ejecta dust is ruled out, as is shock/X-ray-precursor heating of CSM dust.  More plausible sources are (a) an IR-echo from CSM dust driven by the UV/optical peak luminosity, and (b) emission from newly-condensed dust which formed within a cool, dense shell produced by the ejecta shock/CSM interaction. We argue that the evidence favours the condensing dust hypothesis, although an IR-echo is not ruled out. Within the condensing-dust scenario, the IR luminosity indicates the presence of at least 10$^{-3}$~M$_{\\odot}$ of dust in the ejecta, and probably considerably more.  Finally, we show that the late-time $(K-L')_0$ evolution of type~II supernovae may provide a useful tool for determining the presence or absence of a massive CSM around their progenitor stars. ",
        "introduction": "One of the challenges of supernova research is to obtain evidence about the nature and environment of the progenitor.  Type~IIn supernovae (SNe~IIn) are so-called because of the presence of narrow lines in the spectra, originating in a relatively undisturbed circumstellar medium (CSM) (Schlegel 1990).  The progenitors must therefore have undergone one or more mass-loss phases before explosion.  By using the interaction of the supernova explosion with the resulting CSM we can obtain clues about the nature of the progenitor. \\\\  The study of SNe IIn may also help us to understand the origin of dust in the universe.  More than 30 years ago, it was suggested (Cernuschi, Marsicano \\& Kimel 1965; Cernuschi, Marsicano \\& Codina 1967; Hoyle \\& Wickramasinghe 1970) that supernovae could be important sources of interstellar dust. More recent studies of the origins of grains (Gehrz 1989; Tielens 1990; Dwek 1998; Todini \\& Ferrara 2001; Nozawa et al. 2003) still invoke core-collapse SNe as significant contributors of dust to the interstellar medium in the past and present.  Dust grains associated with SNe can be detected both via their attenuation effects on optical/near-IR radiation, and by their thermal re-radiation at longer IR wavelengths of the absorbed energy. Observation of attenuation provides a relatively unambiguous way of demonstrating grain condensation in the ejecta. This process is revealed by the relative suppression of the red wings of the broad spectral lines.  However, apart from SN~1998S (Leonard et al. 2000; Gerardy et al. 2000; this work), only for two events, SN~1987A and SN~1999em, has the condensation of dust in the ejecta been demonstrated via the attenuation of the red wing of the line profiles (Danziger et al. 1991; Lucy, Danziger \\& Gouiffes 1989; Lucy et al. 1991; Elmhamdi et al. 2003). \\\\  The second method of demonstrating grain condensation via the appearance of strong IR radiation (an `IR excess') as a supernova ages, is more challenging.  At least 10 SNe have shown a strong, late-time (i.e. $t>100$~d) NIR excess (relative to that expected by extrapolation from the optical region; see references in Gerardy et al. 2002).  The natural interpretation is that this radiation is produced by hot dust.  However, while such grains may be the result of condensation processes in the expanding SN ejecta, they may also have formed in pre-explosion circumstellar material ejected during one or more mass-loss episodes experienced by the progenitor star. In this latter case, the IR emission is produced as an ``IR echo'' of the maximum-light luminosity from the CSM dust.  Thus, in the interpretation of such strong NIR emission from SNe, a key issue is the location of the source.  Only in the case of SN~1987A has conclusive evidence for dust condensation in the ejecta been presented on the basis of the NIR and mid-IR emission (Moseley et al. 1989; Whitelock et al. 1989; Suntzeff \\& Bouchet 1990; Dwek 1991; Lucy et al. 1991; Dwek et al. 1992; Meikle et al. 1993; Roche, Aitken \\& Smith 1993; Colgan et al. 1994).  Recently, Dunne et al. (2003) have presented evidence based on sub-mm observations that $\\sim2-4~M_{\\odot}$ of cold dust is present in the supernova remnant (SNR) Cassiopeia~A. They argue that the dust must have been produced during the explosion. A similar result for the Kepler SNR was presented by Morgan et al. (2003) although in this case an alternative CSM origin for the dust was not ruled out, which would be more consistent with the most commonly believed origin for this supernova, i.e. a type Ia explosion. Such sub-mm studies may provide important support for the hypothesis that supernovae are a major source of dust. However, at the moment they have the disadvantages that the supernova types are uncertain, and that the number of examples which can be studied in this way is quite small.  In addition, the results of Dunne et al. (2003) and Morgan et al. (2003) have been criticised by Dwek (2004) who points out that, given the mass of condensible elements expected in the ejecta, a very unlikely condensation efficiency of 100\\% would be required to yield even 2~M$_{\\odot}$ of dust.  Moreover thermal sputtering would reduce this mass of dust.  Dwek offers an alternative scenario where the sub-mm emission is from conductive needles (''metallic whiskers'') which may be long enough to be very efficient emitters at sub-mm wavelengths thus requiring a dust mass of no more than $10^{-3}$~M$_{\\odot}$ to account for the observed radiation.\\\\  Lastly, we note that the unique wavelength and temporal coverage of SN~1987A also made it possible to infer dust condensation from (a) spectroscopic evidence for the temporal depletion of gaseous-phase elements in the ejecta e.g., [0\\,{\\sc i}] (6300,~6364~\\AA), [Mg\\,{\\sc i}] (4571 \\AA), Si\\,{\\sc i} (1.2~$\\mu$m), [Si\\,{\\sc i}] (1.65~$\\mu$m) and [Ca\\,{\\sc ii}] (7300~\\AA) (Lucy et al. 1991; Spyromilio et al. 1991; Meikle et al. 1993), and (b) analysis of the luminosity budget of the ejecta (e.g. Whitelock et al. 1989).  As observing facilities improve, these last two methods of dust detection may eventually be applicable to SNe at more typical distances.\\\\  To provide conclusive evidence of SNe as major dust sources, we need to assemble a statistically significant sample of typed supernovae in which grain formation can be reliably tested. In addition, by studying the evolution of recent supernovae in real time we can obtain information about the physics of the dust-formation process.  In order to disentangle the two possible origins of the NIR radiation, it is vital to acquire IR observations spanning as wide a wavelength and temporal range as possible. The longer NIR wavelengths ($HKLM$ bands) are especially important since the flux at these wavelengths is less likely to be contaminated by scattered light or strong H, He emission lines, allowing a simple thermal interpretation. \\\\  \\begin{figure*} \\vspace{15.5cm} \\special{psfile=fig_98S_J.ps hoffset=-230 voffset=-430 hscale=190 vscale=170 angle=0} \\caption[SN~1998S] {A $J$-band image (83'' $\\times$ 83'') showing SN~1998S (centre) at 333 days after shock breakout. North is to the top and east to the left. The image was obtained with the United Kingdom Infrared Telescope, Hawaii.} \\label{sn98S} \\end{figure*}  An opportunity to address both the progenitor problem and the question of supernovae as dust sources was provided by the occurrence of the SN IIn 1998S.  This is the most intensively-studied type~IIn event (Bowen et al. 2000; Fassia et al. 2000, 2001; Gerardy et al. 2000, 2002; Leonard et al. 2000; Lentz et al. 2001; Liu et al. 2000; Roscherr \\& Schaefer 2000; Anupama, Sivarani \\& Pandey 2001; Chugai 2001; Wang et al. 2001; Chugai et al. 2002; Gruendl et al. 2002; Li et al. 2002; Pooley et al. 2002; Meikle et al. 2003).  \\subsection{SN~1998S at early epochs}\\label{early}  SN~1998S is one of the brightest SN IIn ever seen, reaching $V=+12.17$ on 1998 March 18 (Fassia et al. 2000).  It was discovered by Z. Wan (Li \\& Wan 1998) on 1998 March 2.68 UT (JD 2,450,875.18) in NGC 3877, an edge-on Sc spiral galaxy.  The supernova lies about 16'' west and 46'' south of the galaxy nucleus (see Fig.~1).  It was identified as a SN~IIn by Filippenko \\& Moran (1998).  Liu et al. (2000) point out that the $BV$ light curves together with the spectroscopic behaviour at 2--3 months are similar to those of the type IIL SN~1979C.  Hamuy (2003) suggests that SNe IIn are actually a subset of what he terms ``dense wind'' type~II supernovae which includes SN~1979C.  A pre-discovery limiting magnitude of $\\sim$+18 obtained on 1998 February 23.7 (Leonard et al. 2000) indicates that SN~1998S was discovered within a few days of shock breakout.  Following Chugai (2001) we adopt 1998 February 24.7 UT = JD 2,450,869.2, 6~days prior to discovery, as the explosion epoch, $t=0$~days.  Relative to this epoch, the zero epochs adopted by other authors are later by, respectively, +5~days (Leonard et al. 2000), +6~days (Fassia et al. 2000, 2001; Pooley et al. 2002), $\\sim$+23~days (Bowen et al. 2000; Gerardy et al. 2000, 2002).  Throughout this paper, epochs quoted from other authors have been shifted to our zero epoch definition.  We adopt a distance of 17~Mpc (Tully 1988), but note that a smaller value of 15.5 Mpc is sometimes adopted (Sanders \\& Verheijen 1998; Gerardy et al. 2002).\\\\  The earliest optical spectra of SN~1998S show a blue continuum with emission features superimposed. A rough blackbody fit yields T$\\sim$25,000~K, but with a blue excess (Leonard et al. 2000; Fassia et al.  2001).  The emission lines are identified with H\\,{\\sc i} (Balmer series), He\\,{\\sc i}, He\\,{\\sc ii}, C\\,{\\sc iii} and N\\,{\\sc iii}.  The high-ionisation carbon and nitrogen lines are also commonly observed in Wolf-Rayet stars (Leonard et al. 2000). The emission lines have a broad base (e.g., H$\\alpha$ FWZI$\\sim$20,000 km s$^{-1}$), but a narrow `peaked' unresolved centre.  The lines are symmetrical about the local standard of rest.  This is quite surprising since at such an early phase most of the receding part of the supernova should be occulted by the photosphere.  However, this may actually constitute some of the earliest evidence of a strong ejecta-CSM wind interaction. It has been shown theoretically that, at this phase, one might expect the photosphere to reside in a geometrically thin, opaque cool dense shell (CDS) formed by the radiative shock propagating into the extended stellar envelope (Grasberg, Imshennik \\& Nadyozhin 1971; Falk \\& Arnett 1977; Chevalier \\& Fransson 1985; 1994).  Chugai (2001) has argued that these lines originate from the CDS and that the line broadening is caused by the multiple Thomson scattering of line photons on thermal electrons of the circumstellar gas.  The weakness of broad ejecta lines during this early era is attributed to the sharp photosphere at the SN ejecta boundary.  The blue excess can also be attributed to the CDS, since the significant optical depth can yield an increase in continuum absorptive opacity with wavelength, due to both bound-free, (i.e., the Paschen continuum) and free-free processes (Leonard et al. 2000). \\\\  By about 3~weeks after the explosion, the emission lines had almost totally vanished.  This disappearance is attributed to the dense inner-CSM being overrun by the ejecta.  This persisted for $\\sim$3--7 weeks. During this time weak narrow lines are present, superimposed on the continuum.  These are caused by the photoionization and heating of the undisturbed wind by the supernova flash and subsequent X-ray emission from the shocked wind and ejecta.  High-resolution echelle spectra of these lines were obtained by Bowen et al. (2000) and Fassia et al. (2001) on days~23 and 42.  These observations succeeded in resolving the lines.  Between day~23 and day~42 the red wing of the [O\\,{\\sc iii}]~5007~\\AA\\ profile moved redward, broadening the line, so that by the latter epoch the profile was symmetrical about a velocity of +847~km s$^{-1}$.  A high-resolution observation of Gruendl et al. (2002) at 1 year showed that the [O\\,{\\sc iii}]~5007~\\AA\\ profile was still symmetrical about this redshift. This behaviour can be explained as being due to the effect of the finite light travel time across a CSM having a denser outer zone.  As the radiation-induced ionisation (due to the flash + shock radiation) propagates across the CSM, it takes longer for the resulting nebular emission to reach us from the far side.  At any given time, therefore, the observer sees radiation from an ellipsoidal surface propagating through the CSM (Morrison \\& Sartori 1969).  On day~23, the vertex region of the ellipsoid had not yet reached the inner boundary of the denser zone, thus accounting for the weakness of the red wing of the line.  However, the lack of any further shift in the line centre between 46~days and 1~year indicates that the vertex had already reached the dense zone by 46~days.  The early evolution of the line asymmetry can be quantitatively explained if the inner boundary of the denser zone was at a distance of 16.5 light days (2850~AU).  We note that this distance is consistent with the observed persistence of the line for at least a year since, assuming a maximum ejecta velocity of 10,000 km s$^{-1}$ (Leonard et al. 2000; Fassia et al. 2001), at 1~year it would still only have reached 2050~AU and so would not yet have disturbed the dense outer CSM region (but see subsection 4.2 below).  Fassia et al. (2001) adopted +847~km s$^{-1}$ as the centre of mass velocity for the SN~1998S system. Justification for this is provided by the light travel time argument together with the lack of any further shift between day~46 and 1~year.  From the narrow forbidden lines such as [O\\,{\\sc iii}]~5007~\\AA, an undisturbed CSM velocity of about 40~km s$^{-1}$ is obtained, which is characteristic of a red supergiant (RSG) wind.  From the intensity ratio of [O\\,{\\sc iii}] (4959~\\AA\\ + 5007~\\AA) to [O\\,{\\sc iii}] 4363~\\AA, Fassia et al. (2001) infer a CSM density of at least $1.5\\times10^6$~cm$^{-3}$ at 185~AU, implying a CSM mass exceeding 0.005~M$_{\\odot}$, and a mass-loss rate exceeding around $2\\times10^{-5}$~M$_{\\odot}$ yr$^{-1}$.  This is consistent with the radio/X-ray estimate (scaled to a 40~km~s$^{-1}$ wind velocity) of $\\sim5\\times10^{-4}$~M$_{\\odot}$ yr$^{-1}$ (Pooley et al. 2002), and is comparable to that of SN 1979C which had a mass loss rate of $\\sim1.2\\times10^{-4}$~M$_{\\odot}$ yr$^{-1}$ (Lundqvist \\& Fransson 1988) . \\\\  The behaviour of the narrow allowed H~{\\sc i}, He~{\\sc i} CSM lines was more complex.  Not only did they exhibit asymmetric P~Cygni profiles, but there were clearly two velocity components.  The slower component is attributed to the same origin as the forbidden lines {\\it viz.} the photo-ionised, unaccelerated CSM.  The profile of this component was probably a combination of emission from recombination and possibly collisional excitation, together with a classical P~Cygni line due to scattering from the populated excited levels.  The broad absorption component is more difficult to explain.  It extends to a velocity of around 350~km s$^{-1}$ which is too fast for a RSG wind. It may be that, as in the case of SN~1987A, the SN~1998S progenitor went through a fast-wind phase prior to explosion (Fassia et al. 2001).  An alternative explanation is that the CSM close to the supernova was accelerated by photospheric photons, or by relativistic particles from the ejecta/CSM shock (Chugai et al. 2002).  Another possibility is that the faster component arose in shocked clumps within the CSM wind (Chugai et al. 2002).  By day~50, broad emission lines in P${\\beta}$ and He~{\\sc i} 1.083~$\\mu$m had appeared. The earliest second season optical spectra (day~78) showed a similar development in H${\\alpha}$ and the Ca\\,{\\sc ii} triplet (Fassia et al. 2001). The line profiles became increasingly ``square-shaped'', characteristic of strong interaction between the expanding ejecta and a dense CSM (Leonard et al. 2000; Fassia et al. 2000).  This was confirmed by the detection of radio and X-ray emission (Pooley et al. 2002) plus an unusually slow decline in V and I by $\\sim$300 days (Li et al. 2002).  \\\\  In summary, the wind of the SN~1998S progenitor underwent discrete changes in the rate of mass outflow, resulting in several distinct CSM zones.  The fading of the early-time broad emission components by about day~18 indicates that the radius of the first boundary was less than 100~AU (Leonard et al. 2000; Fassia et al. 2001). Beyond this boundary the density presumably fell quite sharply.  The later appearance of broad emission features led Leonard et al.  (2000) and Fassia et al. (2001) to suggest the presence of a second boundary where the CSM underwent a density increase.  For the explosion epoch adopted here, the reappearance occurred between days~42 and 50 corresponding to $260\\pm20$~AU.  However, the reality of this second boundary was contested by Chugai et al. (2002) who suggested the reappearance was due simply to the CDS becoming optically thin. Finally, as argued above, the evolution of the narrow-line profile suggests a third boundary at about 2850~AU where the CSM again underwent a density increase.  For a wind velocity of 40~km s$^{-1}$, these three boundaries correspond to material which left the progenitor at epochs $-12$~yr, $-30$~yr and $-340$~yr relative to the explosion date.  Pooley et al. (2002) also note that the best fit to their radio data requires significant clumping or filamentation in the CSM, a condition already inferred in two other type~II SNe, 1986J and 1988Z. Early-time spectropolarimetry by Leonard et al. (2000) and Wang et al. (2001) indicate asphericity in the ejecta, CSM or possibly both. \\\\  SN~1998S is unique in that it allowed the first-ever good IR spectroscopic coverage of a SN~IIn (Gerardy et al. 2000; Fassia et al. 2001).  In the $J$-band, we see Paschen $\\beta$, Paschen $\\gamma$ and He~{\\sc i} 1.083~$\\mu$m lines.  Their evolution was similar to that seen in the optical.  At the earliest times broad-based, peaked profiles were present. These faded by day~23, being replaced by broad, square-shaped profiles by day~50.  Between days $\\sim$15 and $\\sim$60 a strong, unresolved He~{\\sc i} 1.083~$\\mu$m CSM line was superimposed on the ejecta/CSM broad lines. By day~50, the $HK$-band was dominated by Paschen~$\\alpha$.  By day~114, first-overtone CO emission was clearly present in the $K$-band.  The presence of CO in type~II SNe is increasingly regarded as ubiquitous.  In all cases where $K$-band observations have been carried out in the period 3-6 months post-explosion, CO has been detected, e.g., SN~1987A (IIpec) (Catchpole \\& Glass, 1987; McGregor \\& Hyland, 1987; Spyromilio et al. 1988; Meikle et al. 1993), SN~1995ad (IIP) (Spyromilio \\& Leibundgut 1996), SN~1998S (IIn) (Gerardy et al. 2000; Fassia et al. 2001), SN~1999dl (IIP) (Spyromilio, Leibundgut \\& Gilmozzi 2001), SN~1999em (IIP) (Spyromilio et al.  2001; Gerardy et al. 2002), SN~1999gi (IIP) (Gerardy et al. 2002) and SN~2002hh (IIP) (Pozzo et al., in preparation). In addition, CO has been detected in the type~Ic SN~2000ew (Gerardy et al. 2002).  The rotation-vibration states of CO are a powerful coolant (especially the fundamental band lines).  The presence of CO is suspected to be a necessary condition for dust condensation to occur in the ejecta.  Modelling of the SN~1998S spectra suggests a CO velocity of $\\sim$2000 km s$^{-1}$ (Gerardy et al. 2000; Fassia et al. 2001).  From this, Fassia et al. (2001) deduced a core mass of 4~M$_{\\odot}$ implying a massive progenitor. The actual mass of CO derived was 10$^{-3}$~M$_{\\odot}$.  This is much smaller than the $>0.1$~M$_{\\odot}$ of C and $>1.0$~M$_{\\odot}$ of O likely to be present in the ejecta (Woosley \\& Weaver 1995), implying that most of the C and O remain in the gas phase and are thus available to condense into, respectively, graphite and silicate grains. \\\\  On day~136, Fassia et al. (2000) measured the IR flux out to a wavelength of 3.8~$\\mu$m ($L'$-band).  This revealed a remarkable IR excess of $K-L'=+2.5$.  To produce the observed flux the lowest possible radius of the IR emission region is given by a blackbody with a temperature close to the dust evaporation temperature of $\\sim$1500~K.  For dust condensing in the ejecta to attain a sufficiently large blackbody radius by 136~days would require an expansion velocity of 11,000~km s$^{-1}$ (see Fassia et al. 2000). Such high velocities were seen only in the extreme outer zones of the H/He envelope. No metals were seen at such high velocities.  Fassia et al. (2000) concluded that the IR excess at this epoch cannot, therefore, have been due to grain condensation in the ejecta.  It must instead have been produced by an IR echo of the maximum-light luminosity from pre-existing dust in the CSM. This conclusion is based only on the intensity of the $L'$-band flux.  It is strengthened if we fit a blackbody to both the $K$ and $L'$ fluxes. After correction for the optical photosphere, this yields $K-L'\\approx3.7$ corresponding to a temperature of only 650~K, requiring the dust to lie at an equivalent velocity of at least 67,000~km s$^{-1}$! \\\\  SN~1998S was exceptionally luminous, reaching a de-reddened M$_B = -19.6$ (Fassia et al. 2000). This is around $\\times$10 the typical luminosity of a type~II SN.  The excellent early-time coverage achieved in the optical and NIR allowed Fassia et al. (2000) to examine the bolometric light curve.  Both blackbody and UVOIR (i.e., ultraviolet-optical-infrared range) fits indicate that the total energy radiated in the first 40~days exceeded $10^{50}$ ergs, which is again $\\times$10 the typical value for type~II SNe.  Between days 96 and 136, the bolometric light curve is well-reproduced by the radioactive decay luminosity of 0.15~M$_{\\odot}$ $^{56}$Ni.  However, by this era the ejecta/CSM shock energy must also have been making a contribution.  \\subsection{SN~1998S at later epochs}  Much of the early-time behaviour of SN~1998S can be attributed to the interaction of the supernova with a pre-existing, dusty, perhaps disk-like CSM.  SN~1998S remained observable from X-rays to radio for over 3 years (Pooley et al. 2002; this work) due to the ongoing conversion of the SN kinetic energy to radiation via the ejecta/CSM interaction. We continued regular observations during this phase, to day~1242.  A brief, preliminary presentation of this work was given in Meikle et al. (2003). Other late-time studies have been presented: optical spectroscopy to 499~d (Leonard et al. 2000); high-resolution spectroscopy at $\\sim 370$~d (Gruendl et al. 2002); NIR spectroscopy to $\\sim 380$~d and NIR photometry to 819~d (Gerardy et al. 2000, 2002); X-ray observations to 1055~d and radio observations to 1065~d (Pooley et al. 2002).\\\\  In this paper we present and discuss NIR photometry of SN~1998S spanning 333 to 1242 days post-explosion, optical spectra covering 305 to 1093~days and IR spectra covering 333 to 1191~days.  The NIR photometry presented here extends as far as the $M'$-band (CW 4.7~$\\mu$m), the first time that such a long IR wavelength has been detected in any SN other than the exceptionally close SN~1987A.  The paper is organised as follows. NIR photometry and optical/NIR spectroscopy are presented in Section~2.  The evolution of the H$\\alpha$ and He~{\\sc i}~1.083~$\\mu$m profiles are examined in detail, and for this purpose the data are augmented by spectra from other sources.  In Section~3 we compare the NIR spectral energy distribution (SED) with pure blackbody functions and blackbodies weighted by emissivities $\\lambda^{-1}$ and $\\lambda^{-2}$.  In Section~4 we examine possible energy sources for the post-300~d NIR emission, and discuss the location of the IR-emitting dust. Evidence of episodic mass-loss from the progenitor is discussed.  We also compare the late-time $(K-L')_0$ colour evolution of SN~1998S with that of other type~II SNe and suggest that it provides a useful identifier of supernovae whose progenitors had a massive CSM.  Conclusions follow in Section~5.  \\begin{table*} \\centering \\caption{Log of infrared imaging of SN~1998S at UKIRT} \\begin{minipage}{\\linewidth} \\renewcommand{\\thefootnote}{\\thempfootnote} \\renewcommand{\\tabcolsep}{0.4cm} \\begin{tabular}{clrll} \\hline JD(2450000+) & Date &  Epoch\\footnote{Days after explosion, assumed to be 1998 February 24.7 UT (JD 2450869.2).} & Filters\\footnote{See subsection 2.1 for details.} & Standard star\\footnote{FS standards were used for $JHK$, and HD standards for $L'M'$.} \\\\ \\hline 1202.2  & 1999 Jan 23  &   333.0~   & $J$, $H$, $K$, $L'$                &  FS21, HD105601 \\\\ 1243.2  & 1999 Mar  3  &   374.0~   & $J$, $H$, $L'$, nb$M$              &  FS15, HD84800, HD105601, HD106965 \\\\ 1274.8  & 1999 Apr  4  &   405.5~   & $J$98, $H$98, $K$98              &  FS21  \\\\ 1333.5  & 1999 Jun  2  &   464.2~   & $J$, $H$, $K$, $L'$, nb$M$           &  FS21, HD106965 \\\\ 1533.2  & 1999 Dec 19  &   664.0~   & $J$98, $H$98, $K$98, $L'$98, $M'$98  &  FS130, HD105601 \\\\ 1567.5  & 2000 Jan 23  &   698.2~   & $J$98, $H$98, $K$98, $L'$98, $M'$98  &  FS21, HD105601  \\\\ 1579.0  & 2000 Feb  4  &   709.8~   & $M'$98                       &  HD105601 \\\\ 1696.2  & 2000 May 30  &   827.0~   & $J$98, $H$98, $K$98, $L'$98, $M'$98  &  FS21, HD105601 \\\\ 1916.8  & 2001 Jan  7  &  1047.5~  & $J$98, $H$98, $K$98, $L'$98, $M'$98  &  FS21, HD105601, HD106965 \\\\ 2067.0  & 2001 Jun  5  &  1197.8~  & $J$98, $H$98, $K$98              &  FS131 \\\\ 2111.5  & 2001 Jul 19  &  1242.2~  & $L'$98                       &  HD106965 \\\\ \\hline \\vspace{-0.8cm} \\end{tabular} \\end{minipage} \\end{table*}  \\begin{table*} \\centering \\caption[]{Infrared photometry of SN~1998S} \\begin{minipage}{\\linewidth} \\renewcommand{\\thefootnote}{\\thempfootnote} \\renewcommand{\\tabcolsep}{0.5cm} \\begin{tabular}{crccccc} \\hline JD(2450000+) &  Epoch\\footnote{Days after explosion.}& $J$ & $H$ & $K$ & $L'$ & $M'$ \\\\ \\hline 1202.2  &     333.0  &   ~16.653(12)\\footnote{Figures in brackets give the statistical error, in units of the magnitude's least significant one or two digits.} & 14.809(5)~~ & 13.344(2)~ & 11.560(28) &  -- \\\\ 1243.2  &     374.0  &   16.799(34)   &  14.922(13)  &   -- &   11.52(17)~ &   11.291(86) \\\\ 1274.8  &     405.5  &   17.077(25)   &  15.269(8)~  &   13.431(2)~ &      --      &      --      \\\\ 1333.5  &     464.2  &   17.609(29)   &  15.537(12)  &   13.806(4)~ &   11.69(13)~ &  11.14(16)~  \\\\ 1533.2  &     664.0  &   18.025(33)   &  16.421(14)  &   14.448(4)~ &   11.856(18) &   11.44(14)~ \\\\ 1567.5  &     698.2  &   18.341(35)   &  16.728(15)  &   14.556(4)~ &   11.970(26) &   11.82(14)~ \\\\ 1579.0  &     709.8  &      --        &     --       &      --      &       --     &   12.30(12)~ \\\\ 1696.2  &     827.0  &   18.680(27)   &  17.148(16)  &   14.992(4)~ &   12.200(33) &      $>11.9(2\\sigma)$\\footnote{For this epoch we did not detect the SN in the M' filter band image.}      \\\\ 1916.8  &    1047.5  &   19.280(48)   &  17.956(27)  &   15.843(7)~ &   12.985(39) &   12.84(15)~ \\\\ 2067.0  &    1197.8  &   19.798(83)   &  18.439(77)  &   16.366(13) &   [13.29(3)]\\footnote{Estimated by linear interpolation within $L'$-band light curve.}      &      --      \\\\ 2111.5  &    1242.2  &      --        &     --       &   [16.52(2)]\\footnote{Estimated by linear extrapolation of $K$-band light curve.} &   13.40(15)~ &      --      \\\\ \\hline \\vspace{-0.8cm} \\end{tabular} \\end{minipage} \\end{table*} ",
        "conclusions": "We have presented post-300~d infrared and optical observations of the SN IIn 1998S which extend to longer wavelengths and later epochs than has ever been achieved for this type of supernova. SN~1998S is only the second ever supernova for which $M'$-band has been detected. The IR photometry presented here, together with an earlier measurement by Fassia et al. (2000), has revealed a strong IR-excess during the period 136 to 1242~days.  Spectroscopy in the $HK$ bands show a continuum rising towards longer wavelengths.  Broad H$\\alpha$ and He\\,{\\sc i} 1.083~$\\mu$m line profiles present at $\\sim$4~months had undergone a dramatic change when they were recovered by Gerardy et al. (2000) at about 8~months.  By this time, both lines took the form of a triple peak structure with a central peak close to the supernova rest frame velocity, and with the two other peaks lying at --4860~km s$^{-1}$ and +3400~km s$^{-1}$ respectively. Gerardy et al. argue that this could be indicative of emission from a disk-like or ring-like CSM.  Our observations show that as the supernova evolved, the redward half of the profiles became increasingly weak relative to the blue part so that by the end of the second year the red side had almost completely vanished.  We agree with Gerardy et al. (2000) and Leonard et al. (2000) that this is strongly indicative of dust condensation within the ejecta of the supernova.  The change in the appearance of the H$\\alpha$ profile by 1093~d together with the early-time evolution of the narrow CSM [O\\,{\\sc iii}]~5007~\\AA\\ line may indicate a density discontinuity at about 2850~AU implying that the progenitor mass-loss rate underwent a decline about 340~years ago.  The velocity of the extreme blue wing of the H$\\alpha$ line at 1093~days implies that the CSM extended to at least 3650~AU indicating that mass loss had been taking place for at least 430~years prior to explosion. \\\\  The intensity and rise of the $HK$ continuum towards longer wavelengths together with the relatively bright $L'$ and $M'$ magnitudes suggests strongly that the IR emission was due to hot dust condensing in the ejecta and/or pre-existing in the CSM.  We find that the SED represented by the $HKL'M'$ magnitudes and by the $HK$ continua (when available) can be well-reproduced with single-temperature blackbody functions having temperatures of about 1250~K at 333~days declining to about 1100~K at 464~days.  Subsequent epochs yield temperatures declining to about 900~K at 3~years, although the fits are somewhat poorer.  This probably indicates an increasing range of temperatures in the emission regions.  Fits achieved with blackbodies weighted with a $\\lambda^{-1}$ or $\\lambda^{-2}$ emissivity are almost always less successful than those yielded by pure blackbodies. Newly-condensed iron whiskers having a $\\lambda^{+1}$ emissivity also seem unlikely as they would have to form beyond the CDS.  We conclude that the IR emission arises from dust in the form of IR-opaque clumps or having a large ($a>1\\mu$m) grain size. \\\\  We have considered possible origins for the strong, post-300~d IR emission.  A significant amount of direct heating of condensed ejecta dust by ongoing radioactive decay is ruled out since the observed IR emission exceeds the radioactive luminosity by a substantial factor. We also rule out the possibility of direct heating of pre-existing CSM dust by the shock, since it would not have reached the edge of the dust-free cavity.  X-ray precursor heating of CSM dust is also ruled out since it would imply the additional presence of an IR-echo of much greater luminosity than the observations could support.\\\\  We are therefore left with the IR-echo and condensing dust hypotheses. In favour of the IR-echo explanation, Dwek's (1983) model with a $\\lambda^{-1}$ dust emissivity suggests that the observed IR luminosity could be produced by a slow wind of $\\sim$10$^{-4}$M$_{\\odot}$~yr$^{-1}$ which is similar to that indicated by optical, X-ray and radio observations of SN~1998S.  This assumes the typical ISM gas-to-dust ratio of 160 in the progenitor CSM.  In addition, if we adopt a $\\lambda^{-1}$ emissivity law, then condensing dust would be unable to supply sufficient IR luminosity. However, our measurements tend to disfavour SEDs weighted with a $\\lambda^{-1}$ or $\\lambda^{-2}$ emissivity.  Moreover, for the IR-echo scenario to work for SN~1998S at $t>300$~d, we must invoke a rather extreme case where at least 20\\% of the UV-optical luminosity is absorbed by the CSM dust.  In addition, against the IR-echo scenario is the disagreement in the observed and predicted decline rates and SEDs, although it may be possible to eliminate these discrepancies through more detailed IR-echo modelling, including asymmetric and episodic cases.  We therefore do not rule out the IR-echo scenario.  At the earlier epoch of 136~d, an IR-echo seems to be the only plausible explanation for the IR luminosity. \\\\  In favour of the condensing dust hypothesis is the success of the single temperature pure blackbody fits at about 1~year. Dust condensing in optically thick clumps within a physically relatively thin shell symmetrical about the supernova centre-of mass might be expected to exhibit such temperature uniformity.  (Light travel time effects would be negligible in this case since even at 2~years the ejecta radius would only be about 10~light days.)  Moreover, the coincidence at about 1~year of the blue-peak velocities of the H$\\alpha$ and He\\,{\\sc i} 1.083~$\\mu$m lines with that derived from the blackbody fits to the IR magnitudes suggests a physical connection.  This is supported by the similar values at about 1~year for the optical depths derived from the H$\\alpha$ line profile, the He\\,{\\sc i}~1.083~$\\mu$m line profile, and that obtained from the blackbody-derived dust covering fraction.  We suggest that this connection is the CDS formed as a result of the ejecta-CSM interaction. It is possible that conditions in the CDS would allow the condensation of dust.  Moreover the highest blackbody temperature we derive is 1250~K, which is also believed to be compatible with grain condensation, though only slightly higher temperatures may not be (Gehrz 1988).  Assuming that all the dust at about 1~year is at the same temperature, we estimate that as early as 1~year, over $10^{-3}$~M$_{\\odot}$ had formed.  The apparent decrease in the dust mass after about 700~d together with the reduced success of the blackbody fits to the IR SED may indicate that dust condensation was actually still ongoing but that an increasing proportion had cooled to temperatures where the bulk of the flux was emitted beyond NIR wavelengths.  Given these arguments, we tend to favour dust condensation as the origin of {\\it both} the post-300~d IR luminosity and the line profile attenuation.\\\\  It should be stressed that in the dust condensation scenario, the total dust mass produced could be very much larger than the $\\sim\\times10^{-3}$~M$_{\\odot}$ lower limit. Most of the dust could have been invisible due to the clump opacities being much higher than that necessary to produce a blackbody spectrum, and increasingly large amounts could have gone undetected due to the cooling of an increasingly large proportion of the dust beyond the reach of the NIR observations.  The possibility of very opaque clouds is also raised by Elmhamdi et al. (2003), who obtain a somewhat weaker lower limit of $\\sim10^{-4}$~M$_{\\odot}$ of dust for the type~IIP SN~1999em based on the evolution of the [O~{\\sc i}]~6300~\\AA\\ profile.  In addition, our proposal that dust formed in the CDS does not exclude the possibility that dust also condensed deep within in the metal-rich SN ejecta interior.  Indeed the presence of a CO emission zone on day~115 at a velocity of only 2200~km~s$^{-1}$ (Fassia et al. 2001) suggests that a cool dust-forming environment deep in the ejecta might have developed. Our dust mass lower limit for SN~1998S is consistent with the $\\sim$0.1--0.3~M$_{\\odot}$ of condensed dust in SN~1987A deduced from metal depletion (Lucy et al. 1991; Dwek et al. 1992; Spyromilio \\& Graham 1992). However, SN~1987A was peculiar.  Moreover, SNe~IIn make up only 10--15\\% of all type~II events (Cappellaro \\& Turatto 2001).  Of the three CCSNe (87A, 98S, 99em), only SN~1999em may be described as `typical', but it is only one event.  It is therefore premature to conclude that high rates of dust condensation occur in all type~II events.  Consequently, the exact significance of SNe~II for cosmic dust production is undecided.  For further progress to be made, studies such as presented here should be carried out for a statistically significant sample of supernovae, and extended to even longer IR wavelengths. \\\\  Comparison of the $(K-L')_0$ evolution for a number of type~II SNe, suggests that this may be a useful tool for discriminating between supernovae whose progenitors possessed or lacked a massive CSM.  We propose that the appearance of a strong $(K-L')_0$ excess [$(K-L')_0>1.6$] during the first year is indicative of a massive CSM. The IR flux at this time can be a mixture of emission from CDS dust and an IR-echo i.e., two populations of dust may exist.  The newly-formed ejecta dust would be responsible for the attenuation of the red sides of the line profiles.  Determination of the relative contributions by the two populations to the IR flux requires good wavelength coverage at late-times, extending as far as the $L'$ and $M'$ bands or even further to the mid-IR.  For the progenitors of SNe which do not show an IR-excess in the 100--350~day period (e.g., SNe~1987A, 2002hh), it is much less likely that they were surrounded by a massive CSM in their immediate vicinity.  If such SNe subsequently produce an IR excess this is probably due to dust condensation throughout the ejecta as was inferred for SN~1987A."
    },
    "0404/astro-ph0404019_arXiv.txt": {
        "abstract": "The INTEGRAL Burst Alert System (IBAS) is the software for real time detection of Gamma Ray Bursts (GRBs) and the rapid distribution of their coordinates. IBAS has been running almost continuously at the INTEGRAL Science Data Center since the beginning of the INTEGRAL mission, yielding up to now accurate localizations for 12 GRBs detected in the IBIS field of view. IBAS is able to provide error regions with radii as small as 3 arcminutes (90\\% c.l.) within a few tens of seconds of the GRB start. We present the current status of IBAS, review the results obtained for the GRBs localized so far, and briefly discuss future prospects for using the IBAS real time information on other classes of variable sources. ",
        "introduction": "A new era in the study of Gamma-ray Bursts (GRBs) started with the \\textit{BeppoSAX} observations leading to the discovery of their X--ray, optical and radio afterglows \\citep{costa,vanpa,frail}. The great progress which occurred in the last few years in our understanding of GRBs has been possible thanks to extensive multi-wavelength observations of these unpredictable and rapidly fading events. In this respect, a quick derivation and distribution of accurate sky positions for GRBs is crucial. Here we review the contribution in this field obtained during the first 18 months of the INTEGRAL mission. We concentrate on the GRBs observed within the field of view of the IBIS instrument \\citep{ibis}. Bursts observed with the SPI Anticoincidence Shield (ACS) are described elsewhere in these proceedings \\citep{acs}.  Thanks to its 72 hours orbit, the INTEGRAL satellite is in continuous contact with the ground stations during the observations. This has allowed us to implement a ground-based software, the INTEGRAL Burst Alert System (IBAS), for the search in near real time of GRBs \\citep{ibas}. The IBAS software and its current performances are briefly described in Section 2. In Section 3 we summarize the main results on the twelve GRBs observed to date in the field of view of the INTEGRAL instruments. Finally, in Section 4 we describe the  IBAS capability  to provide real time information also on other classes of transient sources. ",
        "conclusions": ""
    },
    "0404/astro-ph0404369_arXiv.txt": {
        "abstract": "We present a search for water maser emission at 22 GHz associated with young low-mass protostars in six \\Htwo\\ regions --- M16, M20, NGC~2264, NGC~6357, S125, and S140.  The survey was conducted with the NRAO Very Large Array from 2000 to 2002. For several of these \\Htwo\\ regions, ours are the first high-resolution observations of water masers. We detected 16 water masers: eight in M16, four in M20, three in S140, and one in NGC~2264.  All but one of these were previously undetected.  No maser emission was detected from NGC~6357 or S125. There are two principle results to our study.  (1) The distribution of water masers in M16 and M20 does not appear to be random but instead is concentrated in a layer of compressed gas within a few tenths of a parsec of the ionization front.  (2) Significantly fewer masers are seen in the observed fields than expected based on other indications of ongoing star formation, indicating that the maser-exciting lifetime of protostars is much shorter in \\Htwo\\ regions than in regions of isolated star formation.  Both of these results confirm predictions of a scenario in which star formation is first triggered by shocks driven in advance of ionization fronts, and then truncated $\\sim 10^5$ years later when the region is overrun by the ionization front. ",
        "introduction": "The majority of low-mass stars form not in regions of isolated star formation, such as the Taurus-Auriga molecular cloud, but instead in cluster environments that also contain massive stars \\citep[see, e.g.][]{ll03}. The process of star formation may be fundamentally different in these two environments. In low-mass star-forming environments, each protostar is largely responsible for shaping its own natal environment.  In contrast, star-forming environments surrounding \\Htwo\\ regions are shaped by the radiation and mechanical energy of the O and B stars that excite the region.  Much effort has been made to understand the importance of supernova shocks as triggers of low-mass star formation \\citep[see, e.g.][]{fb96}.  But an even more important effect may involve triggering of star formation by radiatively driven implosion of molecular gas surrounding \\Htwo\\ regions \\citep[see, e.g.][]{el77,ber89,bm90}.  \\citet{hes96} pointed out that such regions would be quickly overrun by the advance of the ionization front, exposing cloud cores and protostars to the potentially disruptive environment in the interior of the \\Htwo\\ region.  Recent optical and near-infrared observations of young stellar objects at the edges of \\Htwo\\ regions \\citep[e.g.][]{ma02,osp02,sug02,tsh02} provide additional support for the triggering and truncation of the star formation process in these environments.  Although the stellar population in the low density interior of the \\Htwo\\ region can be studied in the optical and near-infrared part of the spectrum, the youngest protostars are extremely difficult to observe at these wavelengths.  This is due in part to the high foreground extinction associated with the molecular gas in which the protostars are embedded.  This is also due to the fact that protostars are intrinsically cool objects, radiating at centimeter wavelengths and peaking at submillimeter wavelengths \\citep{awb93}.  If we are to address properly the question of low-mass star formation in complex \\Htwo\\ region environments, then we need a way to probe the spatial distribution of young protostars, including low-mass protostars, in the high-density molecular gas bounding \\Htwo\\ regions.  In this paper we present observations of water masers in and around \\Htwo\\ regions, and find such observations to provide a powerful yet previously under-utilized tool with which to conduct such studies.  Previous observations of water maser emission from low-mass protostars have focused on the relatively nearby sample selected from low-mass star-forming regions like Taurus-Auriga.  For example, the sample of \\citet{cla96} was selected to have d $\\lesssim$ 450 pc and 90\\% of the sources of \\citet{fur01} had d $\\lesssim$ 500 pc.  These and other studies have provided a wealth of data on the behavior of water maser emission.  We would like to apply that knowledge to the more distant and more complex case of star formation and water maser activity in regions shaped by the presence of massive stars.  Water maser emission is only known to arise in association with AGN, AGB stars, and protostars.  For example, \\citet{cla99} find that even the most favorable conditions behind supernova remnant shocks fail to produce water maser activity.  Maser emission in the $6_{1,6}\\rightarrow5_{2,3}$ transition of water at 22.235077 GHz is a well-known trait of the youngest protostars, low- and high-mass alike \\citep{tof95,cla96,mee98,fur03}.  Observational studies of both massive and low-mass protostars suggest that maser activity is associated with the youngest stages of star formation. \\citet{fur01,fur03} show that $\\sim$40$\\pm$20\\% of Class 0 and $\\sim$4$\\pm$3\\% of Class I low-mass protostars exhibit water maser emission.  \\citet{fc00} conclude that water masers are likely excited by massive pre-main sequence stars or possibly actively-accreting, massive protostars.  Observations also show that water maser activity lies in close proximity to the exciting object.  Water masers are observed within tens or hundreds of AU of low-mass protostars \\citep[e.g.][]{che95,cla98,fur00}. The separation between maser and massive star can be $\\lesssim10^4$ AU \\citep[e.g.][]{tof95,fc00}. Only in extreme cases (e.g. the water masers in Orion), does the spatial correlation between water masers and the exciting protostar break down. In general, the presence of water masers in a region of star formation, such as the molecular gas surrounding an \\Htwo\\ region, is an unambiguous indication of the presence of a star in the earliest stages of formation. Due to the intrinsic variability and emission geometry of maser emission, this detection method will not uncover all young protostars.  Nonetheless, water maser emission can provide us with a look at the distribution of young protostars within an \\Htwo\\ region environment. Finally, while low- and high-mass protostars may result from the star formation around \\Htwo\\ regions, far more low-mass stars form in this process \\citep{ll03}. Statistically we expect that most of the observed water maser activity in \\Htwo\\ regions is excited by low-mass protostars.  The presence of water masers in \\Htwo\\ region environments is well-known \\citep[see, e.g.][]{gd77,bl79,com90}. These surveys and more recent searches have been performed with single-dish telescopes, which have typically 1 arcminute resolution.  This resolution is inadequate to properly locate protostars relative to ionization fronts and other details of the \\Htwo\\ region environment.  If we want to use water maser emission to study the process of star formation in regions around massive stars, we need positions accurate to 1\\arcsec\\ or better, comparable to optical and near-infrared observations of \\Htwo\\ regions.  The NRAO Very Large Array (VLA)\\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} is the instrument of choice for this type of investigation compared to single-dish telescopes. First, the greater effective aperture of the VLA provides greater sensitivity which allows for weaker masers to be detected in shorter observations compared to single-dish observations. Second, the large primary beam of the VLA allows for fewer pointings to cover the same area of sky compared to large aperture single-dish telescopes. Third, the small synthesized beam of the VLA gives a precise position for water masers that can be compared to optical or infrared observations of the \\Htwo\\ regions.  The observations, calibration, and analysis of this survey are discussed in $\\S2$.  The results of the observations for each \\Htwo\\ region are presented in $\\S3$. In $\\S4$, we discuss our results, and find that we are able to test two specific predictions concerning triggering and subsequent disruption of low-mass star formation in these environments.  A summary of our results and conclusions is presented in $\\S5$. ",
        "conclusions": "We have searched six \\Htwo\\ regions for water maser emission associated with low-mass star formation. We detected a total of 16 water maser systems, with only the emission from S140 IRS~1 reported previously. The distribution of water masers within M16 and M20 is not random. Instead, the masers are strongly concentrated in regions that have recently been compressed by ionization shocks, and will be overrun by ionization fronts on timescales comparable to the maser-exciting lifetimes of protostars.  We have also found that the fraction of protostars that excite masers is significantly lower in \\Htwo\\ region environments than in regions of isolated star formation.  The original motivation for this work was in part to test a scenario for star formation proposed by \\citet{hes96} in their work on M16. Radiative energy from massive O and B stars is deposited at the surface of adjacent molecular clouds, where it drives a stratified photoionized photoevaporative flow away from the interface \\citep[][and references therein]{sh00}, and drives an ionization shock into the molecular gas.  Ongoing star formation that is originally triggered by the ionization shock is disrupted a short time later by the advance of the ionization front.  Stars forming as a result of this process will pass through a well-defined series of stages, beginning as sub-Jeans cores, and ending as exposed young stellar objects, adrift within the ionized interior of the \\Htwo\\ region.  We found that the distribution of maser emission in \\Htwo\\ regions allowed us to test two specific predictions of this scenario for star formation: (1) low-mass star formation, if it is triggered by compression in advance of ionization fronts, should be localized near ionization fronts; and (2) low-mass stars should spend less time enshrouded in high density material in \\Htwo\\ region environments than they do in regions of isolated star formation.  The findings of the present study are in accord with these predictions.  Much work remains before we can claim to have sorted out the complex process of star formation in \\Htwo\\ region environments.  However, our data leave little doubt that pockets of star formation are triggered around \\Htwo\\ regions, and that the accretion phase of the star formation process proceeds differently in \\Htwo\\ region environments than in regions of isolated low-mass star formation.  Since most low mass stars form in \\Htwo\\ region environments, understanding the details of this process is of great importance in answering a diverse set of questions ranging from the origins of the IMF to the environments in which planet formation occurs."
    },
    "0404/astro-ph0404196_arXiv.txt": {
        "abstract": "Recent studies have shown that during their coalescence, binary supermassive black holes (SMBHs) experience a gravitational recoil with velocities of $100~{\\rm km~s^{-1}}\\lsim v_{\\rm kick} \\lsim 600~{\\rm km~s^{-1}}$.  These velocities exceed the escape velocity $v_{\\rm esc}$ from typical dark matter (DM) halos at high--redshift ($z\\gsim 6$), and therefore put constraints on scenarios in which early SMBHs grow at the centers of DM halos.  Here we quantify these constraints for the most distant known SMBHs, with inferred masses in excess of $10^9{\\rm M_\\odot}$, powering the bright quasars discovered in the Sloan Digital Sky Survey at $z>6$.  We assume that these SMBHs grew via a combination of accretion and mergers between pre--existing seed BHs in individual progenitor halos, and that mergers between progenitors with $v_{\\rm esc}< v_{\\rm kick}$ disrupt the BH growth process.  Our results suggest that under these assumptions, the $z\\sim 6$ SMBHs had a phase during which gained mass significantly more rapidly than under an Eddington--limited exponential growth rate. ",
        "introduction": "The gravitational waves (GWs) emitted during the final stages of the coalescence of two merging black holes (BHs) carry linear momentum, implying that the center of mass of the system experiences a recoil (Bonnor \\& Rotenberg 1961; Peres 1962).  The resulting recoil velocity had been known to be large, with values estimated to be of order $v_{\\rm kick} \\sim 1000~{\\rm km~s^{-1}}$ (Fitchett 1983). Favata et al. (2004) has recently revisited this problem and computed recoil velocities, treating the spin and orbital dynamics of the merging BHs, as well as the generation of GWs in the strong gravity regime. They have found the range of possible recoil velocities to be $100~{\\rm km~s^{-1}}\\lsim v_{\\rm kick} \\lsim 600~{\\rm km~s^{-1}}$, with the exact value depending on the mass ratio of the merging BHs, their spin, and orbital parameters.  These velocities are large compared to the escape velocities of dwarf galaxies, and of the typical dark matter halos that existed at the early epochs of galaxy formation ($z\\gsim 6$).  In hierarchical cosmogonies, the SMBHs that are known to exist in the local universe grew via a combination of accretion and mergers between holes residing in individual DM halos.  Merritt et al. (2004) and Madau \\& Quataert (2004) recently considered several consequences of a large recoil that removes a BH from the center of its host galaxy.  In particular, Merritt et al. (2004) pointed out that the ejection of SMBHs from the shallow potentials of DM halos at high redshift implies a maximum redshift at which the progenitors of present--day SMBHs could have started merging (and sticking) with each other.  In this {\\em Letter}, we consider the growth history of SMBHs that are in place at $z>6$, and are thought to power the bright quasars recently discovered (Fan et al. 2000; 2001; 2003) in the Sloan Digital Sky Survey (SDSS).  As discussed in Haiman \\& Loeb (2001; hereafter HL01), relatively little time is available for the growth of these few $\\times10^9~{\\rm M_\\odot}$ SMBHs prior to $z\\sim 6$, and their seed BHs must be present as early as $z\\sim 10$.  A model in which stellar seed BHs appear in small progenitor DM halos is consistent with the presence of a $\\sim 4\\times10^9~{\\rm M_\\odot}$ SMBH at $z\\sim 10$, provided that each seed BH can grow at least at the Eddington--limited exponential rate, and that the progenitor halos can form seed BHs sufficiently eary on.  The ejection of a merger--product SMBH from its host halo severely limits the ability of massive SMBHs to grow at $z>6$, by disrupting the early stages of growth.  In this {\\em Letter}, as an example, we model the growth of the SMBH powering the most distant SDSS quasar, SDSS 1054+1024 at redshift $z=6.43$, with an inferred BH mass of $\\sim 4\\times 10^9~{\\rm M_\\odot}$. Under the assumption that progenitor holes are ejected from DM halos with velocity dispersions $\\sigma<v_{\\rm kick}/2$, and do not contribute to the final BH mass, we find that typical recoil velocities must either be below the {\\it minimum} value $v_{\\rm kick}= 100~{\\rm km~s^{-1}}$ found by Favata et al., or else this SMBH must have had a phase during which it gained mass significantly more rapidly than the Eddington--limited exponential growth rate would imply.  The rest of this {\\it Letter} is organized as follows. In \\S~\\ref{sec:quasar}, we discuss the inferred values of the relevant parameters (halo and BH mass) of SDSS 1054+1024. In \\S~\\ref{sec:model}, we describe the method we use model the growth of the SMBH by accretion and mergers. In \\S~\\ref{sec:results}, we present our main result, showing that excluding seed BHs from low--mass progenitor halos necessitates a faster--than Eddington growth rate. In \\S~\\ref{sec:discuss}, we discuss various uncertainties about our results. In \\S~\\ref{sec:conclude}, we summarize the implications of this work and offer our conclusions.\\footnote{Throughout this paper, we adopt the background cosmological parameters for a flat universe as measured by the {\\it WMAP} experiment, $\\Omega_m=0.29$, $\\Omega_{\\Lambda}=0.71$, $\\Omega_b=0.047$, $h=0.72$, with a power spectrum normalization $\\sigma_{8h^{-1}}=0.9$ and slope $n=0.99$ (Spergel et al. 2003).} ",
        "conclusions": "\\label{sec:conclude}  The recent discovery of luminous quasars at redshift $z>6$ provides evidence that supermassive black holes (SMBH) as large as several $\\times10^9~{\\rm M_\\odot}$ were assembled during the first $\\lsim 10\\%$ of the current age of the universe.  In the context of hierarchical structure formation scenarios, these early SMBHs grow via accretion and mergers of seeds that appear at much earlier epochs, $z\\gsim 10$.  Unless the growth by accretion of individual seed BHs is significantly faster than Eddington--limited accretion at a fiducial radiative efficiency of $\\epsilon\\approx 10\\%$, we find that such scenarios appear inconsistent with the large recoil velocities, $100~{\\rm km~s^{-1}}\\lsim v_{\\rm kick} \\lsim 600~{\\rm km~s^{-1}}$, that were recently calculated to occur during the coalescence of massive BHs.  A natural resolution of this discrepancy would be for individual SMBHs to grow in mass at significantly super--Eddington rates. In this case, the SMBHs can arise from seeds that appear relatively more recently, (i.e. closer to redshift $z\\sim 6$, where the bright quasars exist), in DM halos whose potential wells are deep enough to retain these seeds despite their recoil.  Forthcoming instruments, such as the {\\it James Webb Space Telescope (JWST)} in the infrared, or deep radio instruments, such as the {\\em Allen Telescope Array (ATA)}, {\\em Extended Very Large Array (EVLA)}, and the {\\em Square Kilometer Array (SKA)}, whose sensitivities allow them to detect BHs as small as $10^5~{\\rm M_\\odot}$ at $z\\sim 10$ (see Haiman \\& Quataert 2004 for a recent review) would be able to shed light on whether or not the seeds of $z\\sim 6$ quasar black holes indeed extend out to these redshift."
    },
    "0404/astro-ph0404475_arXiv.txt": {
        "abstract": "I briefly report on the work conducted to probe the gravitational potential of active and non-active disk galaxies using gas and stellar kinematics. ",
        "introduction": "Bars and non-axisymmetric galactic structures imply time variable torques which can help transfering angular momentum. Large-scale bars are important actors in the spatial redistribution of the dissipative component in disk galaxies, and there is indeed a direct observed link between their presence and e.g. star formation or the gas concentration. However, as repeatedly emphasized by many authors (see e.g. Martini's contribution in these Proceedings), there is no significant correlation between the nuclear activity and the presence of such structures. A nice summary has been given by F. Combes (these Proceedings), but I should start by repeating a simple but important argument. ",
        "conclusions": "\\label{sec:concl}  I will conclude by reiterating on a few important issues. First stating once more that building a gas reservoir within the central 1~kpc or 100~pc, is a different thing than fuelling the central engine: spatial and time scales are important in this context. Hence we should make it clear when we write about ''fuelling the AGN''.  In this short paper, I focused on the accumulation of gas first within the central 1000 and 100~pc. Our \\Sauron\\ survey of active and non-active galaxies, supplemented by large-scale HI data, will provide us with a unique probe of the underlying gravitational potential in the inner kpc. Inside this region, inner bars are not rare and may indeed play an important role in accumulating mass within the central 100~pc. However, we need to map both gas and stellar kinematics in order to reveal the dynamical structures. We may then find that $\\sigma$--drops are long-lived signatures of accretion, but this will require to extend the available pool of data quite dramatically."
    },
    "0404/astro-ph0404104_arXiv.txt": {
        "abstract": "Using the Giant Metrewave Radio Telescope (GMRT), we have discovered PSR~J0514$-$4002A, a binary millisecond pulsar in the globular cluster NGC~1851. This pulsar has a rotational period of 4.99 ms and the most eccentric pulsar orbit yet found: $e = 0.89$. The orbital period is 18.8 days, and companion has a minimum mass of 0.9~M$_{\\odot}$; its nature is presently unclear. After accreting matter from a low-mass stellar companion, this pulsar exchanged it for its more massive present companion. This system presents the strongest evidence to date of such a process. ",
        "introduction": "Since 1987, several globular cluster surveys (see review by F. Camilo and papers by A. Possenti, S. Ransom and J. Hessels, these proceedings) have confirmed that most of the binary millisecond pulsars (MSPs) in globular clusters (GCs) have low-mass white dwarf companions and nearly circular orbits, as observed in the Galactic disk. This is an important confirmation of the evolutionary scenarios proposed by Alpar~et~al.~(1982) for the formation of MSPs.  In GCs, {\\em exchange encounters}, which only have a significant probability of occurring in dense stellar environments, occasionally exchange one of the components of a binary system with a typically more massive star.  The exchanges may occur during encounters with either other binaries or with isolated stars.  In GCs such encounters can place isolated neutron stars into binaries with a main sequence (MS) star which eventually evolves, ``recycles'' the neutron star, and finally forms a MSP$-$WD binary system.  Such a process explains the anomalously large numbers of MSPs in GCs (by mass) when compared to the Galaxy. If the companion to the low-mass MS star is a previously recycled neutron star, we observe ``irregular'' eclipsing binary pulsars (see review by P. Freire).  In this paper we present some preliminary results on a new, unique binary millisecond pulsar, PSR~J0514$-$4002A (henceforth NGC~1851A) in the globular cluster NGC~1851. This has been found in a new, sensitive 327-MHz survey of globular clusters carried out using the Giant Metrewave Radio Telescope (GMRT), at Khodad near Pune, India.   \\begin{figure} \\setlength{\\unitlength}{1in} \\begin{picture}(0,2.8) \\put(0.1,-2.2){\\special{psfile=./fig1.eps hscale=70 vscale=70 angle=0}} \\end{picture} \\caption{Discovery observation for NGC~1851A, in the globular cluster NGC~1851. Pulsed emission is persistent throughout the whole scan. The pulse profile is rather narrow (right), corresponding to the time resolution of the system for this DM.} \\end{figure} ",
        "conclusions": "We have discovered a remarkable 5-ms binary pulsar, the first to be found in the GC NGC~1851 and the first pulsar to be discovered with the GMRT. Its orbit is the most eccentric known for any system containing a pulsar, while its rotational period is much shorter than that of any other pulsar in an eccentric binary system.  This combination of characteristics indicates that, after becoming an MSP by accreting matter from a low-mass companion star, this neutron star has almost certainly exchanged it for its present, significantly more massive companion. This is the most effective way of forming a binary system containing a millisecond pulsar and a black hole, provided black holes exist in GCs. Follow-up studies of this object will allow us to determine the nature of the companion and hopefully measure the masses of both components of this binary system."
    },
    "0404/astro-ph0404332_arXiv.txt": {
        "abstract": "{The X-ray source IGR J16318$-$4848 was the first source discovered by INTEGRAL on 2003, January 29. We carried out optical and near-infrared (NIR) observations at the European Southern Observatory (ESO La Silla) in the course of a Target of Opportunity (ToO) programme. We discovered the optical counterpart and confirmed an already proposed NIR candidate. NIR spectroscopy revealed a large amount of emission lines, including forbidden iron lines and P-Cygni profiles. The spectral energy distribution of the source points towards a high luminosity and a high temperature, with an absorption greater than the interstellar absorption, but two orders of magnitude lower than the X-ray absorption. We show that the source is an High Mass X-ray binary (HMXB) at a distance between $\\sim 1$ and $\\sim 6$ kpc, the mass donor being an early-type star, probably a sgB[e] star, surrounded by a rich and absorbing circumstellar material. This would make the second High Mass X-ray Binary (HMXB) with a sgB[e] star after CI Cam, indicating that a new class of strongly absorbed X-ray binaries is being unveiled by INTEGRAL. }    \\addkeyword{stars: circumstellar matter} \\addkeyword{stars: emission-line} \\addkeyword{Be---X-rays: binaries} \\addkeyword{IGR J16318-4848}  \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404042_arXiv.txt": {
        "abstract": "{The Brackett decrement in the Class I source \\object{HH100 IR} has been observed and analyzed to set constraints on the origin of the IR HI emission in this young object. We have used both low resolution (R$\\sim$800) observations of the Brackett lines from \\brg\\, to Br24, and medium resolution (R$\\sim$9000) spectra of the \\brg, Br12 and Br13 lines. The dereddened fluxes indicates that the lines remain moderately thick up to high quantum numbers. Moreover, the profiles of the three lines observed in medium resolution are all broad and nearly symmetric, with a trend for the lines at high n-number to be narrower than the \\brg\\, line. With the assumption that the three lines have different optical depths and consequently trace zones at different physical depths, we interprete the observed profiles as evidence that the ionized gas velocity in the HI emitting region is increasing as we move outwards, as expected in an accelerating wind more than in an infalling gas.\\\\ We have modelled the observed line ratios and velocities with a simplified model for the HI excitation from a circumstellar gas with a velocity law $V=V_0+(V_{max}-V_0)(1-(r_i/r)^{\\alpha})$. Such a comparison indicates that the observations are consistent with the emission coming from a very compact region of 4-6 R$_{\\odot}$, where the gas has been already accelerated to velocities of the order of 200 \\kms, with an associated mass flow rate of the {\\it ionized} component of the order of 10$^{-7}$ M$_{\\odot}$\\,yr$^{-1}$. This implies that the observed lines should originate either from a stellar wind or from the inner part of a disk wind, providing that the disk inner truncation radius is close to the stellar surface. It is also expected that the gas ionization fraction is relatively high as testified by the high rate of ionized mass loss derived. Our analysis, however, does not resolve the problem of how to reproduce the observed symmetrical line profiles, which at present are apparently difficult to model by both wind and accretion models. This probably points to the fact that the real situation is more complicated than described in the simple model presented here.} ",
        "introduction": "The emission from hydrogen recombination lines represents the most direct manifestation of circumstellar activity in young stars. Lines from the Balmer series have been the main spectroscopic tool for identifying classical T Tauri and Herbig AeBe stars. In spite of the large amount of observational data, however, the HI lines are so easily excited that it is very difficult to clearly define their origin. Originally, they were interpreted as being excited in circumstellar ionized winds, on the basis of the often observed P Cygni profile exhibited by the H$\\alpha$ line (Hartmann et al. 1990, Calvet et al. 1992). More recently, such an interpretation has been challanged due to the variety of the observed Balmer line profiles (e.g. Edwards et al. 1994, Reipurth et al. 1996) that are difficult to reproduce by means of wind models, and the HI emission has been interpreted also in the framework of magnetospheric accretion models (Calvet \\& Hartman 1992, Muzerolle et al. 1998a). Studies of HI lines in pre-main sequence stars have been extended also in the near and mid IR (Folha \\& Emerson 2001, Benedettini et al. 1998, Nisini et al. 1995, Natta et al. 1988). While the IR HI line ratios are well reproduced by wind models, they apparently fail to explain the observed profiles of the \\brg\\, and Pa$\\beta$ lines. On the other hand, magnetospheric accretion models were able to reproduce the \\brg\\, profiles showing redshifted absorption components (Muzerolle et al. 1998b) but not all the variety of profiles exhibited by T Tauri stars (Folha \\& Emerson 2001). Finally, very recently, spectro-astrometric observations of the Pa$\\beta$ line in four T Tauri stars, have indicated that the emission originates in the outflowing material (Whelan et al. 2004).\\\\ While there is ample literature on the origin and observations of HI lines in T Tauri and Herbig stars, the interpretation of HI lines in more embedded Class I sources, for which optical observations of the Balmer series are impeded, is much less defined. Observations of spectrally-resolved HI lines in low-mass Class I sources so far have been limited to the Br$\\gamma$ line (Najita et al. 1996, Davis et al. 2001). At variance with H$\\alpha$ in T Tauri stars, the profiles of the \\brg\\, line appear more symmetric, with no evidence of the typical P Cygni signature. On this basis Najita et al. (1996) argue against a wind origin, while in the framework of magnetospheric accretion models, Muzerolle et al. (1998b) use the \\brg\\, luminosity in a sample of Class I sources to infer the mass accretion rates.\\\\ To better constrain the emission region and the mechanism for the excitation of HI lines in Class I sources it is necessary to obtain a larger observational set, including both information on line profiles and flux ratios among different lines. The aim of this work is to start such an analysis using both the ratios from lines in the Brackett series (the Brackett decrement) and high resolution spectroscopy of different Brackett lines. The source investigated here is \\object{HH100 IR}, a low mass ($L_{bol} \\sim$ 10 L$_{\\sun}$, Wilking et al. 1992)  Class I object located in the R CrA star forming core (D=130 pc, Marraco \\& Rydgren 1981). This source already has been recognized as an embedded protostar with high circumstellar activity on the basis of its near IR specrum (Greene \\& Lada 1996, Nisini et al. 2004), and as also suggested by the fact that it is the driving source of an Herbig Haro object. It therefore represents a suitable test case to perform a detailed analysis of its HI emission. ",
        "conclusions": "The analysis of the Brackett decrement in the \\object{HH100 IR} Class I source has shown that the observed Brackett lines remain optically thick up to a high quantum number. In addition, the profiles of the \\brg, Br 12 and Br 13 lines are fairly symmetric with a tendency of the line FWHM to decrease as the n-number increases. On the basis of this evidence we argue that a wind-like mechanism, where the gas is accelerating outwards is  more suited to reproduce the observed features. A comparison with a very simple model suggests that the emission region should be very compact and close to the stellar surface, a fact favouring a stellar wind or the inner region of a disk wind with a small truncation radius as the natural emission region. The mass loss rate of the ionized gas should be fairly high, of the order of $\\dot{M}_{ion} \\sim$ 10$^{-7}$M$_{\\sun}$\\,yr$^{-1}$, to maintain the lines optically thick. This suggests that the ionization fraction in the emitting gas should not be smaller that $\\sim$0.1 to have a total mass loss rate not exceeding typical values expected from Class I objects.\\\\ Our analysis, however, does not solve the problem of how to reproduce the observed symmetrical line profiles of optically thick IR lines, which at present are apparently difficult to modell by both wind and accretion models. This probably points to the fact that the real situation is more complicated than described in the simple model presented here. Another caveat to the derived conclusions is the near IR photometric variability of \\object{HH100 IR}, which shows differences in the $K$ band of more than a magnitude over a period of a few years. Such a variability can also affect the line emission of this source. Indeed, variability not only of HI NIR line fluxes but also of their profiles has been observed in other Class I sources exhibiting photometric variability (Nisini et al. 1994). \\\\ Finally,  our conclusions are valid for the specific case of HH100IR and as such cannot  be generalized to other Class I objects. Indeed, more asymmetric profiles of the \\brg\\, line in similar sources have been succesfully modelled by magnetospheric accretion models, such as e.g. WL16 (Muzerolle et al. 1998b). For a  better understanding of the origin of the HI lines in Class I sources, a more detailed modelling effort is needed, allowing us to fit in a consistent way both the intensity and profiles of the different observed lines."
    },
    "0404/astro-ph0404568_arXiv.txt": {
        "abstract": "The supernova physics may provide a clue of the cosmological dark matter. In the absence of new physics, the supernova calculations do not explain the observed velocities of pulsars.  However, if there exists a singlet fermion with mass in the 1--20~keV range and a small mixing with neutrinos, this particle could be emitted asymmetrically from a cooling neutron star in the event of a supernova explosion.  The asymmetry could explain the long-standing puzzle of pulsar velocities.  The same particle could be the dark matter.  Observations of X-ray telescopes, as well as a future detection of gravitational waves from a nearby supernova can confirm or rule out this possibility. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404274_arXiv.txt": {
        "abstract": "As discussed elsewhere in these proceedings, the double pulsar system is a magnificent laboratory for gravitational physics and for studying pulsar magnetospheres. Here we consider the uses of having two clocks in the system in the context of its origin and evolution. We find that the ``standard'' evolutionary scenario involving spin-up of the first-born neutron star in an X-ray binary phase is consistent with the observed parameters. Equality of the spin-down ages of the two pulsars requires that the post-accretion spin period of A most likely lies in the range $16 \\lapp P_{\\rm 0,A} \\lapp 21$ ms. The likely age of the system is $30-70$ Myr. ",
        "introduction": "An application of the standard binary pulsar evolutionary model (e.g.~Bhattacharya \\& van den Heuvel 1992) to the double pulsar system J0737--3039 identifies the 22.7-ms pulsar (hereafter A) as the first-born neutron star.  Following an initial phase where A existed as a regular radio pulsar with a main-sequence companion, the currently observed spin parameters of A are the result of a subsequent X-ray binary phase where it acquired matter and angular momentum from the secondary star after the secondary evolved off the main sequence and overflowed its Roche lobe.  By processes that are not fully understood, this mass-transfer phase also resulted in a reduction of A's magnetic field (Shibazaki et al.~1989).  Following the supernova explosion of the secondary, a second neutron star was formed which we now observe as the 2.77-s pulsar (hereafter B).  While the spin periods of both these pulsars, and the low inferred magnetic field of A relative to B\\footnote{Due to the interaction of A's wind which penetrates deep into B's magnetosphere, some care should be taken when interpreting B's magnetic field strength.  In spite of this disruption, B is observed to be spinning down due to a steady braking torque which we model in Section 2.} are consistent with this model, a further test is to compare the time since the spin-up phase of A ended (${\\cal T}_{\\rm A}$) with the time since B has been active as a radio pulsar (${\\cal T}_{\\rm B}$). We therefore expect ${\\cal T}_{\\rm A}={\\cal T}_{\\rm B}$.  The simplest means to test this prediction is to use the characteristic ages of A and B, which are based on the observed periods and period derivatives: $\\tau_{\\rm A}=P_{\\rm A}/(2 \\dot{P}_{\\rm A})$ and $\\tau_{\\rm B}=P_{\\rm B}/(2 \\dot{P}_{\\rm B})$. Lyne et al.~(2004) find $\\tau_{\\rm A}=2.1 \\times 10^8$ yr and $\\tau_{\\rm B}=0.5 \\times 10^8$ yr.  Possible explanations for this discrepancy are: (i) the standard evolutionary scenario does not apply; (ii) as observed in other pulsars (see e.g.~Kramer et al.~2003), characteristic ages are not reliable; (iii) both the model and the characteristic ages are wrong!  Given the aforementioned circumstantial evidence in favour of recycling hypothesis, the simplest solution is option (ii). We now briefly investigate the implications for this case. Further details will be given in a forthcoming paper (Lorimer et al.~in preparation). ",
        "conclusions": "The two clocks in the double pulsar system J0737--3039 provide a unique means to constrain the age and birth spin periods of the two pulsars. These simple case studies demonstrate the use of the recycling model to place constraints on the post-accretion spin period of A ($P_{\\rm 0,A}$) and the age since spin-up ceased (${\\cal T}$).  For the spin-down models considered, we find $P_{\\rm 0,A}$ to be in the range $16-21$ ms and ${\\cal T}$ to be in the range $30-70$ Myr.  This is shorter than the age of 100 Myr assumed in the merger rate calculations by Burgay et al.~(2003) and Kalogera et al.~(2004) and hence increases the overall neutron star merger rate by 20--40\\%."
    },
    "0404/astro-ph0404038_arXiv.txt": {
        "abstract": "We report on the results of a BeppoSAX observation of the dipping LMXRB EXO 0748-676 performed in 2000 November. This is the first simultaneous observation of this source over the 0.1-100 keV energy range. The persistent spectrum is complex and shows a soft excess, which requires the inclusion of an ionized absorber (with a sub-solar abundance of iron). A cutoff power-law is a good fit to the high energy part of the spectrum, with a photon index of 1.3 and a cutoff around 50 keV. The 0.1-100 keV luminosity is 8.2$\\times$10$^{36}$ erg~s$^{-1}$. ",
        "introduction": "EXO 0748-676 is a transient LMXRB discovered with EXOSAT (Parmar et al., 1986) with a 3.82~hr orbital period measured from its X-ray eclipses. It is a type I X-ray burster and displays quasi periodic oscillations (discovered with RossiXTE, Homan et al., 1999; Homan \\& van der Klis, 2000) as well as intensity dips and X--ray eclipses (Parmar et al. 1986; Church et al., 1998).  It shows a complex soft X-ray spectrum (Thomas et al., 1997). Church et al. (1998), analysing dip and non-dip ASCA spectra, interpreted the source spectrum as composed by an extended Accretion Disk Corona (ADC), producing the hard emission, plus a point-like soft black-body emission. The spectral evolution during dipping is explained with the  progressive covering of the ADC emission.  A different view on the system came from EPIC/XMM-Newton observations, which revealed a persistent (i.e., no dipping and no eclipsing) spectrum composed by an extended  thermal halo together with a highly absorbed, and more compact, high energy power-law produced in the ADC (Bonnet-Bidaud et al., 2001).  RGS/XMM-Newton observations detected absorption and emission lines from ionized Ne, O and N (Cottam et al.,2001).  The cumulative RGS spectra during type I X-ray bursts revealed iron absorption lines, from which a gravitational redshift of z=0.35 have been measured (Cottam et al., 2002).  Chandra observations confirmed the presence of photoionized plasma, probably located above the accretion disk. The Chandra spectrum during dips was absorbed both by neutral and ionized material (Jimenez-Garate et al., 2003).  A recent analysis of several XMM-Newton observations revealed for the first time the clear detection of eclipses below 2 keV (Homan et al. 2003). ",
        "conclusions": "We have discussed here the first  observation of the broad-band X--ray emission of the dipping LMXRB EXO~0748-676 showing that this source displays significant emission up to 100~keV.  The best-fit model to the ``persistent\" spectrum consists of a cut-off power-law, extra-absorbed at low energies by an ionized absorber with non-solar iron abundance. The need for an ionized absorber  indicates that absorbing matter local to the binary system is located over most or all of the orbital period."
    },
    "0404/astro-ph0404512_arXiv.txt": {
        "abstract": "Some active galactic nuclei, microquasars, and gamma ray bursts may be powered by the electromagnetic braking of a rapidly rotating black hole. We investigate this possibility via axisymmetric numerical simulations of a black hole surrounded by a magnetized plasma.  The plasma is described by the equations of general relativistic magnetohydrodynamics, and the effects of radiation are neglected. The evolution is followed for $2000 G M/c^3$, and the computational domain extends from inside the event horizon to typically $40 G M/c^2$.  We compare our results to two analytic steady state models, including the force-free magnetosphere of Blandford \\& Znajek.  Along the way we present a self-contained rederivation of the Blandford-Znajek model in Kerr-Schild (horizon penetrating) coordinates.  We find that (1) low density polar regions of the numerical models agree well with the Blandford-Znajek model; (2) many of our models have an outward Poynting flux on the horizon in the Kerr-Schild frame; (3) none of our models have a net outward energy flux on the horizon; and (4) one of our models, in which the initial disk has net magnetic flux, shows a net outward angular momentum flux on the horizon. We conclude with a discussion of the limitations of our model, astrophysical implications, and problems to be addressed by future numerical experiments. ",
        "introduction": "A black hole of mass $M$ and angular momentum $J = a G M/c$, $0 \\le a/M < 1$ has a free energy associated with its angular momentum (or ``spin'').  This energy can, in principle, be tapped by manipulating particle orbits so that negative energy particles are accreted \\citep{pen69}.  Spin energy can also be tapped by superradiant scattering of vacuum electromagnetic waves \\citep{pt72}, gravity waves \\citep{hh72,tp74}, or magnetohydrodynamic (MHD) waves \\citep{u97}.  It can also be tapped through the action of force-free electromagnetic fields \\citep{bz77}.  The Blandford-Znajek (BZ) effect-- broadly used here to mean the extraction of energy from rotating holes via a magnetized plasma-- appears to be the most astrophysically plausible exploitation of black hole spin energy.  Relativistic jets in active galactic nuclei, galactic microquasars, and gamma-ray bursts (GRBs) may well be powered by the BZ effect. Despite some hints (e.g. \\citealt{wilms2001}, \\citealt{miller2002}, \\citealt{mt03}) and the general consistency of this idea with the data, however, there is no {\\it direct} observational evidence for black hole energy extraction.  In this paper we take an experimental approach and study the BZ effect through direct numerical simulation of a magnetized plasma accreting onto a black hole.  The energy stored in black hole spin is potentially large.  If $M_{irr}$ is the ``irreducible mass'' of the black hole where, in units such that $G = c = 1$, \\begin{equation} M_{irr}^2 =  {1\\over{2}} M r_+ , \\end{equation} and $r_+ = M (1 + \\sqrt{1 - (a/M)^2})$ is the horizon radius, then the free energy is \\begin{equation} E_{spin} = M - M_{irr} < 5.3 \\times 10^{61} \\left({M \\over{10^8 \\msun}}\\right) {\\rm erg}. \\end{equation} or $\\approx 30\\%$ of the gravitational mass of a maximally rotating hole.  This corresponds to a luminosity of $\\lesssim 4 \\times 10^{10} (M/10^8\\msun) \\lsun $ if released over a Hubble time.  Estimates suggest that black hole accretion is surprisingly efficient, in the sense that the ratio of quasar radiative energy density to supermassive black hole mass density is $\\sim 0.2$ \\citep{yt02,erz02}. During the accretion process some mass-energy is radiated away and the rest is incorporated into the black hole. Through electromagnetic spindown this energy gets a second chance to escape.  A combination of efficient thin disk accretion (in which all radiation is somehow permitted to escape) followed by the Penrose process can in principle extract up to $(1 - 1/\\sqrt{6}) c^2 = 0.59 c^2$ per gram of accreted rest-mass.  In practice, of course, much less energy is likely to be available. One goal of our investigation is to discover how much less. Part of the answer may lie with the calculations already described in \\cite{gsm2004}: if black hole spins are limited by the equilibrium value found there ($a/M \\approx 0.92$) then the nominal thin disk efficiency of the accretion phase is about $\\approx 17\\%$, much less than the 42\\% expected at $a/M=1$.  In this paper we consider the self-consistent evolution of a weakly magnetized torus surrounding a rotating black hole.  The evolution is carried out numerically in the axisymmetric ideal MHD approximation.  As the evolution progresses the computational domain develops matter dominated regions near the equator and electromagnetic field dominated regions near the poles.  To fix expectations for the structure of these regions we review two analytic models for the interaction of a magnetized plasma with a black hole in \\S~\\ref{analyticmodels}.  Along the way we develop the relevant notation and coordinate systems.  In \\S~\\ref{numericalmodels} we describe our numerical model and give a summary of numerical results for a high resolution fiducial model. In \\S~\\ref{parameterstudy} we consider the dependence of our results on model parameters. A discussion and summary may be found in \\S~\\ref{bzdiscussion}. From here on we adopt units such that $G M = c = 1$. Table~\\ref{tbl1} gives a list of commonly used symbols.  \\begin{deluxetable}{lll} \\tablewidth{0pt}  \\tablecaption{Commonly used symbols \\label{tbl1}}  \\tablehead{  \\colhead{Symbol}  & \\colhead{Fiducial Value}  & \\colhead{Description}  }   \\startdata & \\underline{Model Parameters} & \\\\ $a$ & $0.938$ & black hole spin ($J/M^2$) \\\\ $r_+$ &  $1.347$ & radius of the event horizon ($r_+=1+\\sqrt{1-a^2}$) \\\\ $r_{isco}$ &  $2.044$ & radius of the ISCO (innermost stable circular orbit) \\\\ $r_{edge}$ &  $6$ & radius of inner edge of torus\\\\ $r_{max}$ &  $12$ & radius of the pressure maximum \\\\ $\\Omega_H$ &  $\\approx 0.3477$ & spin frequency of zero angular momentum observer at $r_+$ \\\\ $R_{in}$ &  $0.98 r_+$ & inner radial grid location \\\\ $R_{out}$ &  $40$ & outer radial grid location \\\\ $\\beta$ & $100$ & ratio of gas to magnetic pressure (initially $\\frac{p_{gas,max}}{p_{mag,max}}$) \\\\ $\\gamma$ & $4/3$ & $p_{gas}=(\\gamma-1)u$ \\\\ & &\\\\ & \\underline{Diagnostics} & \\\\ $\\mflux$ & see sections~\\ref{goveqns}     \\& \\ref{numericalmodels} & rest-mass flux into the black hole \\\\  $\\eflux$ & see sections~\\ref{goveqns}     \\& \\ref{numericalmodels} & energy flux into the black hole \\\\ $\\efluxem$ & see sections~\\ref{goveqns}   \\& \\ref{numericalmodels} & electromagnetic energy flux \\\\ $\\efluxma$ & see sections~\\ref{goveqns}   \\& \\ref{numericalmodels} & matter energy flux \\\\  $\\lflux$ & see sections~\\ref{goveqns}     \\& \\ref{numericalmodels} & angular momentum flux into the black hole \\\\ $\\lfluxem$ & see sections~\\ref{goveqns}   \\& \\ref{numericalmodels} & electromagnetic angular momentum flux \\\\ $\\lfluxma$ & see sections~\\ref{goveqns}   \\& \\ref{numericalmodels} & matter angular momentum flux \\\\ $\\tilde{L}$ & see sections~\\ref{fiducial}  \\& \\ref{parameterstudy} & $\\tilde{L}=\\efluxem/(-\\eps\\mflux)$ ; $\\eps = 1 - \\eflux/\\mflux$\\\\ & &\\\\ & \\underline{Variables} & \\\\ $b^2/2$ &  see section~\\ref{inflowcomp} & electromagnetic energy density in the fluid frame\\\\ $B^r$,$B^\\theta$,$B^\\phi$ & see section~\\ref{goveqns} & magnetic field components. $B^i = \\dF^{it}$ \\\\ $A_\\phi$ & see section~\\ref{numericalmodels} & azimuthal component of electromagnetic vector potential \\\\ $\\tilde{v}^r$ & see section~\\ref{fiducial}   & asymptotic radial velocity (i.e. $v^r$ at $r=\\infty$) \\\\ $\\omega$ &  see sections~\\ref{bzcomp} \\& ~\\ref{inflowcomp} & spin frequency of electromagnetic field \\\\ $\\Omega$ &  see section~\\ref{inflowcomp} & spin frequency of fluid ($\\Omega=u^\\phi/u^t$)\\\\  \\enddata \\end{deluxetable} ",
        "conclusions": "In Kerr-Schild coordinates, then, the magnetic field components are \\begin{equation}\\label{BZBr} B^r = {C\\over{r^2}} + a^2 \\frac{C}{2r^4}\\left(-2\\cos{\\theta} + r^2(1+3\\cos{2\\theta}) f(r)\\right), \\end{equation} \\begin{equation}\\label{BZBh} B^{\\theta} = -a^2 \\frac{C}{r^2}\\cos{\\theta}\\sin{\\theta} f', \\end{equation} both accurate through second order in $a$, and \\begin{equation} B^{\\phi} = -a {C\\over{8 r^2}}(1 + {4\\over{r}}), \\end{equation} accurate through first order in $a$.  In Boyer-Lindquist coordinates, \\begin{equation} B^r[\\BL] = B^r[\\KS], \\end{equation} \\begin{equation} B^\\theta[\\BL] = B^\\theta[\\KS], \\end{equation} \\begin{equation} B^\\phi[\\BL] = B^\\phi[\\KS] - B^r[\\KS] {(a - 2 r\\omega)\\over{\\Delta}}, \\end{equation} and BZ's toroidal field \\begin{equation} B_T = \\Delta \\sin^2\\theta B^\\phi[\\BL] \\end{equation} (which is different from BZ's $B_\\phi$).  There has been some concern about causality in the application of the force-free approximation (e.g. \\citealt{pun2003}, see also \\citealt{kom02,kom04a}).  The MHD equations are hyperbolic and causal (as are the equations of force-free electrodynamics).  Below we show that a numerical evolution of the MHD equations agrees well with the BZ solution in those regions where $b^2/\\rho_0 \\gg 1$. This is either a remarkable coincidence or else the BZ solution is an accurate representation of the strong-field limit of ideal MHD.  For comparison with computational models, the most relevant aspects of the BZ theory are that: (1) the field is force-free; (2) the field rotation frequency $\\omega = a/8 + \\order{(a^3)}$ in the monopole geometry case and $\\omega = a/8 + \\order{(a^3)}$ at the poles ($\\theta = 0, \\pi/2$) in the paraboloidal field case considered by BZ;\\footnote{ According to the numerical results of \\cite{kom01} and the argument of \\cite{mt82}, $\\omega$ adjusts to $\\approx \\Omega_H/2$ hole even at large $a$.} (3) if the field geometry is nearly monopolar and $a$ is small enough that the expansion to lowest order in $a$ is accurate, then $B^r(\\theta)$ is given by equation (\\ref{BZBr}); and (4) if the field geometry is monopolar and $a$ is small, then the energy flux density $\\deflux \\propto \\sin^2\\theta$ on the horizon.  We compare this analytic BZ model to our numerical models in section~\\ref{bzcomp}.  \\subsection{Equatorial MHD Inflow}\\label{inflowanalytic}  \\cite{gam99} considered a stationary, axisymmetric MHD inflow in the ``plunging region'', between the innermost stable circular orbit (ISCO) and the event horizon.  The flow was assumed to be cold (zero pressure), nearly equatorial, and to proceed along lines of constant latitude $\\theta$.  The latter assumption ignores the requirement of cross-field force-balance.  This model is analogous to the \\cite{wd67} model for the solar wind, only turned inside out so that the wind flows from the disk into the black hole.  The model builds on earlier work by \\cite{tak90}, \\cite{phi83}, and \\cite{cam86}. The analytic model derived here will be used to compare to numerical models in section~\\ref{inflowcomp}.  The MHD inflow model is stationary ($\\del_t \\rightarrow 0$), axisymmetric ($\\del_\\phi \\rightarrow 0$) and nearly equatorial ($\\theta \\approx \\pi/2$) so $\\del_\\theta \\rightarrow 0$ by symmetry.  In addition flow proceeds along lines of constant $\\theta$.  As a result the model is one dimensional with a single independent variable $r$.  The nontrivial dependent variables are the radial and azimuthal four-velocity $u^r$ and $u^\\phi$, the radial and azimuthal magnetic field $B^r$ and $B^\\phi$, and the rest-mass density $\\rho_0$.  With these assumptions the equations of general relativistic MHD can be integrated completely.  The constancy of energy flux \\begin{equation} -\\detg {T^r}_t = const., \\end{equation} and angular momentum flux \\begin{equation} \\detg {T^r}_\\phi = const., \\end{equation} follow from ${T^\\mu}_{\\nu;\\mu} = 0$.  The source-free Maxwell equations imply \\begin{equation} \\detg B^r = const., \\end{equation} which expresses the constraint $\\bnabla \\cdot \\bB = 0$, and the relativistic ``isorotation law'', \\begin{equation} \\detg~\\dF^{r\\phi} = \\detg (u^r b^\\phi - u^\\phi b^r) = const. \\end{equation} where $b^\\mu$ is the magnetic field four-vector (defined above). Finally, conservation of particle number implies \\begin{equation} \\detg \\rho_0 u^r = const. \\end{equation} These five constants yield five constraints on the five nontrivial fundamental variables $u^r$, $u^\\phi$, $B^r$, $B^\\phi$, and $\\rho_0$. Given the constants, and using the constitutive relations that relate the constants and fundamental variables, one can solve the resulting set of nonlinear equations for the fundamental variables at each radius.  The next step is to determine the constants.  The radial magnetic flux and the rest-mass flux are determined by conditions in the disk and can be left as free parameters.  The remaining three degrees of freedom are fixed by imposing boundary conditions. \\cite{gam99} imposed the following conditions: (1) the flow is regular at the fast point (the flow is automatically regular at the Alfv\\'en point-- see \\cite{phi83} for a discussion-- and the slow point is absent because the flow is cold) ; and (2,3) the four-velocity components $u^r$ and $u^\\phi$ match onto a cold disk at the ISCO.  Energy can be extracted from the black hole if the Alfv\\'en point lies inside the ergosphere \\citep{tak90}.  \\cite{gam99} calculated $\\eflux$ and $\\lflux$ as a function of $a$ and $B^r$ and showed that for even modest magnetic field strength these were modified from the values anticipated in classical thin disk theory.  The implications of these modified fluxes for the structure-- particularly the surface brightness-- of a thin disk were explored by \\cite{ak00}.  For comparison with numerical models, the key predictions of the inflow model are: (1) the constancy of the conserved quantities with radius; (2) matching of the flow velocity to circular orbits at the ISCO; (3) modification of the angular momentum and energy fluxes from their thin disk values; and (4) the run of all the fluid variables with radius in the plunging region.  \\label{bzdiscussion}  We have used a general relativistic MHD code, HARM, to evolve a weakly magnetized thick disk around a Kerr black hole.  Our main result is that we find an outward electromagnetic energy flux on the event horizon, as anticipated by \\cite{bz77}.  The funnel region near the polar axis of the black hole is consistent with the Blandford-Znajek model.  The outward electromagnetic energy flux is, however, overwhelmed by the inward flux of energy associated with the rest-mass and internal energy of the accreting plasma. This result essentially confirms work by \\citet{ga97,lop99} that suggested the BZ luminosity should be small or comparable to the nominal accretion luminosity ($\\tilde{L}\\lesssim 1$).  One of our models discussed here, however, begins with a vertical field threading the torus, exhibits a brief episode of outward {\\it net} energy flux.  This appears to be a transient associated with the initial conditions.  The same model exhibits a steady {\\it net} outflow of angular momentum from the black hole.  Of all our models, the vertical field model has the largest negative $-\\efluxem/\\efluxma\\approx 15\\%$ (ratio of the electromagnetic energy flux to ingoing matter energy flux) and largest $\\tilde{L}=\\efluxem/(-\\eps\\mflux)\\approx 80\\%$ (ratio of electromagnetic luminosity to nominal accretion luminosity). This suggests that the BZ effect could play a significant role if the disk has a net dipole moment and accumulates magnetic flux that crosses the horizon.  This possibility will be considered in future work.  Consistent with the results found earlier by \\cite{dhk03}, we find that our models can be divided into four regions: (1) a ``funnel'' region with $b^2/\\rho_0 \\gtrsim 1$ and $\\beta \\ll 1$; (2) a corona with $1 \\lesssim \\beta \\lesssim 3$; (3); an equatorial disk with $\\beta > 3$; and (4) a plunging region between the disk and event horizon with $\\beta \\sim 1$ and a nearly laminar inflow from the disk to the black hole.  We also find no feature in the surface density at or near the ISCO (see figure~\\ref{inflowcompplot}), which agrees with the results by \\citet{kh02,dhk03} and consistent with \\cite{rb97}. This is contrary to the sharp transition predicted by thin disk models and used by XSPEC to fit X-ray spectra.  We have shown that the funnel region is nearly force-free, and is well described by the stationary force-free magnetosphere model of \\cite{bz77}, for which we have presented a self-contained derivation in Kerr-Schild coordinates.  We find agreement between the BZ model and our simulations in measurements of energy flux, magnetic field, efficiency of accretion, and spindown power output.  In all cases we find that in the force-free region the field rotation frequency is about half the black hole spin frequency, $\\Omega_H \\equiv a/(2 r_+)$. This spin frequency maximizes electromagnetic energy output from the hole.  This result is consistent with expectations of \\cite{mt82} and the force-free numerical results of \\cite{kom01}.  We have also compared the time-average of the plunging region in our fiducial model with the stationary MHD inflow model of \\cite{gam99}, which assumes that the flow matches a cold disk at the ISCO.  The inflow model matches the simulated rest-mass flux and electromagnetic flux of energy and angular momentum surprisingly well, particularly considering the strongly variable nature of the simulated flow in the plunging region. The inflow model fails to match other aspects of the flow, such as the radial component of the four-velocity.  This is mainly due to the finite temperature of the simulated flow; the inflow solution assumes zero temperature.  It is slightly surprising that the total angular momentum flux is close to the value predicted by the zero temperature inflow solution, and $20\\%$ less than what is predicted by the thin disk, yet the total energy flux is almost exactly what is predicted by the thin disk.  It is as yet unclear whether this is due to coincidence or conspiracy.  For a set of models similar to the fiducial model, the ratio of electromagnetic to matter energy fluxes is sensitive to the black hole spin, reaching $-7\\%$ for $a \\sim 1$. The evolution is sensitive to the initial field geometry.  Models with a net vertical field are more efficient, and more electromagnetically active than models with comparable field strength but zero net vertical field. Our models have a weak dependence on resolution in the sense that as resolution increases the relative importance of electromagnetic energy fluxes on the horizon diminishes.  With an $R_{out}=400$ model we demonstrate that an outgoing electromagnetic energy flux can reach large radii.  The field in the funnel region does not connect back into the disk.  Rather the poloidal components lie parallel to the polar axis.  The field lines are collimated by a factor of $5/2$ at $r = 40$ and by a factor of $10$ at $r = 400$.  An outflow along the boundaries of the funnel reaches a maximum $\\tilde{v}^r\\approx 0.75c$, but this is sensitive to the value of our artificial density ``floor'': a model with lower density reaches even larger radial velocities at the outer boundary of the computational domain.  \\citet{koide2002} have evolved a cold, highly magnetized uniform density plasma ($\\rho_0/p=0.06$, $b^2/\\rho_0=10$) as it falls into a rapidly spinning ($a=0.99995$) black hole in Boyer-Lindquist coordinates for a time $\\approx 14 G M/c^3$ using the MHD approximation. This initial state does not correspond to an accretion disk system. They demonstrated, however, that a transient {\\it net} energy extraction is possible from a spinning black hole. Because of the short evolution time they are unable to say whether the energy extraction process is possible in steady state. \\citet{koide2003} gives an expanded discussion of the above system.  In contrast to the results of \\citet{koide2002} and \\citet{koide2003}, we model a disk with an initially hydrodynamic equilibrium fluid that is weakly magnetized.  We also use a Kerr-Schild (horizon penetrating) coordinate system that avoids potential problems associated with the treatment of inner boundary condition in Boyer-Lindquist coordinates. In our simulation the Balbus-Hawley instability drives turbulence and accretion in a steady state where we evolve for a time $2000GM/c^3$.  We measure a sustained outward electromagnetic energy flux that is smaller than the inward matter energy flux (i.e. net {\\it inward} energy flux). Their model is evolved for too short a time to observe the unbound mass outflow in the funnel region as seen by us and \\cite{dhk03,dh04}.  \\citet{dh04}(hereafter DH) have also considered the numerical evolution of weakly magnetized tori around rotating black holes.  Their models are quite similar to ours in many respects, although they differ in that: (1) their models are three dimensional while our models are axisymmetric; (2) they use a nonconservative numerical method \\cite{dh03}; (3) DH use Boyer-Lindquist while we use Kerr-Schild coordinates; (4) DH choose $\\gamma = 5/3$ while we use $\\gamma = 4/3$; (5) DH's initial pressure maximum is located at $25 M$, while ours are typically at $12 M$.  Our results for the energy and angular momentum per baryon accreted from table~\\ref{tbl2} can be compared to Table 1 of DH by computing $\\eflux/\\mflux=\\Delta E_i/\\Delta M_i$ and $\\lflux/\\mflux=\\Delta L_i/\\Delta M_i$. For models with $a = (0, 0.5, 0.9)$ DH find $\\eflux/\\mflux=(0.91,0.91,0.84)$ while we find $\\eflux/\\mflux = (0.96,0.93,0.88)$.  For the same models DH find $\\lflux/\\mflux=(3.1,2.6,1.9)$, while we find $\\lflux/\\mflux = (3.1,2.6,1.7)$.  Given the differences in the models and numerical methods, this quantitative agreement is remarkable.  Our models and De Villier and Hawley's models also agree qualitatively in the sense that both show a similar geometry of disk, corona, and funnel and both imply that spin equilibrium is achieved at $a \\sim 0.9$ (see \\citealt{gsm2004}).  \\cite{kom01} finds the BZ solution to be stable in force-free electrodynamics, and \\cite{kom02,kom04a} find the BZ solution to be causal, but inconsistent with the membrane paradigm.  We find our numerical solutions to be consistent with the BZ solution in the low-density funnel region around the black hole. A numerical general relativistic MHD study of strongly magnetized (monopole magnetic field) accretion by \\cite{kom04b} is also consistent with the BZ solution.  For the strong field chosen he finds a considerably faster outflow (Lorentz factors of $\\approx 14$) than found in our models (Lorentz factors of $\\approx 1.5-3.0$).  Komissarov's model does not contain a disk.  The limitations of the numerical models presented here include the assumption of axisymmetry and a nonradiative gas.  The effect of axisymmetry can be tested by comparing our models with the three dimensional models of \\cite{dh04}; the angular momentum and energy per accreted baryon in the two models differs by only a few percent.  In addition the jet structure observed in \\cite{dh04} is nearly axisymmetric.  This is encouraging, although it is unlikely that an axisymmetric calculation can capture the full range of possible dynamical behavior in the accretion flow.  The radiation field, which we have completely neglected here, is likely to play a significant role in the flow dynamics, through radiation force on the outflowing plasma in the wind and through photon bubbles in the disk \\cite{gam98,socrates02}.  It will also, of course, play a significant role in heating and cooling the plasma.  This is clearly the most significant limitation of our calculation-- particularly from the standpoint of comparison with observations-- and clearly the most numerically difficult problem to overcome."
    },
    "0404/astro-ph0404348_arXiv.txt": {
        "abstract": "Recent deep X-ray surveys combined with spectroscopic identification of the sources have allowed the determination of the rest-frame $2-8~keV$ luminosity as a function of redshift.  In addition, an analysis of the $HEAO1~A2~2-10~keV$ full-sky map of the X-ray background (XRB) reveals clustering on the scale of several degrees. Combining these two results in the context of the currently favored $\\Lambda CDM$ cosmological model implies an average X-ray bias factor, $b_x$, of $b_x^2 = 1.12 \\pm 0.33$, i.e., $b_x = 1.06 \\pm 0.16$. These error estimates include only statistical error; the systematic error sources, while comparable, appear to be sub-dominant. This result is in contrast to the large biases of some previous estimates and is more in line with current estimates of the optical bias of $L_*$ galaxies. ",
        "introduction": "An important test of any cosmological model is that it be consistent with the observed distribution of matter in the universe.  Since our primary knowledge of this distribution comes from observations of galaxies, it is essential to understand the extent to which galaxies trace the matter density. This relationship is usually quantified by a bias factor which relates fluctuations in the galaxies to those in the dark matter.  It is complicated by the fact that the relation between the luminosity of a galaxy or groups of galaxies and the underlying distribution of matter can depend on the type of galaxy, the spectral band of the observation, the redshift $z$, and the scale length on which the comparison is made.  However, such complications are also opportunities in that models of galaxy formation must successfully reproduce these differences.  The standard definition of the bias factor, $b$, is the ratio of the fractional galaxy density fluctuations to the fractional matter density fluctuations, i.e., \\begin{equation} b = {\\delta \\rho_g / \\rho_g \\over \\delta \\rho / \\rho} \\end{equation} where $\\rho_g$ is the mean density of galaxies, $\\rho$ is the mean density of matter and $\\delta$ indicates the $rms$ fluctuations of the densities about these means. If galaxies formed early ($z >  1$), as appears to be the case (e.g., Ellis 1997), then there are good reasons to expect that, for linear density perturbations (i.e., $\\delta \\rho / \\rho \\ll 1$) on large scales in the nearby ($z \\ls 1$) universe, galaxies should be relatively unbiased ($b \\rightarrow 1$) tracers of the density field (Fry 1996, Tegmark \\& Peebles 1998); however, this assertion must be tested.  Here we focus on determining the bias of the hard X-ray background (XRB), which is known to be dominated by distant ($z \\ls 2$) active extragalactic galaxies and so provides a probe of the bias on large scales. The observed clustering of the XRB, when combined with what is known about the level of perturbations and the cosmological model from CMB observations, allows us to place relatively strong constraints on the X-ray bias.  Previous determinations of X-ray bias have resulted in a wide range of values, $1 < b_x < 7$ (see Barcons et al. 2000 and references therein).  Some spread in the estimates is to be expected; e.g., at lower energies X-ray emission is dominated by clusters of galaxies, and so are expected to be as highly biased as clusters themselves (Bahcall \\& Soneira 1983). However, another major contribution to the uncertainty in the bias estimates is the lack of accurate determinations of the clustering of various X-ray sources. Two of the lower estimates of X-ray bias are from Treyer et al. (1998) who found that $b_x \\sim 1$ to $2$ from a low order multipole analysis of the HEAO1 A2 data set and Carrera et al. (1998) who found that the ratio of the X-ray bias of AGN to that of IRAS galaxies to be $0.8 \\ls b_x/b_I \\ls 1.7$ from ROSAT observations.  The remaining uncertainty in these determinations arose from uncertainties in both the X-ray luminosity function (LF) and in the cosmological model.  Knowledge of both of these has improved dramatically in the last year and this is largely responsible for the improved accuracy of the estimate of $b_x$ in this paper. ",
        "conclusions": "We have determined the X-ray bias of the hard XRB assuming it is time (i.e., redshift) and scale independent.  These assumptions are probably quite reasonable since the mean redshift weighting of the X-ray ACF is quite low, $z \\sim 0.1$, and the linear scales probed by the ACF are quite large ($10~Mpc$ to $200~Mpc$).  Even if these assumptions are violated to some extent, $b_x$ can still be interpreted as an `average' X-ray bias.  There are several types of sources that contribute to the XRB, including quasars, Seyfert galaxies, LINERS, and clusters of galaxies, and the implied value of the bias must be considered to be an average over all these sources. However, the dominant contribution to the XRB is most likely to be moderately active AGN (Cowie et al. 2003), so $b_x$ should be representative of the bulk of the sources of the XRB.  With these caveats in mind, we find an X-ray bias of $b_x^2 = 1.12 \\pm 0.33$, i.e., $b_x = 1.06 \\pm 0.16$ (statistical error only). This error includes photon shot noise, fluctuations in the XRB from beam smearing, and the clustering of the XRB itself.  The fits of $b_x$ for two extreme models of $dI/dz$ indicate that the uncertainty due to our ignorance of the X-ray luminosity function is likely less than the statistical error.  Other possible sources of systematic error also seem small. We conclude that the hard XRB is a largely unbiased tracer of the matter distribution on large scales.  This is consistent with current models of large-scale, late time galaxy biasing (Benson et al. 2000; Tegmark \\& Peebles 1998). In addition, the latest studies of the clustering of $\\sim L_*$ galaxies on $\\sim 100~Mpc$ scales indicates that these objects are also unbiased tracers of matter.  Verde et al. (2002) found that, on scales of $\\sim 7$ to $\\sim 40~Mpc$, $b = 1.04 \\pm 0.11$ for $1.9~L_*$ galaxies in the $2dF$ survey with a mean redshift of $z = 0.17$.  Using a different analysis of the same data, Lahav et al. (2002) found that $b = 1.20 \\pm 0.11$ on scales of $\\sim 20$ to $\\sim 150~Mpc$. Both of these results are consistent with early findings from the SDSS and 2MASS surveys that imply linear bias factors on the order of unity (Tegmark et al. 2002; Miller et al. 2003).  It should not be surprising that the XRB and galaxy biases are similar since $L_*$ galaxies are closely associated with the moderately active AGN that comprise the bulk of the hard XRB (e.g., Barger et al. 2003; Miller et al. 2003).  If these estimates are accurate, then the X-ray bias factor in the linear regime is now much better determined. The hard XRB background appears to be an excellent tracer of the large-scale distribution of matter, making it a useful tool for understanding the evolution of structure in the universe. One example of the importance of determining galaxy biases (and indeed the driving motivation for this work) is to aid in the interpretation of recent detections of correlations of galaxies with the cosmic microwave background (CMB). We (Boughn \\& Crittenden 2004) have detected a correlation of the $2-10~keV$ XRB with $WMAP$ satellite map of the cosmic microwave background (Bennett et al. 2003), and there have been correlations observed with a number of other galaxy surveys (Nolta et al. 2003; Scranton et al. 2003; Fosalba, Gaztanaga,\\& Castander 2003; Afshordi, Loh, \\& Strauss 2003). These correlations have been interpreted as the detection of the integrated Sachs-Wolfe ($ISW$) effect (Sachs \\& Wolfe 1967).  If confirmed, they would constitute an important test of the $\\Lambda CDM$ cosmological model and provide further evidence of the existence of a substantial amount of ``dark energy'' in the universe (Crittenden \\& Turok 1996)."
    },
    "0404/astro-ph0404381_arXiv.txt": {
        "abstract": "We investigate bulge and disk scaling relations using a volume-corrected sample of early- to intermediate-type disk galaxies in which, importantly, the biasing flux from additional nuclear components has been modeled and removed. Structural parameters are obtained from a seeing-convolved, bulge$+$disk$+$nuclear-component decomposition applied to near-infrared surface brightness profiles spanning $\\sim$10 pc to the outer disk. Bulge and disk parameters, and bulge-to-disk ratios, are analyzed as a function of bulge luminosity, disk luminosity, galaxy central velocity dispersion, and galaxy Hubble type. Mathematical expressions are given for the stronger relations, which can be used to test and constrain galaxy formation models. Photometric parameters of both bulges and disks are observed to correlate with bulge luminosity and with central velocity dispersion. In contrast, for the unbarred, early to intermediate types covered by the sample, Hubble type does not correlate with bulge and disk components, nor their various ratios. In this sense, the early-to-intermediate spiral Hubble sequence is scale-free. However, galaxies themselves are not scale-free, the critical scale being the luminosity of the bulge. Bulge luminosity is shown to affect the disk parameters, such that central surface brightness becomes fainter, and scale-length bigger, with bulge luminosity. The lack of significant correlations between bulge pararmeters (size, luminosity or density)  on disk luminosity, remains a challenge for secular evolution models of bulge growth. ",
        "introduction": "\\label{Sec:Introduction}  In a companion paper (Balcells, Graham, \\& Peletier 2007, hereafter Paper~III), we present high-resolution near-infrared (NIR) surface brightness profiles of early- to intermediate-type disk galaxies.  These are derived from {\\sl Hubble Space Telescope} (\\HST) NICMOS images and extended with NIR ground-based profiles from UKIRT images.  That paper provides profile decompositions into bulge, disk, and nuclear components, and analyzes the nuclear properties of the bulges.  In the present paper, we use the profile decompositions of Paper~III to analyze the global scaling relations for bulges and disks, free from the biasing influence of additional nuclear components.  Global scaling relations provide useful diagnostics on the structure of disk galaxies.  For bulge components, comparison with the scaling relations of elliptical galaxies shows to what degree bulges are either similar or different to ellipticals.  Whether bulges resemble spheroids or disks is still an open question (Wyse, Gilmore, \\& Franx 1997; Kormendy \\& Kennicutt 2004).  To give one example of the complexity of this issue, some apparent bulges in face-on, barred S0 galaxies are in fact inner disks (Kormendy 1993; Erwin et al.\\ 2003), which of course has implications for galaxy formation models.  For both bulges and the outer, large-scale disks, the scaling relations may provide clues on the long-standing question of the origin of the Hubble sequence.  The sample of de Jong \\& van der Kruit (1994) has provided a useful reference for intermediate- to late-type field disk galaxies, and its scaling relations have been analyzed by de Jong (1996) and Courteau, de Jong, \\& Broeils (1996).  Graham (2001, hereafter G01) and  MacArthur, Courteau, \\& Holtzman (2003) provide a more recent study of intermediate- to late-type spirals, using S\\'ersic fits for the bulge.  Hunt, Pierini, \\& Giovanardi (2003, hereafter HPG03) have also analyzed bulge and disk scaling relations from a sample of disk galaxies in the Perseus-Pisces supercluster;  their sample comprises mostly Sb and later-type spirals.  At the other end of the Hubble sequence, S0 galaxies are often discussed together with  ellipticals, but are rarely compared to spiral galaxies.   The study presented in this paper is complementary to the ones above as it samples early- to intermediate-type, S0 through Sbc, disk galaxies.  Moreover, the analysis presented here uses \\HST\\ images to subtract the flux of additional nuclear components such as star clusters and nuclear disks, which have biased previous studies of bulge, and hence disk, parameters.  This is perhaps best evidenced through the reporting of {\\r14}-like bulges with S\\'ersic indices $n\\sim 4-6$ when using low-resolution ground-based data which smear out the flux from the unresolved nuclear components (e.g., Andredakis, Peletier, \\& Balcells 1995, hereafter APB95).  Higher-resolution studies with \\HST\\ have since revealed that the majority of such galaxies have noticeably less concentrated bulges, with $n \\lesssim 3$, but clear additional components (Balcells et al.~2003, hereafter Paper~II; Paper~III).  Bulge and disk scaling relations have traditionally been studied as a function of Hubble type, commonly parametrized with the revised type index T (de Vaucouleurs et al.\\ 1991, hereafter RC3).  The lack of a correlation between the ratio of the bulge effective radius $\\reff$  and the disk scale length $h$ with Hubble type, led to the statement that the spiral Hubble sequence is scale-free (de Jong 1996; Courteau et al.\\ 1996).  In a scale-free situation, an 'ice-berg' model, in which bulge surface brightness rather than effective radius determines how much the bulge protrudes above the disk, might explain the higher prominence of earlier-type bulges (G01). That the Hubble sequence is scale-free has implications for bulge formation models, as secular evolution models (Pfenniger \\& Norman 1990; Norman, Sellwood, \\& Hasan 1996) predict that bulge and disk scale lengths are correlated (Combes et al.\\ 1990).  The parameters used for the scaling relations, derived from NIR images, are less affected by dust extinction than those derived from optical images.  Furthermore, $M/L$ variations with population age and metallicity are small in the NIR, yielding small differences between the photometric length-scales and the stellar-mass scales of the galaxies. We will show that bulge luminosity is indeed a key yardstick that traces the values of bulge, disk, and global galaxy parameters.  Section~\\ref{Sec:Data} describes the sample and the profile decomposition.  Section~\\ref{Sec:Completeness} discusses selection biases and provides a volume correction.  The main results of the paper are contained in Figures \\ref{Fig:BDparMK}, \\ref{Fig:BDparDisMK}, \\ref{Fig:BDparSIG}, and \\ref{Fig:BDparT}, and analyzed in Section~\\ref{Sec:GlobalParameters}, for bulges (\\S\\ref{Sec:BulgeParameters}), disks (\\S\\ref{Sec:DiskParameters}), and bulge-to-disk scaling relations (\\S\\ref{Sec:BDscaling}).  We discuss two specific issues in Section~\\ref{Sec:Discussion}: the trends of the galaxy central surface brightness with spheroid luminosity (\\S\\,\\ref{Sec:DiscussionMu0}), and the scale-free nature of the Hubble sequence (\\S\\,\\ref{Sec:DiscussionHubbleSequence}). A Hubble constant of $H_{0} = 75$ km\\,s$^{-1}$\\,Mpc$^{-1}$ is used throughout. ",
        "conclusions": "\\label{Sec:Discussion}  \\subsection{The trend of central surface brightness with bulge luminosity} \\label{Sec:DiscussionMu0}  Whether the central surface brightness of spheroidal components becomes brighter or fainter with increasing spheroid luminosity provides useful diagnostics on galaxy formation.  The sign of this correlation results from the combined effects of dissipation (e.g., Kormendy 1989), feedback from star formation (e.g., Silk 2005), and dynamical heating by merging black holes (e.g., F97). In mergers between unequal galaxies, it determines which component ends up populating the center of the merger remnant (Zurek et al.\\ 1988; Balcells \\& Quinn 1990). Our trends of $\\mu_0$ with luminosity (Fig.\\,\\ref{Fig:mu0MB}) show that both the inwardly-extrapolated bulge $\\mu_0$, and $\\mu$ measured at a fixed radius, become brighter with luminosity (\\S\\ref{Sec:BulgeMu0}, Fig.\\,\\ref{Fig:BDparMK}c, Fig.\\,\\ref{Fig:mu0MB}).  We have verified that $\\mu$ measured at $r=20$ pc from the center, and at a fixed fraction of $\\reff$, specifically 0.2\\,$\\reff$, follow the same behavior.  Furthermore, the slope defined in Figure\\,\\ref{Fig:mu0MB}b by our bulges (brighter $\\mu_0$ for more luminous objects) becomes steeper the closer to the center we measure $\\mu$.  Jergen et al.\\ (2000) have shown that, at the faint end of spheroid luminosity, the highly resolved, and also common, dwarf galaxies around the Milky Way and M31 have central profiles that do not show central excesses over their outer S\\'ersic profile fits and have very low central surface brightness. The trend of brighter central surface brightness for higher luminosities is therefore well established for spheroids of vastly different luminosities, from giant ellitpicals to dEs (Fig.\\,\\ref{Fig:mu0MB}a).  Our findings contrast with the arguments of F97 who favor of a trend of {\\sl fainter} central surface brightness with increasing luminosity for ellipticals and bulges.  F97 reasoned that, if M32 resided in the Virgo Cluster, its exceedingly high central surface brightness would be smoothed by the PSF to values close to those of other low-luminosity Virgo E's.   By the same argument, bulges and low-luminosity ellipticals at distances comparable to Virgo might have M32-like central profiles which, with sufficient spatial resolution, would yield $\\mu_0$ values following  a faint-ward extrapolation of the trend defined by giant ellipticals, with $M_B + 5\\,\\log h_{50} \\la -20.5$ mag, in Figure~\\ref{Fig:mu0MB}b.  Although Jerjen et al.\\ (2000) have shown that the picture described by F97 breaks down at the low-luminosity end, one could argue that the central {\\sl galaxy} surface density, rather than {\\sl bulge} surface density, of some nucleated, intermediate-luminosity galaxies and bulges may follow the trend argued for by F97 (unresolved components [PS] reside in the centers of 58\\% of our bulges, see Paper~III). Specifically, if the central surface brightness of the PSs, $I_{0,\\mathrm{PS}}$, became sufficiently {\\sl brighter} for fainter PS (and hence bulge) luminosities, the trend between bulge luminosity and {\\sl galaxy} central surface brightness could in principle have an opposite slope to that shown in Figure\\,\\ref{Fig:mu0MB} for the trends between bulge luminosity and {\\sl bulge} central surface brightness.  We find this unlikely. Both the central surface brightness of the underlying S\\'ersic component ($\\IZeroSer \\sim (\\LKBul)^{0.81}$, eqn.\\,\\ref{Eqn:mu0serMK}), and the PS luminosities ($\\LumPt \\sim (\\LKBul)^{0.91}$, Paper~III) rapidly become fainter toward lower luminosities, requiring exceedingly high densities for the unresolved components (over 5 magnitudes in surface brightness) in order to reverse the trend of Figure\\,\\ref{Fig:mu0MB}.  B\\\"oker et al.\\ (2004), using \\HST/PC imaging,  find that the sizes of nuclei do not correlate with nuclei luminosities, hence the central surface brightness \\textsl{of the nuclear clusters} $I_{0,\\mathrm{PS}}$ does not correlate with \\LumPt\\  (for a sample of late-type spiral nuclei.)  The compact elliptical galaxy M32 does provide a counter-example to the previous reasoning.  The ground-based $R$-band surface brightness profile of Kent (1987; see Graham 2002), and the color profiles from Peletier (1993), yield $\\mu_H$(50\\,pc)=14.64 mag.  Given the absolute magnitude of $M_B=-15.74$ for the M32 bulge\\footnote{We assume a bulge-to-disk ratio of $\\log(B/D)=0.22$ from Graham (2002); we use the apparent blue magnitude from LEDA, and a distance modulus of 24.43 (Jacoby et al.\\ 1992).  Ignoring the bulge-disk decomposition in determining the absolute magnitude would not affect our argument.}, M32 lies within the sequence of bulges and low-luminosity ellipticals in Figure\\,\\ref{Fig:mu0MB}.  Yet, the \\HST\\ profiles in F97 yield $\\mu_H$(0.1\\arcsec)=8.52 mag, i.e., five magnitudes above bulges and ellipticals of the same luminosity.  Note that the M32 S\\'ersic profile shape index derived by Graham (2002) is a modest $n=1.5$, well within the range for bulges of similar luminosities; the high central brightness that makes M32 a prototype, compact elliptical galaxy really refers to a region of radius smaller than 37 parsec only (Graham 2002 excluded radii smaller than 10\\arcsec\\ from his fit). This inner region is M32's  'unresolved source', using the term applied to our bulges.  If the unresolved sources in bulges had similar structure to the M32 nucleus, they could reach very high central surface brightnesses despite their low luminosities. However, it is unclear whether M32 is representative of spheroids such as those analyzed in this paper; the bulge of M31 might be a more adequate model; it does not show such high central surface brightness.  \\subsection{Bulge luminosity as a measure of the structure of disk galaxies} \\label{Sec:DiscussionHubbleSequence}  Of the four sets of correlations involving bulge luminosity, disk luminosity, bulge central velocity dispersion, and galaxy Hubble type T, relations involving bulge luminosity show the strongest level of significance.  This was perhaps expected for bulge intrinsic parameters, but it is significant that disk parameters and bulge-to-disk ratios correlate with bulge luminosity as well.  They also do so with the central velocity dispersion $\\sigma_{0}$, as a result of the correlation between $\\sigma_{0}$ and luminosity for bulges.   In contrast, the dependencies with T are weak or non-existent.  Such lack of scaling with T imply a stronger form of the statement from de Jong (1996) and Courteau et al.\\ (1996) that the Hubble sequence is scale-free.  These authors noted that bulge and disk scale lengths are correlated, and that their ratio $\\reff/h$ is independent of Hubble type.  We reproduce those results in our sample, as do HPG03 in theirs.  Here, however, the lack of dependence on Hubble type is not restricted to length scales: the distributions of {\\sl most} bulge, disk, and bulge-to-disk parameters versus T shows an absence of trends (Fig.\\,\\ref{Fig:BDparT}) for our early-type disk galaxy sample.  HPG03 find a similar absence of correlations for a later-type sample, selected with different criteria (galaxies belonging to the Perseus-Pisces supercluster), and analyzed with a different (2D) code, suggesting that this behavior may be common for bulge-disk galaxies.  Although the Hubble sequence is scale-free, galaxies themselves are not. The bulge luminosity (and velocity dispersion) correlate well with the other properties of the bulge, the disk, and the scaling of these two components.  What structural parameters define the Hubble sequence?  Focussing on the early-to-intermediate disk galaxies covered by our sample (S0--Sbc), structural changes with T are only found in the latest, Sbc galaxies (lower B/D,  lower $I_{0,Bul}/I_{0,Disk}$ and marginally lower central velocity dispersions).  For types S0 to Sb, Hubble type must be  dominantly given by the spiral pattern (Block \\& Puerari 1999; Seigar, Chomey, \\& James 2003), dust content, and star formation activity.  Furthermore, the latter must not be affected by the spheroid luminosity or mass, nor by the relative sizes and brightness of disk and bulge. It appears that \\textit{the luminosity and size trends commonly associated with the definition of the Hubble sequence appear when extreme late-type galaxies are compared to early-types; but those trends are not gradual, and are absent among early-to-intermediate disk galaxies}.  That bulge luminosity strongly correlates with the properties of both the disk and the bulge-to-disk ratios may be expected from galaxy formation scenarios in which the bulge precedes the disk.  This includes monolithic collapse models akin to the Eggen, Lynden-Bell, \\& Sandage (1962) model for the formation of the Milky Way, or CDM-based hierarchical galaxy formation scenarios in which disks grow around already formed bulges of merger origin (e.g.\\ Baugh et al.\\ 1996; Abadi et al.\\ 2003; Sommer-Larsen,  G{\\\" o}tz, \\& Portinari 2003).  Our results are probably also compatible with the inside-out galaxy formation model of van den Bosch (1998).  The early formation of bulges does not necessarily imply a classical, slowly-rotating spheroid structure for bulges; if star formation timescales exceed the dynamical time, such bulges might show rapid rotation and some of the disky properties highlighted by Kormendy (1993).  Are our results compatible with secular evolution models in which the bulge grows from instabilities in an already formed disk?  The scaling of bulge and disk parameters is generally taken to support secular evolution models (Courteau et al.\\ 1996; Zhang 2004). Here, we find that most bulge parameters show non-significant correlations with the disk $K$-band luminosity, suggesting that bulge properties (\\reff, mass, central density) do not know how massive is the host disk.  Perhaps more puzzling, the disk parameters show more well-defined trends with bulge luminosity than with disk luminosity.  Indeed, only one galaxy parameter, namely the disk scale length, correlates with the disk luminosity ($\\LumKDis \\sim h^{\\sim 2}$ (eqn.~\\ref{Eqn:hMKDis}).  Explaining these relations, or lack of thereof, between galaxy structural parameters, will remain key challenges for secular evolution models.  This difficulty is compounded by the evidence of stellar population ages.  We show in Paper~I that the bulge populations in the present sample are as old as cluster ellipticals.  Our measurements were performed away from the disk plane, hence ongoing or episodic star formation of secular origin could take place in inner regions of the disk. But  old population ages are inferred for the nuclei of later-type spirals from NIR spectroscopy (Bendo \\& Joseph 2004).  To be compatible with this evidence, secular evolution must have been of little relevance for the growth of the bulge, or else confined to look-back times of order 10 Gyr for our galaxies: the situation essentially reduces to the van den Bosch model; and, timescales need to be known to verify that the term 'secular' applies.  The situation may be different for bulges of later types, and of barred galaxies, which are not addressed in this paper (see, e.g., Fathi \\& Peletier 2003; Castro-Rodr\\'\\i guez \\& Garz{\\' o}n 2004); although, in late-type bulges, young populations are statistically associated with morphological disturbances, pointing to external, rather than internal-secular, processes (Kannappan et al.\\ 2004).  Hence, from the population information in Paper~I and the structural information in this paper, it appears that  bulges of unbarred, early- to intermediate-type disk galaxies cannot have significantly grown from disk instabilities in the last $\\sim$10 Gyr. A small amount of star formation is probably occurring in almost every galaxy nucleus, as evidenced by dust in the very centers of bulges (e.g.\\ Paper~I), and by the presence of young stars (Freeman \\& Bland-Hawthorn 2002).  Bar dynamics could be feeding such activity, but, for early-type disk galaxies, star formation level is too low to affect the general bulge population.  Our conclusions on bulge formation/growth are consistent with the current understanding that many 'mature' galaxies were largely in place by $z=1$ (e.g.\\ Brinchmann \\& Ellis 2000; Simard et al.\\ 2002), and that early-type galaxies in the field became red between redshifts $z=1$ and 2 (Eliche-Moral et al.~2006)."
    },
    "0404/hep-ph0404033_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{I}   The standard model (SM) of electroweak interactions meets all the requirements --commonly referred to as Sakharov conditions~\\cite{Sakharov}-- to generate a baryon asymmetry during the electroweak phase transition (EWPT), provided that this last be of first order. However, it is also well known that neither the amount of CP violation within the minimal SM nor the strength of the EWPT are enough to generate a sizable baryon number~\\cite{{Gavela},{Kajantie}}. Supersymmetric extensions of the SM, with a richer particle content, contain new sources of CP violation~\\cite{CKN}. They also allow a stronger first order phase transition~\\cite{Carena}. In spite of these improvements with respect to the SM, the minimal supersymmetric model (MSSM) is severely constrained from experimental bounds on the chargino properties~\\cite{Cline} leaving only a small corner of parameter space for the MSSM as a viable candidate for baryogenesis. Further possibilities to accommodate an explanation for the generation of baryon number during the EWPT include non-minimal supersymmetric models which, nonetheless, all share the unappealing feature of containing an even larger set of parameters than the already extensive number contained in the MSSM.  Though it might appear tempting to abandon the idea of electroweak baryogenesis (EWB) given the above difficulties, in recent years this possibility has been revisited due to the observation that magnetic fields are able to generate a stronger first order EWPT~\\cite{{Giovannini},{Elmfors},{Giovannini2}}. The situation is similar to what happens to a type I superconductor where an external magnetic field modifies the nature of the superconducting phase transition, changing it from second to first order due to the Meissner effect.  In spite of this observation, it has also been realized that the sphaleron bound becomes more restrictive due to the interaction between the sphaleron's magnetic dipole moment and the external field~\\cite{Comelli}. Nevertheless, these arguments are either classical or resort to perturbation theory to lowest order. However, in the absence of magnetic fields, it is well known that the phase transition picture is influenced by non-perturbative effects cast in terms of the resummation of certain classes of diagrams~\\cite{Carrington} and it might also be expected that the same is true in the presence of magnetic fields. The situation with regards to the strength of the phase transition in the presence of magnetic fields is thus far from being settled and requires further research.  However, the influence on the magnetic fields on the enhancement of CP violation has received much less attention. In a series of recent papers~\\cite{Ayala1, Ayala2, Ayala3}, it has been shown that the external field is able to produce an axially asymmetric scattering of fermions off first order phase transition bubbles during the EWPT. This CP violating reflection is due to the chiral nature of the couplings of right- and left-handed modes with the external field in the symmetric phase . This mechanism produces an axial charge segregation which can then be transported in the symmetric phase where sphaleron induced transitions can convert it into baryon number~\\cite{Nelson}. The main purpose of our work is the description of the mechanism for the generation of this axial charge segregation.  As was briefly presented in this introduction, the process of baryogenesis involves several physical ingredients which all deserve to be addressed in order to  cover the entire topic; here we will concentrate on those aspects related to hypermagnetic fields, which are also one of the most recent and less explored parts of this field. For other aspects on the subject, we refer the reader to recent reviews~\\cite{Trodden, Riotto, Rubakov}.  The work is organized as follows: in Sec.~\\ref{II} we present the basic framework of electroweak baryogenesis. Sec.~\\ref{III} is devoted to generalities of hypermagnetic fields and phase transitions. In Sec.~\\ref{IV} we summarize the current ideas for the origin of large scale magnetic fields as well as the experimental bounds on their strength set by different observations. In Sec.~\\ref{V} we describe the mechanism whereby the asymmetric reflection of fermions off first order EWPT bubbles in the presence of external magnetic fields leads to an axial charge segregation that can then be converted into baryon number in the symmetric phase. Finally, in Sec.~\\ref{VI} we summarize and give a brief account of some possible research venues to further explore the influence of primordial magnetic fields in the generation of the baryon asymmetry of the universe. ",
        "conclusions": "\\label{VI}  In this work we have given a quantitative outline of the CP violating scattering of fermions off (a simplified picture of) EWPT bubbles in the presence of hypermagnetic fields. This scattering produces an axial asymmetry built on either side of the bubble walls. The origin of this asymmetry is the chiral nature of the fermion coupling to the hypermagnetic field in the symmetric phase. We have shown how to compute reflection and transmission coefficients and also that these differ for left and right-handed incident particles.  Primordial hypermagnetic fields thus provide with a much needed ingredient, namely, additional CP violation, for the possible generation of baryon number during the EWPT. A second ingredient, the strengthening of the order of the phase transition and thus the avoidance of the sphaleron bound seems at the moment a difficult problem to surmount. Nonetheless, it is important to bear in mind that so far, the calculations that provide insight into the effect of the hypermagnetic fields on the order of the EWPT do not account for the non-perturbative effects, cast in the language of resummation, which are otherwise well known to play a very important role for the dynamics of the phase transition in the absence of magnetic fields. Much work is needed in this direction. This is for the future."
    },
    "0404/hep-ph0404175_arXiv.txt": {
        "abstract": "In this review article, we discuss the current status of particle dark matter, including experimental evidence and theoretical motivations. We discuss a wide array of candidates for particle dark matter, but focus on neutralinos in models of supersymmetry and Kaluza-Klein dark matter in models of universal extra dimensions. We devote much of our attention to direct and indirect detection techniques, the constraints placed by these experiments and the reach of future experimental efforts. ",
        "introduction": "\\subsection{Overview} A great deal of effort has been made since 1687, the year of publication of Issac Newton's classic work {\\it``Philosophiae Naturalis Principia Mathematica''}, towards explaining the motion of astrophysical objects in terms of the laws of gravitation. Since then, the deviations of observed motions from expected trajectories have proved very effective in deepening our understanding of the Universe. Whenever anomalies were observed in the motion of planets in the Solar system, the question arose: should such anomalies be regarded as a refutation of the laws of gravitation or as an indication of the existence of unseen (today we would say ``dark'') objects?  The second approach proved to be correct in the case of the anomalous motion of Uranus, which led the French astronomer U.~Le Verrier and the English astronomer John Couch Adams to conjecture the existence of Neptune, eventually discovered in 1846 by J.G.~Galle. Conversely, the attempt to explain the anomalies in the motion of Mercury as due to the  existence of a new planet, called Vulcan, failed, and the final solution had to wait for the advent of Einstein's theory of general relativity, \\ie the introduction of a more refined description of the laws of gravitation.  The modern problem of dark matter is conceptually very similar to the old problem of unseen planets. We observe in large astrophysical systems, with sizes ranging from galactic to cosmological scales, some ``anomalies'' that can only be explained either by assuming the existence of a large amount of unseen, {\\it dark}, matter, or by assuming a deviation from the known laws of gravitation and the theory of general relativity.  About ten years ago, Jungman, Kamionkowski and Griest wrote a review of supersymmetric dark matter for Physics Reports~\\cite{Jungman:1995df}. This article, although incredibly useful, complete and popular, has gradually become outdated over the last decade. With this in mind, we have endeavored to write a new review of particle dark matter. As with the Jungman {\\it et al.} article, our review is intended to be suitable for a wide range of readers. It could be used as an introduction for graduate students interested in this subject or for more experienced scientists whose research focuses in other areas. It is also intended to be a useful reference in day-to-day research for particle physicists and astrophysicists actively working on the problem of dark matter. Unlike the review by Jungman {\\it et al.}, we do not limit our discussion to supersymmetric dark matter.  The article is organized as follows: we first present, in this chapter, a brief review of the Standard Model of particle physics and cosmology, and review our present understanding of the history of the Universe. We focus in particular on the freeze-out of dark matter particles and on the calculation of their relic abundance, and discuss the possible relationship between dark matter and physics beyond the Standard Model of particle physics.  Chapter 2 is devoted to the compelling evidence for dark matter at all astrophysical length scales. We review the key observations and discuss the theoretical predictions (from N-body simulations) for the distribution of dark matter, focusing in particular on the innermost regions of galaxies, and discuss how they compare with observations. Particular attention is devoted to the galactic center, where the presence of a supermassive black hole could significantly modify the dark matter distribution.  Dark matter candidates are presented in chapter 3. We start with an introduction to the ``dark matter zoo'', \\ie a description of the many candidates that have been proposed in the literature. We then focus on two particularly interesting dark matter candidates: the supersymmetric neutralino and Kaluza-Klein dark matter. For each of these candidates, we give a brief introduction to the physical motivations and underlying theories. We conclude chapter 3 with a review of the constraints put on dark matter from collider experiments, and discuss the prospects for future experiments.  The second part of this review is devoted to the astrophysical constraints on particle dark matter. We begin in chapter 4 with a review of  existing and next-generation experiments that will probe the nature of dark matter. This chapter is propedeutical to chapter 5, which discusses the many possible direct and indirect searches of dark matter and which constitutes the heart of this review. We give our conclusions in chapter 6. Some useful particle physics details are given in the appendices.      \\subsection{Standard Cosmology}  Although the exact definition of the {\\it Standard} cosmological model evolves with time, following the progress of experiments in measuring the cosmological parameters, most cosmologists agree on a fundamental picture, the so-called {\\it Big Bang} scenario, which describes the Universe as a system evolving from a highly compressed state existing around $10^{10}$~years ago.  This picture has its roots in the discovery of Hubble's law early in  the past century, and has survived all sorts of cosmological observations, unlike alternative theories such as the ``steady state cosmology'', with continuous creation of baryons, which, among other problems, failed to explain the existence and features of the cosmic microwave background.  We now have at our disposal an extremely sophisticated model, allowing us to explain in a satisfactory way the thermal history, relic background radiation, abundance of elements, large scale structure and many other properties of the Universe. Nevertheless, we are aware that our understanding is still only partial. It is quite clear that new physics is necessary to investigate the first instants of our Universe's history (see section \\ref{beyond}).  To ``build'' a cosmological model, in a modern sense, three fundamental ingredients are needed: \\begin{itemize} \\item {\\it Einstein equations}, relating the geometry of the Universe with its matter and energy content \\item {\\it metrics}, describing the symmetries of the problem \\item {\\it Equation of state}, specifying the physical properties of the matter and energy content \\end{itemize}  The Einstein field equation can be derived almost from first principles, assuming that: 1) the equation is invariant under general coordinate transformations, 2) the equation tends to Newton's law in the limit of weak fields, and 3) the equation is of second differential order and linear in second derivatives~\\cite{Ohanian:uu}. The resulting equation reads\\index{Einstein equations} \\be R_{\\mu \\nu}- \\frac{1}{2}g_{\\mu \\nu}R = -\\frac{8 \\pi G_N}{c^4}T_{\\mu \\nu}+ \\Lambda g_{\\mu \\nu}, \\ee where $R_{\\mu \\nu}$ and $R$ are, respectively, the Ricci tensor and scalar (obtained by contraction of the Riemann curvature tensor). $g_{\\mu \\nu}$ is the metric tensor, $G_N$ is Newton's constant, $T_{\\mu \\nu}$ is the energy-momentum tensor, and $\\Lambda$ is the so--called cosmological constant.  Ignoring for a moment the term involving the cosmological constant, this equation is easily understood. We learn that {\\it the geometry of the Universe}, described by the terms on the left-hand-side, {\\it is determined by its energy content}, described by the energy-momentum tensor on the right-hand-side. This is the well known relationship between the matter content and geometry of the Universe, which is the key concept of general relativity.  The addition of the cosmological constant term, initially introduced by Einstein to obtain a stationary solution for the Universe and subsequently abandoned when the expansion of the Universe was discovered, represents a ``vacuum energy'' associated with space-time itself, rather than its matter content, and is a source of gravitational field even in the absence of matter. The contribution of such ``vacuum energy'' to the total energy of the Universe can be important, if one believes recent analyses of type Ia supernovae and parameter estimates from the cosmic microwave background (for further discussion see section~\\ref{cmb}).  To solve the  Einstein equations one has to specify the symmetries of the problem. Usually one assumes the properties of statistical {\\it homogeneity} and {\\it isotropy} of the Universe, which greatly simplifies the mathematical analysis. Such properties, made for mathematical convenience, are confirmed by many observations. In particular, observations of the Cosmic Microwave Background (CMB) have shown remarkable isotropy (once the dipole component, interpreted as due to the Earth motion with respect to the CMB frame, and the contribution from the galactic plane were  subtracted). Isotropy alone, if combined with the Copernican principle, or ``mediocrity'' principle, would imply homogeneity. Nevertheless, direct evidence of homogeneity comes from galaxy surveys, suggesting a homogeneous distribution at scales in excess of  $\\sim 100$ Mpc. More specifically, spheres with diameters larger than $\\sim 100$ Mpc centered in any place of the Universe should contain, roughly, the same amount of matter.  The properties of isotropy and homogeneity imply a specific form of the metric: the line element can in fact be expressed as \\be ds^2=-c^2\\mbox{d}t^2+a(t)^2\\left({\\mbox{d}r^2\\over 1-kr^2}+r^2\\mbox{d} \\Omega^2\\right), \\label{ds2} \\ee where $a(t)$ is the so-called {\\it scale factor} and the constant $k$, describing the spatial curvature, can take the values $k= -1,\\,0,\\,+1$. For the simplest case, $k=0$, the spatial part of Eq.~\\ref{ds2} reduces to the metric of ordinary (flat) Euclidean space.  The Einstein equations can be solved with this metric, one of its components leading to the Friedmann equation \\index{Friedmann equation} \\be \\left({\\dot a\\over a}\\right)^2+{k\\over a^2}={8\\pi G_N\\over 3}\\rho_{tot}, \\label{eq:friedmann} \\ee where $\\rho_{tot}$ is the total average energy density of the universe. It is common to introduce the Hubble parameter \\be H(t) ={\\dot a(t)\\over a(t)} \\;\\;. \\ee  A recent estimate~\\cite{Contaldi:2003hi} of the present value of the Hubble parameter, $H_0$, (also referred to as the Hubble {\\it constant}) is $H_0=73 \\pm 3$ km s$^{-1}$ Mpc$^{-1}$. We see from Eq.~\\ref{eq:friedmann} that the universe is flat ($k=0$) when the energy density equals the {\\em critical density}, $\\rho_{c}$: \\be \\rho_{c}\\equiv {3H^2\\over 8\\pi G_N} \\;\\;. \\label{rhocrit} \\ee In what follows we will frequently express the abundance of a substance in the Universe (matter, radiation or vacuum energy), in units of $\\rho_{c}$. We thus define the quantity $\\Omega_i$ of a substance of species $i$ and density $\\rho_i$ as \\be \\Omega_i \\equiv {\\rho_i \\over \\rho_{c}} \\;\\;. \\label{omega} \\ee It is also customary to define \\be \\Omega =\\sum_i \\Omega_i \\equiv \\sum_i {\\rho_i \\over \\rho_{c}}, \\ee in terms of which the Friedmann equation (Eq.~\\ref{eq:friedmann}) can be written \\be \\Omega-1={{k}\\over {H^2 a^2}} \\ . \\ee The sign of $k$ is therefore determined by whether $\\Omega$ is greater than, equal to, or less than one (see table 1).  Following Ref.~\\cite{Bergstrom:2000pn}, we note that the various  $\\Omega_i$ evolve with time differently, depending on the equation of state of the component. A general expression for the expansion rate is \\ba {H^2(z)\\over H_0^2}=\\left[ \\Omega_X\\left(1+z\\right)^{3(1+\\alpha_X)} +\\Omega_K\\left(1+z\\right)^2\\right. \\nonumber \\\\ \\left. +\\Omega_M\\left(1+z\\right)^3 +\\Omega_R\\left(1+z\\right)^4 \\right] \\label{eq:general} \\ea where $M$ and $R$ are labels for matter and radiation, $\\Omega_K={-k\\over a_0^2H_0^2}$ and X refers to a generic substance with equation of state $p_X=\\alpha_X\\rho_X$ (in particular, for the cosmological constant, $\\alpha_\\Lambda=-1$). $z$ is the redshift.  We discuss in Sec.~\\ref{cmb} recent estimates of cosmological parameters using CMB measurements, combined with various astrophysical observations.  \\begin{center} \\begin{table} \\centering \\begin{tabular}{|c|ccc|} \\hline $\\rho<\\rho_{c}$ & $\\Omega < 1$ &$k=-1$ & ${\\rm open}$\\\\ $\\rho=\\rho_{c}$ & $\\Omega = 1$ &$k=0$ & ${\\rm flat}$\\\\ $\\rho>\\rho_{c}$ & $\\Omega > 1$ &$k=1$ & ${\\rm closed}$\\\\ \\hline \\end{tabular} \\label{tabUniv} \\caption{Classification of cosmological models based on the value of the average density, $\\rho$, in terms of the critical density, $\\rho_{c}$.} \\vskip 0.3cm \\end{table} \\end{center}  \\subsection{The Standard Model of Particle Physics} \\label{sec:SM} \\index{Standard Model of particle physics} The Standard Model (SM) of particle physics has, for many years, accounted for all observed particles and interactions \\footnote{It is a matter of definition whether one considers neutrino masses as part of the SM or as physics beyond the SM.}. Despite this success, it is by now clear that a more fundamental theory must exist, whose low-energy realization should coincide with the SM.  In the SM, the fundamental constituents of matter are fermions: {\\it quarks} and {\\it leptons}. Their interactions are mediated by integer spin particles called {\\it gauge bosons}. Strong interactions are mediated by gluons $G_a$, electroweak interaction by  $W^{\\pm}$, $Z_0$, $\\gamma$ and the Higgs boson $H^0$. The left-handed leptons and quarks are arranged into three generations of $SU(2)_L$ doublets \\begin{equation} \\label{su2doublets} \\left(\\begin{array}{c} \\nu_e \\\\ e^- \\end{array}\\right)_L\\qquad \\left(\\begin{array}{c} \\nu_\\mu \\\\ \\mu^- \\end{array}\\right)_L\\qquad \\left(\\begin{array}{c} \\nu_\\tau \\\\ \\tau^- \\end{array}\\right)_L \\end{equation} \\begin{equation}\\label{2.66} \\left(\\begin{array}{c} u \\\\ d^\\prime \\end{array}\\right)_L\\qquad \\left(\\begin{array}{c} c \\\\ s^\\prime \\end{array}\\right)_L\\qquad \\left(\\begin{array}{c} t \\\\ b^\\prime \\end{array}\\right)_L \\end{equation} with the corresponding right-handed fields transforming as singlets under $ SU(2)_L $. Each generation contains two flavors of quarks with {\\it baryon number} $B=1/3$ and {\\it lepton number} $L=0$ and two leptons with $B=0$ and $L=1$. Each particle also has a corresponding antiparticle with the same mass and opposite quantum numbers.  The quarks which are primed are {\\it weak eigenstates} related to {\\it mass  eigenstates} by the Cabibbo-Kobayashi-Maskawa (CKM) matrix \\be \\label{ckm} \\left(\\begin{array}{c} d^\\prime \\\\ s^\\prime \\\\ b^\\prime \\end{array}\\right)= \\left(\\begin{array}{ccc} V_{ud}&V_{us}&V_{ub}\\\\ V_{cd}&V_{cs}&V_{cb}\\\\ V_{td}&V_{ts}&V_{tb} \\end{array}\\right) \\left(\\begin{array}{c} d \\\\ s \\\\ b \\end{array}\\right)=\\hat V_{\\rm CKM}\\left(\\begin{array}{c} d \\\\ s \\\\ b \\end{array}\\right). \\end{equation}   Gauge symmetries play a fundamental role in particle physics. It is in fact in terms of symmetries and using the formalism of gauge theories that we describe electroweak and strong interactions. The SM is based on the  $ SU(3)_C\\otimes SU(2)_L\\otimes U(1)_Y$ gauge theory, which undergoes the spontaneous breakdown: \\be SU(3)_C\\otimes SU(2)_L\\otimes U(1)_Y \\to SU(3)_C\\otimes U(1)_Q \\label{brkdown} \\ee where $Y$ and $Q$ denote the weak hypercharge and the electric charge generators, respectively, and $SU(3)_C$ describes the strong (color) interaction, known as Quantum ChromoDynamics (QCD). This spontaneous symmetry breaking results in the generation of the massive $W^{\\pm}$ and $Z$ gauge bosons as well as a massive scalar Higgs field.   \\subsection{A very brief history of the Universe}  Our description of the early Universe is based on an extrapolation of known physics back to the Planck epoch, when the Universe was only $t=10^{-43}$seconds old, or equivalently up to energies  at which the gravitational interaction becomes strong (of the order of the Planck mass, $M_{Pl}=10^{19}$~GeV). Starting at this epoch we take now a brief tour through the evolution of the Universe:  \\begin{itemize} \\item $T \\sim 10^{16}$~GeV. It is thought that at this scale, some (unknown) grand unified group, $G$, breaks down into the Standard Model gauge group, $SU(3)_C\\otimes SU(2)_L\\otimes U(1)_Y$. Little is known about this transition, however.  \\item $T \\sim 10^{2}$~GeV. The Standard Model gauge symmetry breaks into $SU(3)_C \\otimes U(1)_Q$ (see Eq.~\\ref{brkdown}). This transition, called electroweak symmetry breaking, could be the origin of baryogenesis (see \\eg Ref.~\\cite{Affleck:1984fy}) and possibly of primordial magnetic fields (\\eg Ref.~\\cite{Joyce:1997uy}).  \\item $T \\sim 10^1 - 10^3$~GeV. Weakly interacting dark matter candidates with GeV-TeV scale masses {\\it freeze-out}, as discussed in next section. This is true in particular for the {\\it neutralino} and the $\\bone$ Kaluza-Klein excitation that we discuss in Chapter~\\ref{candidates}.  \\item $T \\sim 0.3$~GeV. The QCD phase transition occurs, which drives the confinement of quarks and gluons into hadrons.  \\item $T \\sim 1$~MeV. Neutron freeze-out occurs.  \\item $T \\sim 100$~keV. Nucleosynthesis: protons and neutrons fuse into light elements (D, $^3$He, $^4$He, Li). The standard Big Bang nucleosynthesis (BBN) provides by far the most stringent constraints to the Big Bang theory, and predictions remarkably agree with observations (see Fig.~\\ref{bbn}).  \\item $T \\sim 1$~eV. The matter density becomes equal to that of the radiation, allowing for the formation of structure to begin.  \\item $T \\sim 0.4$~eV. Photon decoupling produces the cosmic background radiation (CMB), discussed in Sec.~\\ref{cmb}.  \\item $T = 2.7 K \\sim 10^{-4}$~eV. Today. \\end{itemize}  \\begin{figure}[t!] \\centering \\includegraphics[width=0.7\\textwidth,clip=true]{figures/bbn_2003.eps} \\caption[Big Bang nucleosynthesis predictions for the abundances of \\qua , \\deu, \\tro\\ and \\sep] {Big Bang nucleosynthesis predictions for the abundances of light elements as a function of the baryon over photon ratio $\\eta$ or $\\Omega_b h^2$~\\cite{Coc:2003ce}. From Ref.~\\cite{subirbbn}.} \\label{bbn} \\end{figure}   \\subsection{Relic Density} \\label{sec:relic} \\index{relic density}  We briefly recall here the basics of the calculation of the density of a thermal relic. The discussion is based on Refs.~\\cite{Kolb:vq, Gondolo:1990dk,Servant:2002hb} and we refer to them for further comments and details.  A particle species in the early Universe has to interact sufficiently or it will fall out of local thermodynamic equilibrium. Roughly speaking, when its interaction rate drops below the expansion rate of the Universe, the equilibrium can no longer be maintained and the particle is said to be {\\it decoupled}.  \\subsubsection{The standard calculation} The evolution of the phase space distribution function, $f({\\bf p},{\\bf x})$, is governed by the Boltzmann equation \\be {\\bf L}[f] = {\\bf C}[f], \\ee where ${\\bf L}$ is the Liouville operator, and ${\\bf C}$ is the collision operator, describing the interactions of the particle species considered.  After some manipulation, the Boltzmann equation can be written as an equation for the particle number density $n$: \\be \\frac{dn}{dt}+3Hn=-\\langle\\sigma v\\rangle\\left(n^2-({n^{eq}})^2\\right) , \\label{BE} \\ee where $\\sigma v$ is the total annihilation cross section multiplied by velocity, brackets denote thermal average, $H$ is Hubble constant, and $n^{eq}$ is the number density at thermal equilibrium. For massive particles, \\ie in the non-relativistic limit, and in the Maxwell-Boltzmann approximation, one has \\be n^{eq}=g\\left(\\frac{mT}{2\\pi}\\right)^{3/2}e^{-m/T} , \\label{exponentialdecrease} \\ee where $m$ is the particle mass and $T$ is the temperature. We next introduce the variables \\be Y \\equiv \\frac{n}{s} \\;\\;\\;\\; , \\;\\;\\;\\; Y^{eq} \\equiv \\frac{n^{eq}}{s} \\ee where $s$ is the entropy density $s={2\\pi^2 g_* T^3}/{45}$ and $g_*$ counts the number of relativistic degrees of freedom. Using the conservation of entropy per co-moving volume ($sa^3=$constant), it follows that $\\dot{n}+3Hn=s\\dot{Y}$ and Eq.~\\ref{BE} reads \\be s\\dot{Y}=-\\langle\\sigma v\\rangle s^2\\left(Y^2-({Y^{eq}})^2\\right) . \\label{BE2} \\ee If we further introduce the variable $x \\equiv \\frac{m}{T}$, Eq.~\\ref{BE2} can be expressed as \\be \\label{Boltz} \\frac{dY}{dx}=-\\frac{\\langle\\sigma v\\rangle s}{H x} \\left(Y^2-({Y^{eq}})^2\\right) . \\ee For heavy states, we can approximate $\\langle\\sigma v\\rangle$ with the non-relativistic expansion in powers of $v^2$ \\be \\langle\\sigma v\\rangle= a+b \\langle v^2 \\rangle + {\\cal O}(\\langle v^4 \\rangle ) \\approx a+6 \\ b/x \\; , \\label{approsig} \\ee which leads to our final version of Eq.~\\ref{Boltz} in terms of the variable $\\Delta=Y-Y^{eq}$: \\be \\label{delta} \\Delta^{\\prime}=-{Y^{eq}}^{\\prime}-f(x)\\Delta(2Y^{eq}+\\Delta) , \\ee where prime denotes $d/dx$ and \\be f(x)=\\sqrt{\\frac{\\pi g_*}{45}}\\  m \\ M_{Pl}\\ (a+6 \\ b/x) \\  x^{-2} . \\ee  Following Ref.~\\cite{Kolb:vq} we introduce the quantity $x_F \\equiv m/T_F$, where $T_F$ is the freeze-out temperature of the relic particle, and we notice that Eq.~\\ref{delta} can be solved analytically in the two extreme regions $x \\ll x_F$ and $x \\gg x_F$ \\ba \\Delta =& -\\frac{{Y^{eq}}^{\\prime}}{2f(x)Y^{eq}} \\; \\mbox{for } x \\ll x_F \\\\ \\Delta^{\\prime} =& -f(x) \\Delta^{2} \\; \\mbox{for } x \\gg x_F. \\label{latetimes} \\ea These regions correspond to long before freeze-out and long after freeze-out, respectively. Integrating the last equation between $x_F$ and $\\infty$ and using $\\Delta_{x_F} \\gg \\Delta_{\\infty}$, we can derive the value of $\\Delta_{\\infty}$ and arrive at \\be \\label{Y_infty} Y_{\\infty}^{-1}=\\sqrt{\\frac{\\pi g_*}{45}}M_{Pl} \\ m \\ x_F^{-1}(a+3b/x_F). \\ee The present density of a generic relic, $X$, is simply given by $\\rho_X=m_X n_X=m_X s_0 Y_{\\infty}$, where $s_0=2889.2$ cm$^{-3}$ is the present entropy density (assuming three Dirac neutrino species). The relic density can finally be expressed in terms of the critical density (see Eq.~\\ref{omega}) \\be \\label{eq:relic} \\Omega_X h^2 \\approx \\frac{1.07 \\times 10^9 \\, \\rm{GeV}^{-1}} {M_{Pl}}\\frac{x_F}{\\sqrt{g_*}}\\frac{1}{(a+3b/x_F)} , \\ee where $a$ and $b$ are expressed in GeV$^{-2}$ and $g_*$ is evaluated at the freeze-out temperature. It is conventional to write the relic density in terms of the Hubble parameter, $h = H_0 / 100 \\, \\rm{km} \\, \\rm{s}^{-1}  \\, \\rm{Mpc}^{-1}$.  To estimate the relic density, one is thus left with the calculation of the annihilation cross sections (in all of the possible channels) and the extraction of the parameters $a$ and $b$, which depend on the particle mass. The freeze-out temperature $x_F$ can be estimated through the iterative solution of the equation \\be \\label{XF} x_F=\\ln \\left[c (c+2) \\sqrt{\\frac{45 }{8}} \\frac{g}{2\\pi^3} \\frac{m \\ M_{Pl} (a+6b/x_F)}{g^{1/2}_* x^{1/2}_F}\\right], \\ee where $c$ is a constant of order one determined by matching the late-time and early-time solutions.  It is sometimes useful to perform an order-of-magnitude estimate using an approximate version of Eq.~\\ref{eq:relic}~\\cite{Jungman:1995df}: \\be \\Omega_X h^2 \\approx \\frac{3 \\times 10^{-27} \\mbox{cm}^3 \\mbox{s}^{-1}} {\\langle\\sigma v\\rangle} \\;\\; . \\ee  We note that the approximation introduced in Eq.~\\ref{approsig} is not always justified (see \\eg Ref.~\\cite{Jungman:1995df}). For example, Ref.~\\cite{Salati:2002md} suggests a scenario where the presence of a scalar field in the early Universe could significantly affect the value of the relic density. Furthermore, a dramatic change in the relic density can be induced by resonance enhancements or so-called {\\it coannihilations}. We discuss the effects of coannihilations in the next section.   \\subsubsection{Including coannihilations} \\index{coannihilations} Following earlier works (see Ref.~\\cite{Binetruy:1983jf}), Griest and Seckel~\\cite{Griest:1990kh} noticed that if one or more particles have a  mass similar to the relic particle and share a quantum number with it, the standard calculation of relic density fails.  Let us consider $N$ particles $X_i$ ($i=1,\\ldots,N$) with masses $m_i$ and internal degrees of freedom (statistical weights) $g_i$.  Also assume that $m_1 \\leq m_2 \\leq \\cdots \\leq m_{N-1} \\leq m_N$, and that the lightest particle is protected against decay thanks to some symmetry ({\\it i.e.} R-parity or KK-parity, for neutralinos or Kaluza-Klein particles, respectively. See section~\\ref{candidates}). We will also denote the lightest particle by $X_1$.  In this case, Eq.~\\ref{BE} becomes \\begin{equation} \\frac{dn}{dt} = -3Hn - \\sum_{i,j=1}^N \\langle \\sigma_{ij} v_{ij} \\rangle \\left( n_{i}n_{j} - n_{i}^{\\rm{eq}}n_{j}^{\\rm{eq}} \\right), \\end{equation} where $n$ is the number density of the relic particle and $n= \\sum_{i=1}^N n_{i}$, due to the fact that the decay rate of particles, $X_i$, other than the lightest is much faster than the age of the Universe. Here, \\begin{eqnarray} \\sigma_{ij}  & = & \\sum_X \\sigma (X_i X_j \\rightarrow X_{SM}) \\end{eqnarray} is the total annihilation rate for  $X_i X_j$ annihilations into a Standard Model particle. Finally, \\begin{equation} v_{ij} = \\frac{\\sqrt{(p_{i} \\cdot p_{j})^2-m_{i}^2 m_{j}^2}}{E_{i} E_{j}} \\end{equation} is the relative particle velocity, with $p_{i}$ and $E_{i}$ being the four-momentum and energy of particle $i$.  The thermal average $\\langle\\sigma_{ij}v_{ij}\\rangle$ is defined with equilibrium distributions and is given by \\begin{equation} \\langle \\sigma_{ij}v_{ij} \\rangle = \\frac{\\int d^3{\\bf p}_{i}d^3{\\bf p}_{j} f_{i}f_{j}\\sigma_{ij}v_{ij}} {\\int d^3{\\bf p}_{i}d^3{\\bf p}_{j}f_{i}f_{j}}, \\end{equation} where $f_{i}$ are distribution functions in the Maxwell-Boltzmann approximation.  The scattering rate of supersymmetric particles off particles in the thermal background is much faster than their annihilation rate. We then obtain \\begin{equation} \\label{eq:Boltzmann2} \\frac{dn}{dt} = -3Hn - \\langle \\sigma_{\\rm{eff}} v \\rangle \\left( n^2 - n_{\\rm{eq}}^2 \\right), \\end{equation} where \\begin{equation} \\label{eq:sigmaveffdef} \\langle \\sigma_{\\rm{eff}} v \\rangle = \\sum_{ij} \\langle \\sigma_{ij}v_{ij} \\rangle \\frac{n_{i}^{\\rm{eq}}}{n^{\\rm{eq}}} \\frac{n_{j}^{\\rm{eq}}}{n^{\\rm{eq}}}. \\end{equation}  Edsjo and Gondolo \\cite{Edsjo:1997bg} reformulated the thermal average into the more convenient expression \\begin{equation} \\label{eq:sigmavefffin2} \\langle \\sigma_{\\rm{eff}}v \\rangle = \\frac{\\int_0^\\infty dp_{\\rm{eff}} p_{\\rm{eff}}^2 W_{\\rm{eff}} K_1 \\left( \\frac{\\sqrt{s}}{T} \\right) } { m_1^4 T \\left[ \\sum_i \\frac{g_i}{g_1} \\frac{m_i^2}{m_1^2} K_2 \\left(\\frac{m_i}{T}\\right) \\right]^2}, \\end{equation} where  $K_{i}$ are the modified Bessel functions of the second kind and of order $i$. The quantity $W_{\\rm{eff}} $ is defined as \\begin{equation} \\label{eq:weff} W_{\\rm{eff}} = \\sum_{ij}\\frac{p_{ij}}{p_{11}} \\frac{g_ig_j}{g_1^2} W_{ij} = \\sum_{ij} \\sqrt{\\frac{[s-(m_{i}-m_{j})^2][s-(m_{i}+m_{j})^2]} {s(s-4m_1^2)}} \\frac{g_ig_j}{g_1^2} W_{ij}, \\end{equation} where $ W_{ij} = 4 E_{i} E_{j} \\sigma_{ij} v_{ij}$ and $p_{ij}$ is the momentum of the particle $X_i$ (or $X_j$) in the center-of-mass frame of the pair $X_i X_j$, and $s  =  m_{i}^2+m_{j}^2 + 2E_{i}E_{j}-2 |{\\bf p}_{i}| |{\\bf p}_{j}| \\cos \\theta$, with the usual meaning of the symbols.  The details of coannihilations in the framework of supersymmetric models are well established (see \\eg the recent work of Edsjo \\etal~\\cite{Edsjo:2003us}), and numerical codes now exist including coannihilations with all supersymmetric particles, \\eg MicrOMEGAs \\cite{Belanger:2002sq} and the new version of DarkSusy \\cite{darksusynew,Gondolo:2002tz}, publicly released in 2004. The case of coannhilations with a light top squark, such as the one required for the realization of the electroweak baryogenesis mechanism, has been discussed in Ref.\\cite{Balazs:2004bu}.  \\subsection{Links with Physics Beyond the Standard Model} \\label{beyond}  The concepts of dark energy and dark matter do not find an explanation in the framework of the Standard Model of particle physics. Nor are they understood in any quantitative sense in terms of astrophysics. It is interesting that also in the realm of particle physics, evidence is accumulating for the existence of physics {\\it beyond} the Standard Model, based on theoretical and perhaps experimental arguments.  On the experimental side, there is strong evidence for oscillations of atmospheric neutrinos (originating from electromagnetic cascades initiated by cosmic rays in the upper atmosphere) and solar neutrinos. The oscillation mechanism can be explained under the hypothesis that neutrinos {\\it do} have mass, in  contrast to the zero mass neutrinos of the Standard Model (see Ref.~\\cite{Maltoni:2003da} for a recent review).  On the theoretical side, many issues make the Standard Model unsatisfactory, for example the {\\it hierarchy problem}, \\ie the enormous difference between the weak and Planck scales in the presence of the Higgs field (this will be discussed in some detail in Sec.~\\ref{whysusy}), or the problem of unification addressing the question of whether there exists a unified description of all known forces, possibly including gravity.  The list of problems could be much longer, and it is natural to conjecture that our Standard Model is the low-energy limit of a more fundamental theory. Two examples of popular extensions of the Standard Model include: \\begin{itemize} \\item \\underline{{\\it Supersymmetry.}} As a complete symmetry between fermions and bosons, supersymmetry's theoretical appeal is very great \\cite{wesszumino}. So great, in fact, is this appeal, that it appears to many as a necessary ingredient of future extensions of the Standard Model. Many interesting features make it attractive, including its role in understanding the fundamental distinction between bosons and fermions, and the problems of hierarchy and unification discussed above. Last, but not least, it provides an excellent dark matter candidate in terms of its lightest stable particle, the {\\it neutralino}. We will present the basics of supersymmetry and the properties of the neutralino in Sec.~\\ref{sec:susy}.  \\item \\underline{{\\it Extra dimensions.}} In the search of a fundamental theory with a unified description of all interactions, physicists developed theories with extra spatial dimensions, following an early idea of Kaluza~\\cite{Kaluza:tu}, who extended to four the number of space dimensions to include electromagnetism into a ``geometric'' theory of gravitation. In theories with {\\it unified} extra dimensions, in which all particles and fields of the Standard Model can propagate in the extra dimensions, the lightest Kaluza-Klein particle, \\ie the lightest of all the states corresponding to the first excitations of the particles of the Standard Model, is a viable dark matter candidate, as we discuss in Sec.~\\ref{sec:kk}.  \\end{itemize}  Despite the fact that neutrinos are thought to be massive, they are essentially ruled out as dark matter candidates (see Sec.~\\ref{sec:zoo}). Consequently, the Standard Model does not provide a viable dark matter candidate. This is further supported by the fact that most of the dark matter is non-baryonic (see section~\\ref{cmb}). Dark matter is therefore a motivation to search for physics beyond the Standard Model (others might  say that this is {\\it evidence} for physics beyond the Standard Model).  This is a typical example of the strong interplay between particle physics, theoretical physics, cosmology and astrophysics. From one side, theoretical particle physics stimulates the formulation of new theories predicting new particles that turn out to be excellent dark matter candidates. On the other side, cosmological and astrophysical observations constrain the properties of such particles and consequently the parameters of the new theories. ",
        "conclusions": "There is compelling evidence for the existence of dark matter. Although our understanding of its nature and distribution is still incomplete, many independent observations suggest that about 30\\% of the total energy density of the Universe is made of some sort of non-baryonic dark matter. We have reviewed such observations and discussed how they compare with theoretical predictions, and in particular with the results of N-body simulations.   The dark matter problem is not only relevant to astrophysicists but also to the particle and high-energy physics community. In fact, some of the best dark matter candidates come from possible extensions of the Standard Model of particle physics. There is certainly no shortage of particle dark matter candidates found in such models. Among those proposed in literature, we have focused on the dark matter particles found in models of supersymmetry (the lightest neutralino) and models with Universal Extra Dimensions (Kaluza-Klein dark matter). Although many simple models of supersymmetry, extra dimensions or other scenarios are widely discussed by the particle and astroparticle communities, the phenomenology of the actual physical theory could be more rich and complex. Collider experiments are probing significant regions of the parameter space of these hypothetical particles. Conversely, a positive astrophysical detection of dark matter would provide invaluable information regarding the physics ``beyond the Standard Model''.  The astroparticle community has started a vigorous and broad  program of experiments that may be able to shed new light on the physics and astrophysics of dark matter. Before discussing the results of direct and indirect searches, we have reviewed the present and future experiments on which they are based.  Among the most promising dark matter searches appears to be direct detection. The current situation is complicated by the claim of a positive detection by the DAMA experiment, which have been contradicted by several other experiments. It is unclear, but more and more improbable, whether it is possible to find a theoretical scenario that accommodates all the experimental findings. The much higher (several orders of magnitude) sensitivity of future experiments should be able to solve this controversy.  Indirect dark matter detection via annihilations in the Galactic center region is also an exciting possibility, although the prospects for the observation of gamma-rays, neutrinos and synchrotron radiation from that direction strongly depend on astrophysical parameters, such as the profile of dark matter in the innermost regions, which unfortunately are poorly known. Nevertheless, the development of next-generation gamma-ray and neutrino telescopes will allow us to test many scenarios, especially if effects such as the adiabatic accretion onto the central black hole significantly enhance the dark matter density and corresponding annihilation signal. If the Galactic center turns out to contain less dark matter, observations of dwarf galaxies, external galaxies and local dark substructure may play an important role for indirect searches.  Indirect searches for dark matter through the observation of high energy neutrinos produced in dark matter annihilations in the Sun are also promising. These rates do not depend strongly on the dark matter halo distribution and are thus fairly unambiguous probes of particle dark matter models. Measurement of the positron and anti-proton spectra, which are soon to improve dramatically, can also provide an opportunity to observe products of dark matter annihilations in the galactic halo.  Collectively, the direct, indirect and collider searches for particle dark matter have incredible prospects for discovery in the coming years. We hope that this review can be a useful tool in guiding members of the scientific community closer to the goal of dark matter identification which has eluded us for so long."
    },
    "0404/astro-ph0404395_arXiv.txt": {
        "abstract": "Young massive clusters differ markedly from old globular clusters in featuring extended, rather than tidally truncated, envelopes.  Their projected-luminosity profiles are well fit by Elson-Fall-Freeman models with core radii $0.3\\ {\\rm pc} \\la r_{\\rm c}\\la 8$~pc and {\\it power-law envelopes} of negative exponent $2.0\\la \\gamma\\la 3.8$.  These envelopes form within the first few 10$^6$\\,yr and last $\\sim$10$^8$\\,--\\,10$^{9.5}$\\,yr, depending on the environment. Many YMCs show clumpy substructure that may accelerate their initial relaxation.  The cores of Magellanic-Cloud clusters show universal expansion from $r_{\\rm c}< 1$~pc at birth to $r_{\\rm c}= 2$\\,--\\,3~pc after 10$^8$\\,yr, but then seem to evolve along two bifurcating branches in a $r_{\\rm c}$\\,--\\,log(age) diagram. The lower branch can be explained by mass-loss driven core expansion during the first 10$^9$\\,yr, followed by slow core contraction and the onset of core collapse due to evaporation.  The upper branch, which shows continued core expansion proportional to logarithmic age, remains unexplained.  There is strong evidence for rapid mass segregation in young clusters, yet little evidence for top-heavy IMFs or primordial mass segregation. Finally, YMCs show similar structure throughout the Local Group and as far away as we can resolve them ($\\la$20 Mpc). ",
        "introduction": "Luminosity profiles have long been used to study the structure and evolution of old globular clusters (e.g., King 1966;  Illingworth \\& Illingworth 1976). Their application to the study of young massive clusters is of more recent origin, having begun in earnest with the classic study of ten rich, 8\\,--\\,300~Myr old clusters in the LMC by Elson, Fall, \\& Freeman (1987, hereafter EFF87). As this study demonstrated, the luminosity profiles of YMCs differ significantly and systematically from those of old globulars.  The observed differences contain important clues concerning the formation and dynamical evolution of massive star clusters.  Among the various motivations for studying luminosity profiles (hereafter LPs) of YMCs is our desire to understand the effects of mass segregation as a function of time and to separate any possible primordial mass segregation from the longer-term segregation caused by energy equipartition.  Young massive clusters offer two main advantages: They still contain stars over most of the mass range of $\\sim$0.1\\,--\\,120 $M_{\\sun}$, and they allow us to trace evolutionary effects directly as a function of cluster age.  Luminosity profiles can also yield information on core collapse, both slow and fast (``equipartition instability,'' Spitzer 1969), and on the possible formation of central black holes in clusters. Finally, such profiles can help us evaluate the impact of mass loss on cluster formation and disruption, yielding estimates of the fraction of field stars that may have originated in clusters. ",
        "conclusions": ""
    },
    "0404/astro-ph0404440_arXiv.txt": {
        "abstract": "{We present the results of new XMM observations of FR-I and FR-II radio galaxies that show the importance of jet/environment interactions for both radio-source structure and the properties of the surrounding gas. The FR-I observations reveal that the distribution of hot gas determines the radio-lobe morphology, and provide evidence that the subsonic expansion of the lobes heats the surrounding gas. The FR-II observations show that the sources are at equipartition and in pressure balance with their surroundings. Finally, we present an analysis of a sample of elliptical-dominated groups, showing that radio-source heating is common.} ",
        "introduction": "Interactions between radio galaxies and their surrounding X-ray-emitting hot gas regulate the expansion of the radio plasma and transfer large amounts of energy to the gas environment. X-ray observations of radio-galaxy environments are therefore important both for developing our understanding of radio-galaxy properties and dynamics, and for studying the impact of radio galaxies on groups and clusters.  In the early Universe, radio galaxies and quasars are a possible source of the energy input needed to reconcile cold dark matter models of structure formation and the observed properties of clusters (e.g. Sanderson et al. 2003, Roychowdhury et al. 2004). It has been argued for some time that the energy injected by AGN into cluster gas is likely to be important in balancing radiative cooling in cluster cores (e.g. Binney \\& Tabor, 1995), and it is now increasingly accepted that the energy input of radio galaxies can reconcile X-ray observations with cooling-flow models (Churazov et al. 2002, Fabian et al. 2003). X-ray observations can provide information about the heating mechanism used to transfer energy from the radio source to the gas; for example, {\\it XMM-Newton} and {\\it Chandra} observations of the nearest radio galaxy Centaurus A (Kraft et al. 2003, and these proceedings; Worrall et al. these proceedings) have shown the importance of shock-heating by young sources. Insight into other heating mechanisms has also come from X-ray observations, such as those of ``ripples'' in the Perseus cluster (Fabian et al. 2003).  Both low-power (FR-I) and powerful (FR-II) radio galaxies are expected to heat their environments. FR-I radio galaxies are thought to expand subsonically for most of their lives, so that they should be in approximate pressure balance with their surroundings; they will do $P$d$V$ work on the gas as they expand, but we would not expect shock-heating, except in young sources. In contrast, the standard model of FR-II expansion is that the jets inflate a supersonically expanding cocoon of material that shocks and compresses the surrounding gas. In this model, we might expect to be able to observe a rim of hotter gas around the radio source. X-ray observations can test the validity of these models, as they constrain properties of the radio source, such as particle content, magnetic field strength and internal pressure, and of the environment, e.g. density and pressure.  In the following sections we describe our {\\it XMM-Newton} observations of FR-I and FR-II radio galaxies. Details of these observations and further discussion are presented in Croston et al. (2003, 2004). We also discuss a survey of {\\it ROSAT}-observed groups that provides evidence that radio-source heating is common. ",
        "conclusions": "Our X-ray observations of FR-I radio galaxies have shown that: \\begin{itemize} \\item differences in the distribution of hot gas around the source lead to the development of the varied types of radio-lobe morphologies seen in FR-Is. \\item FR-Is cannot be in the minimum energy condition, and the additional pressure required for pressure balance with the surrounding gas cannot come from an increase in the density of synchrotron-emitting electrons or from entrained gas at the temperature of the surroundings. Heated, entrained material is perhaps the most plausible candidate to provide the required pressure. \\item FR-Is can heat their environments gently via subsonic expansion. \\end{itemize}  Our X-ray observations of FR-II radio galaxies have shown that: \\begin{itemize} \\item 3C~223 and 3C~284 are in the minimum energy condition and in pressure balance with their surroundings. \\item Older FR-IIs are likely to come into pressure balance and no longer shock-heat their surroundings via a supersonically expanding cocoon. \\end{itemize}  We have shown that by studying the group atmospheres of individual sources, we can investigate the energy transfer mechanisms for different types of source and at different stages in source evolution. The presence of shock-heating in sources like Cen A (Kraft et al. 2003) and more gentle heating in 3C~66B are consistent with a change of energy transfer mechanism in FR-Is as they expand and come into pressure balance. Our FR-II observations suggest that a similar picture is also true for these objects, although the overpressured, supersonic stage probably lasts for a larger fraction of the source's life.  Finally, our study of a sample of groups has shown that radio-source heating is common, and that the effects of radio-source heating on groups are significant and long-lived."
    },
    "0404/astro-ph0404506_arXiv.txt": {
        "abstract": "We present the results of UV spectroscopy of AGNs with double-peaked Balmer emission lines. In 2/3 of the objects, the far-UV resonance lines are strong, with single-peaked profiles resembling those of Seyfert galaxies. The \\ion{Mg}{ii} line is the only UV line with a double-peaked profile. In the remaining objects, the far-UV resonance lines are relatively weak but still single-peaked. The latter group also displays prominent UV absorption lines, indicative of a low-ionization absorber. We interpret the difference in the profiles of the emission lines as resulting from two different regions: a dense, low-ionization accretion disk (the predominant source of the Balmer and \\ion{Mg}{ii} lines), and a lower density, higher-ionization wind (the predominant source of the far-UV resonance lines). These results suggest a way of connecting the double-peaked emitters with the greater AGN population: in double-peaked emitters the accretion rate onto the black hole is low, making the wind feeble and allowing the lines from the underlying disk to shine through. This scenario also implies that in the majority of AGNs, the wind is the source of the broad emission lines. ",
        "introduction": "Broad, double-peaked Balmer lines are found in the spectra of approximately 20--25\\% of radio-loud AGNs at $z<0.4$ (\\cite{eracleous94}, \\cite[2003]{eracleous03a}) and in the spectra of 4\\% of {\\it all} AGNs at $z<0.33$ from the Sloan Digital Sky Survey (\\cite{strateva03}). Half of these profiles are well described by kinematic models of axisymmetric, relativistic, Keplerian disks (\\cite{chen89a}; \\cite{chen89b}), while the other half require a modified version of this model in which the disk is not axisymmetric (\\cite{eracleous95}; \\cite{strateva03}).  The hosts of double-peaked emission lines (hereafter the ``double-peaked emitters'') share a number of distinguishing characteristics which make them stand out from the average AGN and suggest that these objects represent the extreme end of the distribution of AGN properties, namely: (a) Balmer line widths that are, on average, twice as large as those of other AGNs, (b) steep broad-line Balmer decrements, (c) unusually strong low-ionization emission lines (i.e., [\\ion{O}{i}] and [\\ion{S}{ii}]), as well as unusually large [\\ion{O}{i}]/[\\ion{O}{iii}] ratios, and (d) a great deal of starlight in the optical continuum. These characteristics support the scenario of \\cite[Chen \\& Halpern (1989)]{chen89b} in which the inner accretion in these objects is an ion torus (\\cite{rees82}; known today as a radiatively inefficient accretion flow). Some of the spectroscopic properties of double-peaked emitters are reminiscent of LINERs and the relation between the two groups is underscored by the detection of double-peaked Balmer lines in several LINERs.  \\begin{figure} \\centerline{\\psfig{file=eracleous1_f1.ps,width=4.8in}} \\caption{Optical, near-UV, and far-UV emission-line profiles of three double-peaked emitters. A circular-disk model is superposed on the Balmer and \\ion{Mg}{ii} profiles for comparison. The H$\\alpha$ profile of Arp~102B after subtraction of the model is included, to show that the residual resembles the Ly$\\alpha$ profile. Absorption lines marked with a ``G'' come from the gas in the Milky way, while all others arise in gas associated with the AGN.} \\end{figure} ",
        "conclusions": ""
    },
    "0404/astro-ph0404260_arXiv.txt": {
        "abstract": "The HH\\,24-26 star forming region within the Lynds 1630 (L1630) dark cloud in Orion contains a remarkable collection of rare Class~0 and Class~I protostars, collimated molecular and ionized jets, and a luminous but spatially unresolved ASCA X-ray source. To study the X-ray properties of the embedded protostar population of that region, we  have obtained a deep X-ray image with the ACIS-S camera on board the {\\it Chandra X-Ray Observatory}.  A number of H$\\alpha$ emission-line objects were detected in the areas surrounding HH\\,24--26, of which the weak-line T~Tauri star SSV\\,61 was the brightest source, at a steady luminosity of $L_x$(0.3--10 keV)$ = 10^{31.9}$ \\ergs. Two Class~I protostars aligned with optical jets in HH\\,24, SSV\\,63E and SSV\\,63W, were also detected, as was the continuum radio source SSV\\,63NE, which is very likely an extreme Class~I or Class~0 object. We observed no X rays from the Class~0 protostars HH\\,24-MMS and HH\\,25-MMS, nor any from regions of the cloud bounded by HH\\,25 and HH\\,26, at a $2\\sigma$ upper limit of $L_x \\sim 10^{30.0}$ \\ergs. HH\\,26-IR, the Class~I object thought to be the origin of the HH\\,26 flow, was not detected. Near-infrared spectroscopy obtained at the NASA IRTF reveals 3\\,$\\micron$ ice bands in the spectra of SSV\\,59, 63E, 63W, and HH\\,26-IR, and 2.3\\,$\\micron$ CO overtone absorption bands for SSV\\,61. SSV\\,60, which lies astride one end of the great arc of nebulosity forming HH\\,25, exhibits a deep infrared ice band and CO absorption, but is not an X-ray source, and is most likely a distant background giant of late spectral type. ",
        "introduction": "The Lynds 1630 (L1630) dark cloud is part of the Orion B Molecular Cloud.  At a distance of approximately 450 pc, it is one of the nearest giant molecular clouds.  Its mass is estimated to be more than $\\sim1 \\times 10^5$ \\msun\\  (Maddalena et al.  1986).  L1630 shows many signs of on-going star formation: ({\\it a}) dense cores of molecular gas and high velocity outflows (Snell \\& Edwards 1982; Gibb \\& Heaton 1993; Gibb et al. 1995); ({\\it b}) bright reflection nebulae (e.g., NGC 2068 and NGC 2071); ({\\it c}) T Tauri stars (Herbig \\& Kuhi 1963; Wiramihardja et al. 1989); ({\\it d}) Herbig-Haro (HH) objects, including the morphologically complex HH\\,24--26 (Herbig 1974; Solf 1987; Mundt, Ray, \\& Raga 1991; Eisl\\\"{o}ffel \\& Mundt 1997); ({\\it e}) luminous far-infrared and submillimeter sources (Cohen et al. 1984; Cohen \\& Schwartz 1987; Gibb \\& Heaton 1993; Schwartz, Burton, \\& Herrmann 1997; Mitchell et al. 2001); ({\\it f}) radio continuum sources (Bontemps, Andr\\'{e}, \\& Ward-Thompson 1995; Verdes-Montenegro \\& Ho 1996; Reipurth et al. 2002); and ({\\it g}) dense clusters of optically obscured near-infrared sources (Strom et al. 1975; Strom, Strom, \\& Vrba 1976). The number density of young stars in L1630, 450 pc$^{-3}$, is comparable with that of the $\\rho$~Ophiuchus star-forming region (Lada et al. 1991).  Starting with the work of the Stroms, L1630 has been studied in ever increasing detail at infrared wavelengths. Particular attention has been focused recently on its many outflows and HH features.  Evidence for emission in the 2\\,$\\micron$ near-infrared lines of H$_2$ from parsec-scale HH flows and jets has been presented by Zealey et al. (1992), and more extensively by Davis et al. (1997) and Chrysostomou et al. (2000). Images of HH\\,24 in narrow emission lines at optical and infrared wavelengths trace out at least three intersecting, collimated, bipolar flows (e.g., Mundt et al. 1991:  Fig. 17--19; Davis et al. 1997:  Fig. 2). The flows emanate from the vicinity of HH\\,24B, at the location of the bright near-infrared sources SSV\\,63E and SSV\\,63W. Both objects are presumed from their steep spectral energy distributions to be Class I protostars (Cohen et al. 1984), that is, they are in a very early stage of evolution preceding the T Tauri star (TTS) or so-called Class II phase, and in the process of actively accreting material from their protostellar disks. The dynamical age of HH\\,24 is of the order of $\\sim10^4$~yr (Eisl\\\"{o}ffel \\& Mundt 1997; Snell \\& Edwards 1982), which implies that the Class I objects powering those flows must be very young, and also in a very short-lived phase of evolution.  In the past, the bright HH\\,24A knot was interpreted as the ``working surface'' or bow shock for one of those outflows (Mundt et al. 1991; Solf 1987), but it is now recognized (e.g., Eisl\\\"{o}ffel \\& Mundt 1997) to be the terminus of an independent outflow, which originates in a region immediately to the southwest, at the location of a deeply embedded Class~0 protostar, HH\\,24-MMS. Class~0 objects like HH\\,24-MMS represent the main accretion phase for protostars and are extremely rare because their lifetimes are so short. HH\\,24-MMS is observed directly at near-infrared wavelengths only in shocked H$_2$ (Bontemps, Ward-Thompson, \\& Andr\\'{e} 1996; Davis et al. 2002). It was discovered as a bright source at submillimeter and radio wavelengths (Chini et al. 1993; Ward-Thompson et al. 1995; Bontemps et al. 1995). Greybody fits suggest a color temperature of 20~K and a total luminosity of $\\sim$5--20 \\lsun, of which $\\sim2.5$ \\lsun\\ originates from an unresolved and possibly disk-like central source (Ward-Thompson et al. 1995). The compact source is surrounded by a much more extended envelope, from which it is most likely accreting material at a very high rate (Chandler et al. 1995b). The detection of an H$_2$ jet, blue-shifted CO emission (Bontemps et al. 1996), and indications of disk structure suggest that HH\\,24-MMS, even at its remarkably young age, has already begun to drive a powerful collimated outflow into the surrounding molecular cloud.  Various theories and mechanisms have been advanced in the past to explain such highly focused (and frequently bipolar) jets. In the class of magnetic models, for example, it has been proposed that the outflow is confined and driven by interactions between the magnetosphere of the protostar and the primordial field that threads its accretion disk (e.g., Shu et al. 1994).  South of HH\\,24, high resolution radio, millimeter, and submillimeter maps have revealed a number of strong molecular peaks and bipolar outflows in the proximity of HH\\,25 and HH\\,26 (Gibb \\& Heaton 1993; Gibb et al. 1995; Gibb \\& Davis 1998).  The surrounding areas have been extensively mapped in the $v = 1-0\\ S$(1) line of H$_2$ at 2\\,$\\micron$ (Davis et al. 1997; Chrysostomou et al. 2000; Davis et al. 2002). A Class~0 object, HH\\,25-MMS, was discovered $10\\arcsec$ south of the optically-visible knot HH\\,25A as both a 1.3-mm and a 3.6-cm source (Bontemps et al. 1995). Strong 450 and 850\\,$\\micron$ continuum emission has been observed from the same location (Mitchell et al. 2001; Phillips, Gibb, \\& Little 2000). The brightest near-infrared objects in the vicinity of HH\\,25--26 are SSV\\,59, SSV\\,60, and HH\\,26-IR (Strom et al. 1976; Davis et al. 1997). None of the three has been shown conclusively to be the origin of a large-scale flow, although the weight of the evidence for HH\\,26 (in particular, the discovery of a compact H$_2$ jet:  Davis et al. 2002) now points most convincingly to HH\\,26-IR. The latter, as well as SSV\\,59, is generally regarded as a Class~I protostar (e.g., Bontemps et al. 1995, Davis et al. 2002). The status of SSV\\,60 is unknown; unlike the other two sources, it is associated with neither a submillimeter peak (Lis, Menten, \\& Zylka 1999; Mitchell et al. 2001; Phillips et al. 2001) nor a molecular core (e.g., Torrelles et al. 1989; Gibb \\& Heaton 1993).  The same process that drives the outflows and jets of Young Stellar Objects (YSOs) is expected to result in the repeated shearing and reconnection of  the magnetic field lines that attach such a star to its circumstellar disk, and thereby to generate a substantial amount of X-ray emission  (Hayashi, Shibata, \\& Matsumoto 1996; Shu et al. 1997; Goodson, Winglee, \\& B\\\"{o}hm 1997). The numerical simulations of Hayashi et al. and Goodson et al. in particular demonstrated that as a consequence of reconnection events some of the magnetically entrapped gas that is close to an accreting YSO (not necessarily its outflowing jet) will reach temperatures as high as 100 MK and will be flung outwards, away from the star. Such high temperatures should give rise to hard X-ray spectra. X-ray emission at equally high temperatures can also be produced in an accretion shock at the footpoints of the magnetic field, where infalling material from the disk streams down onto the surface of the forming star (Uchida \\& Shibata 1984). The HH\\,24--26 region was observed in X-rays with ASCA for a net exposure time of 30 kiloseconds (ks) by Ozawa et al. (1999). The strongest source in the soft X-ray band (0.7--2.0 keV) in that observation was SSV\\,61, a T~Tauri star (TTS) which is embedded in an unusual patch of nebulosity situated mid-way on the sky between HH\\,24 and HH\\,25--26 (e.g., Scarott, Gledhill, \\& Warren-Smith 1987). Strom et al. (1975) classified the star as an M2--M4 giant from image-tube spectra. A CCD spectrum of SSV\\,61,  kindly made available to us by G. Herbig, exhibits H$\\alpha$ emission and \\ion{Li}{1} $\\lambda6707$ absorption equivalent widths that are typical of the class of weak-lined T Tauri stars (wTTS), EW(H$\\alpha$) = $-5$ \\AA\\ and EW(\\ion{Li}{1}) = 480 m\\AA\\ (cf. Walter 1986; Mart\\'{\\i}n et al. 1994). The brightest emission in the hard ASCA band (2.0--8.0 keV) is located north of SSV\\,61, in the vicinity of HH\\,24. The high-energy peak lies nearest SSV\\,63, but cannot be identified unambiguously with either of its infrared components due to the low spatial resolution of ASCA. Ozawa et al. noted the great difficulty of separating the X-ray flux of SSV\\,63 from that of SSV\\,61, given their offset of just $1\\farcm8$, at roughly the half-power radius of ASCA.  The quiescent luminosity $L_x \\simeq 10^{32}$ \\ergs\\ they derive for SSV\\,63 places it above the top end of the range thus far observed for all but a few Class I objects, by factors of 50--100 (Carkner, Kozak, \\& Feigelson 1998; Preibisch, Neuh\\\"{a}user, \\& Stanke 1998). Whether SSV\\,63 is truly one of the most X-ray luminous YSOs in the sky, however, has awaited confirmation from high resolution imaging that resolves not only SSV\\,61 from SSV\\,63, but also the individual infrared components of SSV\\,63 itself. That task, including the requisite high precision astrometry, has now been accomplished by the {\\it Chandra} observations reported here.  As for the Class~0 sources in L1630, Ozawa et al. found no sign of X-ray emission at the position of HH\\,24-MMS, located in the saddle between the soft and hard X-ray peaks in their image, nor any emission at the position of HH\\,25-MMS, which was also in the ASCA field of view. Their $3\\sigma$ upper limit of $L_x \\lesssim 1.2\\times 10^{31}$ \\ergs\\ is not a stringent limit on the emission, however, since YSOs elsewhere in Orion and in other star forming regions have been detected at much fainter levels, $L_x \\simeq \\mbox{few} \\times 10^{30}$ \\ergs\\ (e.g., Feigelson et al. 2002; Imanishi, Koyama, \\& Tsuboi 2001). Using the {\\it Chandra Observatory}, we have undertaken a far more sensitive search for the onset of X-ray emission in the Class 0 objects within the L1630 cloud. Our objective was to establish an evolutionary progression in X-ray properties within a single cloud, and to examine the possible relationship between protostellar X rays, accretion disks, magnetically confined jets, and large-scale molecular outflows.  Taking advantage of the lower background, greater sensitivity, and unsurpassed spatial resolution of {\\it Chandra}, we are now able to set a 10 times smaller bound on possible X-ray emission from the three established Class 0 protostars in the L1630 cloud, HH\\,24-MMS, HH\\,25-MMS, and LBS\\,17-H. ",
        "conclusions": "We have acquired a deep X-ray image of the HH\\,24--26 region in the L1630 cloud with the ACIS-S camera on {\\it Chandra}, which has resulted in the detection of a number of anonymous 2MASS infrared sources and optically visible H$\\alpha$ emission-line stars, including the X-ray luminous coronal source, the wTTS SSV\\,61. Also detected at what appear to be their quiescent levels of activity are two deeply embedded Class~I protostars (SSV\\,63E and W), which have been identified in earlier work as the source of energetic molecular and ionized flows in the region. The two were blended together in an earlier detection by ASCA, but are well separated at the much higher spatial resolution of {\\it Chandra}. Two other HH-associated Class~I objects (SSV\\,59 and HH\\,26-IR) as well as three Class~0 YSOs in the same cloud (HH\\,24-MMS, HH\\,25-MMS, and LBS 17-H) were not detected. Although X-ray observations of several different star forming regions thus far reveal a striking lack of Class~0 detections, no firm conclusions can be drawn from our small sample of YSOs regarding a possible evolutionary trend in the intrinsic X-ray luminosities of the Class~0 and Class~I protostars. X-ray emission was detected from a deeply embedded continuum radio source (HH\\,24NE), whose location within HH\\,24 suggests that it may be the highly obscured driving source for one of the long HH filaments in this area. The object has no visible or near-infrared counterpart, and therefore is potentially a Class~0 or extreme Class~I object.\\\\    This research has made use of the {\\sc simbad} database, operated at {\\it CDS}, Strasbourg, France, and is based in part on observations made with the \\chandra. Support for this work was provided by \\nas\\ through Chandra Award Number GO2-3007X issued by the \\chandra\\ Center, which is operated by the Smithsonian Astrophysical Observatory for and on behalf of \\nas\\ under contract NAS8-39073. The Digitized Sky Surveys (DSS) were produced at the Space Telescope Science Institute under NASA Grant NAG W-2166 and are based on photographic data obtained using the Oschin Schmidt Telescope on Palomar Mountain and the UK Schmidt Telescope with the permission of those institutions. The Guide Star Catalogue-II is a joint project of the Space Telescope Science Institute and the Osservatorio Astronomico di Torino. Space Telescope Science Institute is operated by the Association of Universities for Research in Astronomy, for the National Aeronautics and Space Administration under contract NAS5-26555. The participation of the Osservatorio Astronomico di Torino is supported by the Italian Council for Research in Astronomy. Additional support is provided by  European Southern Observatory, Space Telescope European Coordinating Facility, the International GEMINI project and the European Space Agency Astrophysics Division.  The Two Micron All Sky Survey (2MASS) is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center (IPAC)/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation.   \\newpage"
    },
    "0404/astro-ph0404056_arXiv.txt": {
        "abstract": "We have calculated the radiation field, dust grain temperatures, and far infrared emissivity of  numerical models of turbulent molecular clouds. When compared to a uniform cloud of the same mean optical depth, most of the volume inside the turbulent cloud is brighter, but most of the mass is darker. There is little mean attenuation from center to edge, and clumping causes the radiation field to be somewhat bluer. There is also a large dispersion, typically by a few orders of magnitude, of all quantities relative to their means.  However, despite the scatter, the 850$\\mu$m emission maps are well correlated with surface density. The fraction of mass as a function of intensity can be reproduced by a simple hierarchical model of density structure. ",
        "introduction": "Observations of molecular clouds indicate that they are inhomogeneous in their internal structures. Although the mean extinction through such a cloud may be quite high, the extinction along selected lines of sight can be quite low. The result is that even dense gas in a starless cloud can find itself brightly illuminated by the ambient Galactic radiation field, with consequences for grain heating, ionization balance, and photochemistry (e.g. Spitzer 1978, Tielens \\& Hollenbach 1985) within the cloud. Although astrophysicists have long been aware that extinction in clouds is highly variable - Chandrasekhar \\& M\\\"unch (1950) proposed that the statistics of extinction could be used  as a probe of interstellar turbulence - the relative paucity of detailed cloud models has made it infeasible to explore these effects and their observational consequences. Only in the past few years have  dynamical models of turbulent molecular gas, based on numerical simulations, become available and some studies of extinction and radiative transport in the context of these models have been carried out (Padoan \\& Nordlund 1999, Juvela \\& Padoan 2003).  In this paper we use such models to explore the range of mean intensities, $J_\\lambda$, in clumpy clouds exposed to an interstellar radition field (ISRF) appropriate to the solar neighborhood.  Although the principle application of our results is to show the effects of a clumpy gas distribution on grain temperatures and far-infrared emissivity, the general problem of how radiation penetrates a clumpy, dusty medium is important in other situations in astrophysics such as protostellar disks (Wolf, Henning \\& Stecklum 1999), the dusty interstellar medium (Witt \\& Gordon 2000), the Galactic center (Morris \\& Serabyn 1996), and AGN (Krolik 1998).  Because dynamical models are not always available or easy to use, we attempt to identify features of our results which can be reproduced by a simpler model, namely a hierarchical (fractal) density distribution.  In \\S 2, we introduce the cloud models, comparing their basic physical attributes, as well as the representative dust mixture and interstellar radiation field used in the calculations. In \\S 3, we describe the approach used to calculate the penetration of an isotropic, monochromatic ambient radiation field into a molecular cloud. We decribe tests of the numerical technique and consider resolution effects due to the model clouds. Section 4 concerns the results of radiative transfer calculations applied to the model clouds. We show that there is little diminution of mean radiative intensity from the edge of the cloud to its center, and that although there is a correlation between local gas density and local radiation field, there is a large dispersion about the mean.  In principle our results are easily extended to point sources of radiation within the clouds (Natta et al. 1981).  However, lacking a dynamical model which includes star formation self consistently we have chosen not to do so.  In advance of our grain temperature calculations in \\S 5 we consider the distribution of colors in the intracloud radiation field.  Section 5 deals with the detailed calculation of grain temperatures, as well as the resuling far-infrared emission, for a spectrum of grain sizes exposed to a representation of the interstellar radiation field propagated through the model clouds. Following the results of \\S 4, we find only a weak correlation between the infrared spectrum and local conditions in the cloud, although the detailed consideration of grain temperatures in \\S 5 leads us to the conclusion that the 850$\\mu$m continuum surface brightness is very well correlated with column density, considerably more so than maps made at 100$\\mu$m.  Although these results are obtained for the particular grain models of Draine \\& Lee (1984, hereafter 'DL84') we believe they are qualitatively robust. Section 6 is a summary and discussion. ",
        "conclusions": "We have studied the attenuation of the interstellar radiation field by dust in models of clumpy clouds and calculated the dust temperatures and far-IR emissivities. The main conclusions are as follows:  1. Inhomogeneity of the density continuum dramatically increases the radiant energy that penetrates into the cloud, making the \\textit{volume} markedly brighter when compared to a uniform cloud of comparable optical depth (Fig. 4c).  The mass contained in clumpy structures provides enough local attenuation that most of the \\textit{mass} is in fact associated with lower intensities (Fig. 4b) when compared to the uniform cloud.  2.  The above two effects imply an anti-correlation of average $J$ with $n_H$, but from point to point there exists an extremely large scatter about this trend, more so at wavelengths at which the cloud is optically thick (Figs. 4a and 5). Despite this large scatter, the mean intensity \\textit{averaged} by spherical shells is largely independent of radial position within the inner 50\\% of the cloud volume (Fig. 6).  3.  The cloud's inhomogeneity allows small amounts of weakly unattenuated blue ISRF to penetrate the cloud which, when combined with the predominant reddened field, \\textit{makes the intracloud field somewhat bluer than in the uniform cloud}.  The variability of this effect introduces a scatter in the color (Fig. 7) which in accordance with $J$ is also on average quite uniform throughout the model (Fig. 8).  4. The differential mass per intensity distribution of the MHD models can be qualitatively reproduced by fractal clouds generated according to a simple prescription and augmented by a uniform density background. In both types of inhomogeneous model, the peak of the mass distribution is shifted occurs at lower intensity than in a uniform cloud, suggesting that this is a robust feature of clumpiness (Figure 9). These results are relatively insensitive to fractal dimension or magnetic fieldstrength over the range explored ($D$ = 2.3 - 2.6; $\\beta$ = 0.05 - 4.04).    5. For one grain model (DL84) bathed in the attenuated ISRF, the sensitivity of dust grains of different radii to the overall spectral form (i.e color) of the intracloud radiation field $J_{\\lambda}$ produces average grain temperatures which do not depend strongly on the grain size (Table 2).  \\textit{Importantly, the smaller grains, absorbing the most variable part of the intracloud radiation field, exhibit a greater scatter in their temperatures}.  When considering how the dust mass is distributed with temperature there is a slight preponderance for low temperature ($T\\sim 10$K) material (Fig. 10), primarily due to the abundance of silicates in the DL84-MRN grain ensemble.  6. Small, relatively hot grains emit the majority of the $100\\mu$m emission, larger grains emitting longward of this (Fig. 11).  The temperature variations in these two grain populations, considered in the context of the Rayleigh-Jeans and Wien spectral regimes, yield emissivities $j_{\\lambda,H}$ exhibiting very different point to point variations, shown in Figure 12.  The $850 \\mu$m emissivity per H atom varies by less than a factor of 5 whereas the $100 \\mu$m emissivity per H atom varies by several orders of magnitude.  7.  In constructing brightness maps (Fig. 13) the intrinsic scatter in $j_{\\lambda,H}$, the emissivity per H atom, that ultimately arises from the cloud's inhomogeneity increasingly decorrelates maps made at $\\lambda<100\\mu$m when compared with the surface density.  Conversely, maps made at longer wavelengths show increasingly higher correlations between brightness and surface density maps, and should be preferred if one wishes to infer surface densities and cloud masses from brightness maps."
    },
    "0404/astro-ph0404326_arXiv.txt": {
        "abstract": "We investigate the broadband optical emission of diffuse intergalactic stars in galaxy clusters at $z\\sim0.25$ by means of an image stacking technique. The images of 654 clusters, selected with the max-BCG algorithm from a subsample of the SDSS-Data Release 1 (DR1), have been stacked after masking all the sources detected down to very low S/N. The resulting images in the g, r and i bands provide reliable photometric data at the level of $\\gtrsim30$~mag$/\\square$\" (r band), out to $\\gtrsim 600$~kpc from the Brightest Cluster Galaxy (BCG). \\\\Our analysis shows that:\\\\ i) the IntraCluster Light (ICL) is much more concentrated than the galaxy light, contributing $\\sim$30\\% of the total cluster optical emission at $R=100$~kpc, but less than 10\\% at $R>500$~kpc;\\\\ ii) the ICL contributes between 15 and 20\\% of the total cluster optical luminosity between the optical radius ($\\mu_r=25$~mag$/\\square$\") of the BCG and 500 kpc;\\\\ iii) the colours of the ICL are consistent with the global colours of the cluster galaxies, with little evidence for redder g-r. ",
        "introduction": "Firstly proposed by \\cite{1951PASP...63...61Z}, the presence of a diffuse population of intergalactic stars in galaxy clusters is now a well established observational fact (see e.g. \\cite[Arnaboldi 2003]{2003IAUS..217E..20A} and references therein, Krick \\& Bernstein this conference). Such a population is commonly believed to originate from the disruption of small galaxies and from the tidal stripping of the outskirts of large galaxies in high-density environments. Studying the properties of the ICL as a function of redshift and of the richness and morphology of the parent cluster can give a significant insight in the interaction processes that drive galaxy evolution in dense environments.\\\\ Observing the intracluster light (ICL) is extremely challenging, due to the very low surface brightness involved ($>26$~mag$/\\square$\"), which is less than 1\\% of the typical surface brightness of the sky. During the last years the increasing sensitivity of CCD detectors and the development of new observational techniques has made it possible to start a number of optical surveys aiming at the study of the ICL, by means of both broadband imaging of unresolved stellar populations (e.g. \\cite[Feldmeier \\etal~2002]{2002ApJ...575..779F}) and of narrow band detection of intracluster planetary nebulae (e.g. \\cite[Arnaboldi \\etal~2003]{2003AJ....125..514A}). Nevertheless, for a statistically representative number of clusters of different morphology and richness and located at different redshifts we still lack reliable determinations of the basic properties of the ICL, such as the fractional amount with respect to the total optical emission of the cluster, the radial distribution, the link with (possible) cD envelopes, or the colours. In fact, observations of the ICL in individual clusters are extremely demanding in telescope time and require an incredibly high control of the flat fielding and of the sources of light pollution, such as fore- and background sources and the internal reflections of the camera.\\\\ We have chosen to tackle the problem with a purely statistical approach, adopting an image stacking technique, similar to that used by \\cite{2004MNRAS.347..556Z}. The Sloan Digital Sky Survey (SDSS, \\cite[York \\etal~2000]{2000AJ....120.1579Y}) is providing a one-quarter-of-sky coverage with 5-band imaging (\\cite[Fukugita \\etal~1996, Gunn \\etal~1998]{1996AJ....111.1748F,1998AJ....116.3040G}) and spectroscopy. From a subsample of its first data release (DR1, \\cite[Abazajian \\etal~2003]{2003AJ....126.2081A}), we have extracted a catalogue of 654 clusters and rich groups in the redshift range 0.2-0.3, whose imaging in the g, r, and i bands was processed and stacked, as described in Sec.\\ref{stacking}. The results are presented and discussed in Sec.\\ref{results}. ",
        "conclusions": ""
    },
    "0404/astro-ph0404110_arXiv.txt": {
        "abstract": "We present an initial sample of 19 intermediate-mass black hole candidates in active galactic nuclei culled from the First Data Release of the Sloan Digital Sky Survey.  Using the linewidth-luminosity-mass scaling relation established for broad-line active nuclei, we estimate black hole masses in the range of $M_{\\mathrm{BH}} \\approx 8 \\times 10^4 - 10^6 \\, M_{\\sun}$, a regime in which only two objects are currently known. The absolute magnitudes are faint for active galactic nuclei, ranging from $M_g \\approx -15$ to $-18$ mag, while the bolometric luminosities are all close to the Eddington limit.  The entire sample formally satisfies the linewidth criterion for so-called narrow-line Seyfert 1 galaxies; however, they display a wider range of \\ion{Fe}{2} and [\\ion{O}{3}] $\\lambda$5007 line strengths than is typically observed in this class of objects. Although the available imaging data are of insufficient quality to ascertain the detailed morphologies of the host galaxies, it is likely that the majority of the hosts are relatively late-type systems.  The host galaxies have estimated $g$-band luminosities $\\sim 1$ mag fainter than $M^*$ for the general galaxy population at $z \\approx 0.1$.  Beyond simply extending the known mass range of central black holes in galactic nuclei, these objects provide unique observational constraints on the progenitors of supermassive black holes.  They are also expected to contribute significantly to the integrated signal for future gravitational wave experiments. ",
        "introduction": "Dynamical studies have established the existence of supermassive black holes (BHs) with masses $M_{\\rm BH} \\approx 10^6-10^9$~$M_{\\sun}$ in the centers of most, if not all, local galaxies with bulges (Magorrian et al.  1998; Richstone 2004).  A significant challenge to any model of cosmological BH growth is the nature of seed BHs.  Observations of intermediate-mass BHs ($M_{\\rm BH} \\approx 10^3-10^6$~$M_{\\sun}$) in the local Universe would provide the most direct empirical constraints on such seeds.  The \\msigma\\ relation (Ferrarese \\& Merritt 2000; Gebhardt et al. 2000a; Tremaine et al.  2002) strongly indicates feedback in the coevolution of galaxies and BHs, but it is still an open question whether the relation is established at early or late times.  If, as some models suggest (e.g., Di~Matteo et al. 2003), the \\msigma\\ relation is established late in the evolution of a galaxy, then we may expect to see deviations from this relation in a population of intermediate-mass BHs that is not yet fully grown. There are other motivations to characterize the lowest end of the nuclear BH mass distribution.  Certainly their number and mass distribution will impact the expected integrated background in gravitational radiation (e.g., Hughes 2002).  Presumably they will pass through an active phase of accretion and contribute at some level to the observed integrated background radiation, particularly at X-ray energies.  The range of accretion phenomena they display will depend on the major accretion modes for BHs of such low mass.  The accretion states and, by extension, radiative growth mechanisms of this population will therefore have direct implications for the growth history of primordial BHs.  Apart from their physical implications, intermediate-mass BHs have practical value in that they offer tremendous leverage in anchoring local BH-galaxy scaling correlations such as the \\msigma\\ relation.  Recently discovered intermediate-mass BH candidates in galactic nuclei offer tantalizing hints as to the nature of this population.  NGC 4395 is a very late-type (Sdm) spiral with no bulge component, whose central stellar velocity dispersion is \\sigmastar$<30$ \\kms\\ (Filippenko \\& Ho 2003).  Nevertheless, it has the emission properties of a type 1 active galactic nucleus (AGN), with broad permitted optical and UV emission lines (Filippenko \\& Sargent 1989; Filippenko, Ho, \\& Sargent 1993), a compact radio core (Ho \\& Ulvestad 2001) with a nonthermal brightness temperature (Wrobel, Fassnacht, \\& Ho 2001), and a pointlike, hard X-ray source (Ho et al. 2001) that is highly variable (Moran et al. 1999, 2004; Shih, Iwasawa, \\& Fabian 2003).  Mass estimates based on the \\hbeta\\ linewidth-luminosity-mass scaling relation and X-ray variability suggest a BH mass of $10^4 - 10^5$~\\msun\\ (Filippenko \\& Ho 2003).  Interestingly, these agree with the limit of $< 10^5$~\\msun\\ derived from the \\msigma\\ relation (Filippenko \\& Ho 2003).  POX 52, on the other hand, has a dwarf elliptical host.  Recently revisited by Barth et al. (2004), POX 52 was first observed by Kunth, Sargent, \\& Bothun (1987), who noted the Seyfert-like characteristics of its optical spectrum, including a broad component to the \\hbeta\\ emission line.  The galaxy has a central velocity dispersion of 36~\\kms, which yields a mass estimate of $M_{\\rm BH} = 1.4 \\times 10^5$~\\msun, again consistent with the value of $1.6 \\times 10^5$ \\msun\\ derived from the \\hbeta\\ linewidth-luminosity-mass scaling relation.  The [\\ion{O}{3}] $\\lambda$5007 linewidth, $\\sigma~\\approx~37$ \\kms, is also fully consistent with the stellar velocity dispersion (Barth et al. 2004).  Unfortunately, it is currently technically impossible to obtain direct mass measurements for $M_{\\rm BH} \\lesssim 10^6$ \\msun, because of our inability to resolve the BH sphere of influence, for all but the nearest galaxies in the Local Group (e.g., M33; Gebhardt et al. 2001). The only direct mass determinations for intermediate-mass BHs come from observations of massive star clusters (Gebhardt, Rich, \\& Ho 2002; Gerssen et al. 2002), which, in any case, have been controversial (Baumgardt et al. 2003a, b).  Barring evidence from spatially resolved kinematics, the best indicator for nuclear BHs in nearby galaxies comes from AGN activity as revealed through optical emission-line intensity ratios (Baldwin, Phillips, \\& Terlevich 1981; Veilleux \\& Osterbrock 1987; Ho, Filippenko, \\& Sargent 1993), kinematically distinct, broad emission lines (e.g., Ho et al. 1997c), or the presence of compact, nonstellar radio and X-ray cores (e.g., Ho et al. 2001; Nagar et al. 2002; Terashima \\& Wilson 2003; see Ho 2004a for a review).  Thus, the best hope to find intermediate-mass BHs is through an AGN survey.  However, there are at least two practical challenges.  First, it may be that objects like NGC 4395 are quite rare.  In the Palomar survey of $\\sim 500$ nearby galaxies, NGC 4395 is the only example of its kind found (Ho, Filippenko, \\& Sargent 1997a, b).  POX 52 was discovered serendipitously in a fairly limited objective prism survey for emission-line objects (Kunth et al.  1987).  These two are the only known cases of intermediate-mass BHs in galactic nuclei.  Thus, to accumulate a significant sample will require a much larger survey. Second, even radiative signatures of accretion in these objects will be challenging to uncover.  The host galaxies are expected to be either very late-type systems with low mass, and therefore intrinsically faint, or else regular, luminous disk-dominated galaxies with a tiny or nonexistent bulge.  Even at a high accretion rate a $10^5$~\\msun\\ BH will itself have a low luminosity.  As a result, detection of the faint AGN will be limited by sensitivity and host galaxy contamination.  A large-area, sensitive, and uniform optical spectroscopic galaxy survey, such as the Sloan Digital Sky Survey (SDSS; York et~al. 2000), offers the best opportunity for finding a significant number of new intermediate-mass BH candidates.  We present the first sample of candidates culled from the SDSS.  We discuss the data in \\S 2, the sample definition and analysis in \\S 3, and the main results in \\S 4.  In this paper we assume the following cosmological parameters to derive distances: $H_0 = 100h = 72$~\\kms~Mpc$^{-1}$, $\\Omega_{\\rm m} = 0.3$, and $\\Omega_{\\Lambda} = 0.7$. ",
        "conclusions": "We employ a simple technique to exploit the breadth of the SDSS galaxy spectroscopy to extend the demography of central BHs at least 1 order of magnitude below the $10^6$ \\msun\\ threshold, a regime hardly explored previously.  We present the first sizable sample of AGNs with candidate intermediate-mass BHs.  The sample contains 19 objects with virial BH masses ranging from $M_{\\rm BH} \\approx 8\\times 10^4$ to $10^6$ \\msun\\ (Fig. 5).  We now briefly examine some of the properties of these objects.  \\subsection{Spectral Properties}  In the literature to date, the closest analog to our sample (aside from NGC 4395 and POX 52) is the class of AGNs known as narrow-line Seyfert 1s (NLS1), which are thought to possess relatively low-mass BHs emitting at a high fraction of their Eddington rates (e.g., Boroson 2002). Since our sample selection required a relatively prominent AGN component relative to the host galaxy, the sample should preferentially single out AGNs in a relatively ``high'' state.  First classified as a group by Osterbrock \\& Pogge (1985), NLS1s are typically characterized by broad permitted lines with FWHM $\\lesssim 2000$ \\kms, prominent \\ion{Fe}{2} emission, and relatively weak [\\ion{O}{3}] $\\lambda \\lambda$4959, 5007 lines.  NLS1s also tend to emit copiously in the soft X-ray band (e.g., Boller, Brandt, \\& Fink 1996; but see Williams, Pogge, \\& Mathur 2002), presumably as a consequence of their high accretion rates (Pounds, Done, \\& Osborne 1995).  Our sample conforms to some, but not all, of these expectations.  By the linewidth criterion, certainly all of the objects would technically qualify as ``narrow-line'' sources: the broad \\halpha\\ component has FWHM ranging from 464 to 1730 \\kms, with an average value of 836 \\kms. However, we note that our sample shows a greater diversity of \\ion{Fe}{2} and [\\ion{O}{3}] strengths than in previous samples of NLS1s.  \\ion{Fe}{2} emission is common, but not ubiquitous\\footnote{The same comment applies to NGC 4395 and POX 52, both of which may be formally regarded as NLS1s.  Fe~{\\tiny II} multiplets can be seen in the UV spectrum of NGC 4395 (Filippenko et al. 1993), but they are quite weak in its optical spectrum (Filippenko \\& Sargent 1989).  Fe~{\\tiny II} emission is also not prominent in the optical spectrum of POX 52 (Barth et al. 2004).}.  For example, the intensity ratio of \\ion{Fe}{2} $\\lambda$4570 relative to total (broad + narrow) \\hbeta, where \\ion{Fe}{2} $\\lambda$4570 is the complex between 4434 and 4680 \\AA, varies from $\\sim$0.465 to 3.17, with a mean of $1.32 \\pm 0.16$, calculated using the Kaplan-Meier product-limit estimator to account for upper limits (Feigelson \\& Nelson 1985). This is to be compared with $\\langle$\\ion{Fe}{2} $\\lambda$4570/\\hbeta(total) $\\rangle$ = $0.67\\pm0.04$ for the 56 NLS1s presented by V\\'eron-Cetty, V\\'eron, \\& Gon{\\c c}alves (2001). Similarly, in our sample [\\ion{O}{3}]/\\hbeta(total) = 0.335 to 12.0, with a mean of $2.24 \\pm 0.72$.  In contrast V\\'eron-Cetty et al. (2001) find $\\langle$[\\ion{O}{3}]/\\hbeta(total)$\\rangle$ = $0.65 \\pm 0.09$, and Constantin \\& Shields 2003 find $\\langle$[\\ion{O}{3}]/\\hbeta(total)$\\rangle$ = 0.39 for their composite spectrum of 22 NLS1s. Thus, compared to classical NLS1s, we find stronger [\\ion{O}{3}] and \\ion{Fe}{2} emission. The distribution of power-law slopes for the AGN-dominated spectra are also quite interesting (Table 3). While $1.5 < \\beta < 1$ is a standard range for luminous AGNs (e.g., Malkan 1988), we find an average value of $\\beta = 0.42$.    \\begin{figure*} \\begin{center} \\epsfig{file=f6_120.eps,angle=-90,width=0.9\\textwidth,keepaspectratio=true} \\vskip -0.1mm \\figcaption[f6.eps]{ Diagnostic diagram plotting log [O~{\\tiny III}] $\\lambda$5007/\\hbeta\\ versus log [N~{\\tiny II}] $\\lambda$6583/\\halpha\\ ({\\it a}) and versus log [S~{\\tiny II}] $\\lambda\\lambda$6716, 6731/\\halpha\\ ({\\it b}).  The line ratios have not been corrected for reddening, but this should not matter because of the close wavelength separation of the lines. Object \\#15 was omitted because [N II], [S II], and narrow H$\\alpha$ were not detected.  The dotted lines mark the boundaries of the three main classes of emission-line nuclei, according to the convention of Ho et al. (1997a): ``H'' = H~{\\tiny II} nuclei, ``S'' = Seyferts, and ``L'' = LINERs. \\label{fig6}} \\vskip -5mm \\end{center} \\end{figure*}     \\begin{figure} \\begin{center} \\epsfig{file=f7_120.eps,width=2.8in,height=5.6in} % \\vskip -0.1mm \\figcaption[f7.eps]{ Distributions of $g$-band absolute magnitudes.  The magnitudes pertain to the entire system ({\\it top}), the host galaxy alone ({\\it middle}), and the AGN alone ({\\it bottom}). The AGN contribution was estimated from \\lf, which was then converted to a $g$ magnitude assuming an $f_\\lambda \\propto \\lambda^{-\\beta}$ spectrum, where $\\beta$ is shown in Table 3. \\label{fig7}} \\vskip -5mm \\end{center} \\end{figure}    \\begin{deluxetable}{rcccc} \\tablecolumns{5} \\tablewidth{0pc} \\tabletypesize{\\scriptsize} \\tablecaption{ X-ray Detections\\label{tablexr}} \\tablehead{ \\colhead{ID} & \\colhead{$C$} & \\colhead{log $N_{\\mathrm{H}}$} & \\colhead{log $f_{\\rm X}$} & \\colhead{log $L_{\\rm X}$}\\\\ \\colhead{(1)} & \\colhead{(2)} & \\colhead{(3)} & \\colhead{(4)} & \\colhead{(5)} }  \\startdata  1 & 0.0483 & 20.59 & $-$12.21 & 42.92\\\\ 7 & 0.0276 & 20.59 & $-$12.46 & 42.86\\\\ 8 & 0.113 & 20.37 & $-$12.86 & 42.32\\\\ 12 & 0.100 & 20.30 & $-$12.92 & 42.51\\\\ 14 & 0.0564 & 20.40 & $-$12.16 & 42.07\\\\ 17 & 0.0174 & 20.50 & $-$12.66 & 42.70\\\\  \\enddata  \\tablecomments{ Col. (1): Identification number assigned in this paper. Col. (2): {\\it ROSAT} count rate (counts s$^{-1}$). Col. (3): Galactic column $N_{\\rm H}$ (cm$^{-2}$), calculated following Dickey \\& Lockman (1990). Col. (4): X-ray flux in the 0.5--2 keV band (\\flux) assuming a power-law spectrum with $\\Gamma = 3$ and $N_{\\rm H}$ from Col. (3). Col. (5): X-ray luminosity in the 0.5--2 keV band (erg~s$^{-1}$). }  \\end{deluxetable}   Only six of the 19 sources (30\\%) were detected in the {\\it ROSAT} All-Sky Survey (Table 4), confirming the finding of Williams et~al. (2002) that not all NLS1s have detectable soft X-ray (0.5--2 keV) emission.  The data are not of sufficiently high quality to make spectral analysis possible, so to convert count rate to flux we assume a power-law spectrum with a fixed photon index of $\\Gamma$ = 3, which is roughly the average value observed in previous {\\it ROSAT}\\ studies of NLS1s (e.g., Boller et al. 1996; V\\'eron-Cetty et al. 2001; Williams et al. 2002).  We use the online program {\\tt webPIMMS}\\footnote{{\\tt http://heasarc.gsfc.nasa.gov/Tools/w3pimms.html}} to do the conversion, assuming that the only source of absorption is due to the Galaxy along the line of sight.  The derived soft X-ray luminosities range from $\\sim 10^{42.1}$ to $10^{42.9}$ \\lum\\ (Table 4), which probably can be attributed mostly to the AGN. The host galaxies of our sample are relatively faint (\\S 4.3), with absolute $g$-band magnitudes that roughly translate to $L_{B} \\approx 10^{9.4}-10^{10.1} \\, L_{\\sun}$ (Fukugita et al. 1996).  If these galaxies follow the $L_{B}-L_{X}$ relation of normal galaxies (e.g., Fabbiano 1989), they are expected to have $L_X \\approx 10^{38.5}$ erg~s$^{-1}$.  Only one of the objects was detected in the VLA FIRST survey (Becker, White, \\& Helfand 1995), although the entire sample overlapped with the FIRST survey area.  Object \\#10 has a 20~cm flux density of 1.3 mJy, which corresponds to a spectral power of $2.0 \\times 10^{22}$ W Hz$^{-1}$. This level of radio emission is not atypical of normal galaxies (Condon 1992), so we cannot conclude, without additional evidence, that it is associated with AGN activity.  Figure 6 places the sources on two commonly used line intensity ratio diagnostic diagrams. The dotted lines schematically denote the boundaries used by Ho et al. (1997a) to demarcate the main classes of emission-line objects. Although all the objects are clearly AGNs by the presence of their prominent broad lines, some of them in the [\\ion{O}{3}]/\\hbeta\\ vs.  \\ [\\ion{N}{2}]/\\halpha\\ and [\\ion{O}{3}]/\\hbeta\\ vs. [\\ion{S}{2}]/\\halpha\\ diagrams formally lie outside of the AGN zone ([\\ion{N}{2}]/\\halpha\\ $\\geq$ 0.6 and [\\ion{S}{2}]/\\halpha\\ $\\geq$ 0.4). This is rarely seen in local type 1 AGNs (Ho et al. 1997a).  Since the SDSS fibers subtend a significant fraction of each galaxy, the narrow emission-line ratios of the nuclei are, at least in part, contaminated by host galaxy emission, which would have the effect of lowering the [\\ion{S}{2}]/\\halpha\\ and [\\ion{N}{2}]/\\halpha\\ ratios.  Thus, one expects, a priori, that AGNs not selected by their narrow lines will, in some cases, exhibit narrow-line ratios that fall below the AGN threshold.  \\subsection{Nuclear Luminosities and Eddington Ratios}  Since our sample was selected to have low BH mass, we expect that the AGNs will generally be of relatively low luminosity.  Figure 7 shows the distribution of nuclear $g$-band absolute magnitudes for the sample.  To estimate this quantity, we first compute \\lf\\ (as described above), which is then transformed to $M_g$.  For those spectra with direct power-law fits, we use the best-fit spectral index, while for the remainder we assume an $f_\\lambda \\propto \\lambda^{-1}$ spectrum.  The AGN component has absolute magnitudes ranging from $M_g \\approx -15$ to $-18$ mag. Note that these nuclear luminosities, though generally not as low as many local Seyferts (Ho \\& Peng 2001), are still significantly lower than in many previously studied low-redshift type 1 AGNs (e.g., K\\\"{o}hler et al. 1997).  It is of interest to compare the nuclear luminosities with respect to the Eddington luminosities.  The selection criteria---low BH mass and a prominent AGN spectrum---are expected to favor objects with high Eddington ratios.  A major uncertainty in this exercise is that we do not have a robust estimate of the bolometric luminosity, but only a measurement of the optical luminosity (\\lf).  Since the spectral energy distributions of AGNs are known to vary significantly with accretion rate (Pounds et al. 1995; Ho 1999), it is doubtful that a single bolometric correction is always applicable. Nonetheless, assuming $L_{\\rm bol}$ = 9.8\\lf\\ (McLure \\& Dunlop 2004), the derived Eddington ratios mostly cluster around $L_{\\rm bol}/L_{\\rm Edd} \\approx 1$ (Fig. 8; Table 3), as expected.    \\begin{figure} \\begin{center} \\epsfig{file=f8_120.eps,width=2.4in,keepaspectratio=true,angle=-90} \\vskip -0.1mm \\figcaption[f8.eps]{ Distribution of Eddington ratios.  The bolometric luminosity was estimated from \\lf\\ (see text). \\label{fig8}} \\vskip -5mm \\end{center} \\end{figure}   \\subsection{Properties of the Host Galaxies}  The morphology of the host galaxies of intermediate-mass BHs is of fundamental importance for understanding the origin of this class of objects and their relationship to the overall demography of central BHs in galaxies.  The two currently known nuclear intermediate-mass BHs are found in anomalous hosts compared to the local AGN population (Ho et al. 1997b, 2003; Kauffmann et al. 2003), which is invariably affiliated with massive, bulge-containing galaxies.  Similarly, essentially every galaxy known to contain a supermassive BH has a bulge.  Yet, NGC 4395, which without a doubt lacks a bulge, evidently hosts a nuclear BH (Filippenko \\& Ho 2003).  The same could be said of POX 52 (Ho 2004b).  Although POX 52 appears to be a dwarf elliptical galaxy (Barth et al.  2004), this class of spheroids should be regarded as physically distinct from classical bulges (Bender, Burstein, \\& Faber 1992), and they may have originated as tidally stripped late-type disk galaxies (e.g., Moore et al. 1996).  Thus, both NGC 4395 and POX 52 provide compelling evidence that a classical bulge is {\\it not}\\ necessary for the formation of a central BH.  We have attempted to estimate the luminosities of the host galaxies by subtracting the contribution of the AGN from the total (Petrosian) $g$-band magnitude of the AGN plus host.  As illustrated in Figure 7 (see also Table 3), in most cases the luminosities of the AGNs are sufficiently low that they hardly perturb the distribution of total magnitudes; thus, the host galaxy magnitudes should be quite robust. Interestingly, the hosts are relatively low-luminosity and therefore more likely to be late-type galaxies: $M_g$(host) ranges from $-17.8$ to $-20.6$ mag, with an average value of $-19.2$.  According to Blanton et al.  (2003), the luminosity function of galaxies at $z = 0.1$ has $M_g^* = -19.39 + 5 \\log h$.  For $h = 0.72$, as assumed here, $M_g^* = -20.1$ mag.  Thus, the hosts of our sample are $\\sim 1$ mag fainter than $M_g^*$.  Our selection technique is unbiased with respect to galaxy morphology. Unfortunately, we have very limited information on the morphological types of the host galaxies.  The low resolution of the SDSS imaging makes morphological classification difficult for these objects.  By selection, the galaxy centers are dominated by an AGN core, which makes it impossible to say anything meaningful about the central structure of the galaxies.  If the low BH masses are accompanied by a low-mass (low-luminosity) galaxy or bulge, this makes it doubly hard to detect. We show a selection of the more interesting images in Figure 9.  Based on the SDSS imaging, one-third of the galaxies clearly have a disklike component.  Of the six that show disks, five have significant galaxy light in their spectra, which required PCA decomposition.  Given the low luminosities of these galaxies, perhaps they are analogs of NGC 4395.  The remainder are barely resolved enough to tell, but could be consistent with compact spheroids. We caution that the images are of low resolution and sensitivity, so in many cases we may be insensitive to an extended stellar component.  Higher resolution, deeper imaging, for example with the {\\it Hubble Space Telescope}, is essential to determine the true morphology and detailed structural parameters of these hosts.  In lieu of resolved morphological information, the SDSS collaboration has developed statistical proxies that describe morphological trends (Shimasaku et al. 2001; Strateva et al. 2001).  We consider the ``eclass'' (a PCA statistic), the \\sur\\ color, and the ``inverse concentration index.''  The eclass ranges from $-0.35$ for early-type galaxies to +0.5 for late-type galaxies. The mean eclass of our sample is $0.24$. This statistic is difficult to interpret because of the strong contamination by the AGN.  For $\\sim 150,000$ SDSS galaxies, Strateva et al. (2001) find a bimodal distribution in \\sur\\ color, with the division at \\sur\\ = 2.22 mag, which roughly divides early and late types.  Our galaxies have a mean color of $u-r \\approx 1.51$, consistent with early-type galaxies.  An inverse concentration index of $\\sim 0.3$ corresponds to a de~Vaucouleurs (1948) $r^{1/4}$ profile, while an exponential profile has an index of $\\sim 0.43$. The mean inverse concentration index of our objects is 0.41.   \\begin{figure} \\begin{center} \\epsfig{file=f9_50dpi.eps,angle=-90,width=3.1in,keepaspectratio=true} \\vskip -0.1mm \\figcaption[f9.eps]{ Sample SDSS $g$-band images of six objects. The vast majority of the sample are barely resolved and look similar to object \\#1.  In all images North is up and East is to the left. Each panel is $30^{\\prime\\prime} \\times 30^{\\prime\\prime}$. \\label{fig9}} \\vskip -5mm \\end{center} \\end{figure}   \\subsection{Implications and Future Work}  While the discovery of the \\msigma\\ relation established a compelling link between the evolution of galaxies and central BHs, models of this process are faced with significant uncertainties.  Possible models include seed BHs formed from Population III stars, which then gain mass through mergers (Schneider et al. 2002; Volonteri, Haardt, \\& Madau 2003) or through accretion (Adams, Graff, \\& Richstone 2001; Islam, Taylor, \\& Silk 2003).  Alternatively, seed BHs may form directly from a collapsing gas cloud (Koushiappas, Bullock, \\& Dekel 2004; for a more complete discussion, see the review by van~der~Marel 2004).  However, there are few direct observational constraints on the nature of seed BHs, the path by which nuclear BHs grow, or the form of the BH mass function.  The low end of the BH mass function, in particular, provides critical input to check theoretical models of quasar formation (e.g., Haehnelt, Natarajan, \\& Rees 1998; Bromley, Somerville, \\& Fabian 2004), and it affects predictions of future gravitational wave experiments like {\\it LISA}\\ (Hughes 2002).  A useful approach, first proposed by So{\\l}tan (1982), matches the local BH mass density inferred from velocity dispersion measurements with an accretion history inferred from observations of luminous quasars. Yu \\& Tremaine (2002) conclude that, for standard assumptions about radiative efficiency and quasar lifetimes, accretion during the luminous quasar phase can account for the majority of the local BH mass density.  Aside from the possibility of significant obscured accretion (Fabian 2004), this method is limited by its use of the SDSS velocity dispersion function, which at the moment is only well determined for early-type galaxies with $\\sigma_{\\star} \\gtrsim 200$ \\kms\\ (Sheth et al. 2003), corresponding to $M_{\\rm BH} \\gtrsim 10^8$ \\msun.  Much of the current observational and theoretical effort on intermediate-mass BHs has focused on studies of star clusters (Gebhardt et al. 2002; Gerssen et al. 2002) and so-called ultraluminous X-ray sources (see van~der~Marel 2004).  The existence of BHs in either class of objects is still far from proven, and, in any event, their relationship to galactic nuclei is unclear.  With the recognition that galaxies like NGC 4395 and POX 52 do harbor central intermediate-mass BHs, the most straightforward way to make progress would be to find many more objects like them.  This study represents the first step in this direction.  By sifting through the massive SDSS database, we have identified a sample of AGNs that we believe are promising candidates for hosting BHs with masses below the $10^6$ \\msun\\ threshold.  Our sample opens up several new avenues for further study.  While well-defined, our sample is likely to have missed intermediate-mass BH candidates such as NGC 4395 or POX 52, whose spectrum is dominated by narrow emission lines, with the broad-line component visible only under very high signal-to-noise ratio.  Moreover, if either AGN were embedded in a much more luminous host galaxy, the weak broad lines would be even harder to discern.  Furthermore, we know that pure narrow-line (type 2) AGNs do exist, either because the BLR is obscured or intrinsically absent, and there is no reason to suppose that this situation does not persist to the regime of intermediate-mass BHs.  A complimentary approach, then, is to use the width of the narrow emission lines to select potential candidates.  While the detailed geometry and kinematics of the narrow-line region (NLR) in AGNs are complex and not well understood, it has been argued that its velocity field is largely dominated by the gravitational potential of the bulge of the host galaxy (Smith, Heckman, \\& Illingworth 1990; Whittle 1992a, b).  If so, then the velocities in the NLR are roughly virial, and on average the velocity dispersion of the line-emitting gas, $\\sigma_{\\rm g}$, should be approximately equal to the velocity dispersion of the bulge stars, \\sigmastar.  The most detailed comparison between NLR and stellar kinematics for a large sample of Seyfert galaxies has been done by Nelson \\& Whittle (1996), who indeed find that in most cases $\\sigma_{\\rm g}\\, \\approx$ \\sigmastar, where $\\sigma_{\\rm g}$ was measured using [\\ion{O}{3}] $\\lambda5007$.  The exception are objects with powerful, jetlike radio structures where $\\sigma_{\\rm g}$ can be much larger than \\sigmastar, presumably because of extra nongravitational motions induced by the radio outflows.  Thus, in general one expects that $\\sigma_{\\rm g}$, as measured in [\\ion{O}{3}] $\\lambda5007$ or other optical forbidden lines, can be taken as a rough proxy, or perhaps more realistically as an upper limit, for \\sigmastar.  We are in the process of enlarging the current sample through a selection based on this strategy.  Another important next step is to obtain stellar velocity dispersions for these objects.  Such measurements will test whether the BH mass estimates we derive obey the \\msigma\\ relation.  Conversely, if we assume that the \\msigma\\ relation is universally valid, the stellar velocity dispersion measurements will test the reliability of the virial mass estimator for AGNs. It would also be valuable to compare the velocity dispersions derived from stars with those derived from gas.  Another area requiring attention is to obtain high-resolution, deep imaging.  Much better data are needed to quantify the morphologies and detailed structural parameters of the host galaxies.  Do these objects have bulges, and if so, what kind and where do they fall on galaxy scaling relations like the fundamental plane?  Finally, much work remains to be done, especially at non-optical wavelengths, to further characterize the properties of the AGNs themselves."
    },
    "0404/astro-ph0404606_arXiv.txt": {
        "abstract": "We address the question of why low-luminosity radio sources with similar flat radio spectra show a range of optical activity. The investigation is based on the spectral energy distributions (SEDs) of objects from the 200 mJy sample. We gathered new data from the VLA at 43 GHz, from SCUBA in the JCMT at 2000, 1350 and 850 $\\mu$m and from the ISOPHOT instrument on {\\it ISO} at 170, 90, 60 and 25 $\\mu$m. There is considerable diversity amongst the SEDs of the objects: there are objects with steep broadband spectra  between centimetre and millimetre bands (14\\% of the sample), there are those with flat broadband spectra over most of the spectral range (48\\% of the sample) and there are those which show pronounced sub-mm/infrared excesses (27\\% of the sample). Some objects of the first group have two sided radio morphology indicating that their pc-scale emission is not dominated by beamed jet emission. Amongst the objects that have smooth broadband spectra from the radio to the infrared there are passive elliptical galaxies as well as the expected BL Lacs objects. The most pronounced sub-mm/infrared excesses are shown by the broad emission line objects. ",
        "introduction": "Active galaxies are distinguished by the broad range of the electromagnetic spectrum over which they radiate and by the fact that most of the luminosity comes from the compact nuclear region of the galaxy. Considering the emission in the radio domain, AGNs come, broadly, in two flavours: those with weak radio emission, the ``radio-quiet'', and those with strong radio emission, the ``radio-loud''. The radio-loud objects are generally divided further into core-dominated objects and lobe-dominated objects. The former generally have flat radio spectra and it is upon this sub-set of objects we concentrate in this paper. At optical wavelengths, amongst the high-luminosity flat-spectrum AGN, we find objects with strong optical emission lines, the quasars. At lower luminosities we find objects with very weak or absent emission lines and polarised and variable continuum, the BL Lacs, Seyfert-like objects with strong emission lines, and objects with no sign of optical activity, the passive elliptical galaxies. Our primary objective is to investigate the origin of the panoply of optical properties displayed by objects with very similar radio properties, in particular the flat-spectrum, core-dominated, low-luminosity radio sources.  The investigation is based on the 200 mJy sample (\\citealt{marchatese}, \\citealt{mar95}, hereinafter M96). The sample has been slightly revised as a result of better available data (\\citealt{antontese}) and the revised sample is presented here. All the sources are core-dominated when observed with the VLA with 0.2~arcsec resolution at 8.4 GHz and have flat radio spectra between 1.4 and 5 GHz. Their ``radio homogeneity'', however, disappears at the optical wavelengths: about 30\\%\\ of the objects show optical properties of BL Lacs, but in the remaining $\\sim$~70\\%\\ of the objects the optical properties split into broad and narrow emission line objects and apparently inactive elliptical galaxies (M96). The range of optical activity amongst the 200 mJy objects begs the question: why in some objects does the non-thermal component, clearly present in the radio, disappear at optical frequencies?  The origin of the split in the optical properties must either arise from differences in the AGN emission and/or differences in the host galaxy properties. In order to try to identify that origin we have embarked in a multi-wavelength study of the 200 mJy sample. This work includes the study of the large-scale properties such as the structure of the host galaxies and the small scale properties, in particular the analysis of the spectral energy distributions of the AGNs. In March\\~a et al. (in prep) we present the surface photometry of approximately half of the sample, and we find that the large-scale properties of the sources are very homogeneous and also similar to those found in radio galaxies in general. In this paper we present the SEDs of the 200 mJy sample, and we pose the following questions:\\\\ \\noindent $\\bullet$ Is the lack of optical activity in some objects due to extinction?\\\\ \\noindent $\\bullet$ Do all low luminosity flat spectrum radio sources have beamed BL Lac emission?  \\\\ Extinction could hide core optical non-thermal emission. Radio core-dominated objects are expected to be viewed down the radio axis and thus, according to unification models, should be least likely to be affected by extinction. These models, however, only take into account central obscuring regions, but more extended dusty regions may also be present and prevent us from detecting the AGN emission. The far-infrared band offers the best chance to trace dust emission, and for this reason new infrared data were gathered with the {\\it Infrared Space Observatory} ({\\it ISO}).\\\\ In what concerns the second question, the fact that an object has a flat radio-spectrum does not necessarily prove that the emission originates in relativistic jets and is beamed. If the emission comes from self-absorbed lobes, as for example in Compact Symmetric Objects (CSOs), then the flat radio spectrum synchrotron emission is not expected  to extend to higher frequency bands. To investigate this, new sub-mm data were gathered with SCUBA. The main astronomical objective of these observations is to try and look how far the synchrotron component extends into the mm/sub-millimetre part of the spectrum. We use VLBI observations (Bondi et al, in prep.) to distinguish between beamed core-jet sources from intrinsically compact and young CSOs.\\\\ The framework in which we will discuss the SEDs of the objects is the following:\\\\ \\noindent 1. The centimetre wavelength properties of the 200 mJy objects (flat radio spectra plus compact radio structure) are similar to other objects that have dominant synchrotron cores at centimetre wavelengths. Thus we assume that their core radio emission is synchrotron in origin.\\\\ \\noindent 2. This synchrotron emission may extend into the sub-mm, infrared and even the optical part of the spectrum. In this latter case one would call the object a BL Lac. \\\\ \\noindent 3. Cold dust may contribute to the sub-mm emission. \\\\ \\noindent 4. If the optical synchrotron photons are not detected because they are absorbed by dust, an excess of infrared emission (due to reprocessed energy) showing above the extrapolation of the non-thermal broadband spectra might be present.\\\\ \\noindent This paper is organised as follows. In Section ~\\ref{section-obs} the properties of the 200 mJy sample are summarised, the new observations and data reduction are presented. In Section ~\\ref{section-results} the shape of the SEDs are discussed. The important points  of this work are summarised in Section ~\\ref{section-conclusions}. Throughout the paper we assume H$_o$=50 kms$^{-1}$Mpc$^{-1}$ and $q_o$=0.   \\begin{table} \\caption{200 mJy selection criteria} \\label{200mjy} \\begin{tabular}{@{}ccccc} \\hline S$_{\\mbox{\\tiny 5GHz}}$ & $\\alpha_{1.4-5 GHz}$ & R  & $\\delta_{\\tiny 1950}$ & $|b|$\\\\ $\\geq$ 200 mJy & $\\geq -$ 0.5 & $\\leq$ 17 mag & $\\geq$ 20$^o$ &$\\geq$ 12$^o$  \\\\ \\hline \\end{tabular} \\medskip   The spectral index $\\alpha$ ($S\\propto\\nu^{\\alpha}$) is defined between 1.4 GHz (NVSS catalogue) and 5 GHz (GB6 catalogue), R magnitudes are from APM, corrected by the calibration presented in \\citet{antontese}. \\end{table} ",
        "conclusions": "\\label{section-conclusions}  The main goal of this work was to investigate the origin of the diverse optical activity amongst similar radio-loud low-luminosity flat-spectrum objects. The investigation was based on the 200 mJy sample, a radio-selected sample showing a range of optical activity. To achieve that goal, multi-wavelength observations were obtained (millimetre data with VLA, sub-mm data with SCUBA and infrared data with {\\it ISO}). Adding  data from the literature allowed us to produce the SEDs of 64 objects. The SEDs were analysed with two main goals: to characterize the non-thermal synchrotron component and to search for dust emission. Our findings may be summarised as follows: \\begin{enumerate} \\item There is considerable diversity amongst the SEDs but objects fall broadly into three groups: there are objects with steep spectra between 1.4 GHz and 43 GHz (14\\% of the sample), there are those with flat spectra over most of the spectral range (48\\% of the sample) and there are those which show pronounced sub-mm/infrared excesses (27\\% of the sample). \\item Sub-mm observations indicate that a major fraction of the 200 mJy objects have smooth and flattish synchrotron spectra up to the sub-mm band. This includes BL Lacs, BL Lac Candidates, and some Seyfert1-like and PEG objects. In only few cases the dominant sub-mm emission mechanism is clearly thermal dust emission  (for example, Arp 220 and Mrk 231). \\item The broadband spectra between radio and infrared wavelengths suggest that different objects have different dominant emission mechanisms in the infrared band. Some objects show a smooth broadband spectrum  which can be explained by synchrotron emission -- all BL Lac and some broad-emission line objects are in this group.  Others show moderate infrared excess over the extrapolated radio/submillimetre emission, suggesting that both synchrotron emission and dust emission contribute to the infrared -- this group includes broad and narrow emission line objects and PEGs. The far-infrared bump present in the majority of the PEGs indicates the presence of cold dust, which mass we estimate to be similar to that found in typical early-type galaxies. Even though there is dust emission the analysis of Serote Roos \\& Gon\\c calves (2004) suggests that it is not strongly affecting the line emission regions. Thus extinction is unlikely to account for the low level of observed activity on PEGs. \\item Correlating the broadband SED shapes with the radio morphology (through high resolution radio images available in the literature) allow us to conclude that the sources with two sided morphologies nearly all have falling spectra between 1.4 GHz and 43 GHz. Not surprisingly, all the sources showing core-jet radio morphology have flat broadband spectra often extending well into the sub-mm/infrared range. The vast majority of these are spectroscopically classified as BL Lacs or Seyfert-like objects. In a subsequent paper we will discuss in more detail the SEDs of those objects that are synchrotron emitters, independently of the optical classification. We will compare their spectra with those of blazars, in general, and suggest that there may be many more objects  with low frequency peaks in their synchrotron spectra than has been previously thought.  \\end{enumerate}"
    },
    "0404/astro-ph0404026_arXiv.txt": {
        "abstract": "{ We present the kinematics of the molecular gas in the barred spiral galaxy M\\,83 (NGC5236). The study is based on $^{12}$CO(\\mbox{$J$=1--0} and \\mbox{2--1}) observations with the Swedish-ESO Submillimetre Telescope (SEST). Iso-velocity maps of the entire optical disk, 10$\\arcmin$\\,$\\times$\\,10$\\arcmin$ or 13\\,$\\times$\\,13\\,kpc, are produced. They show the pattern of an inclined, rotating disk, but also the effects of streaming motions along the spiral arms. A dynamical mass of about 6\\,$\\times$\\,10$^{10}$\\,M$_\\odot$ is estimated by fitting the rotation curve of an exponential disk model to these data. The gas constitutes about 13\\% of the disk mass. The pattern speed is determined from the residual velocity pattern The locations of various resonances are discussed. The molecular gas velocity dispersion is determined, and a trend of decreasing dispersion with increasing galactocentric radius is found. A total gas (H$_2$+\\ion{H}{i}+He) mass surface density map is presented, and compared to the critical density for star formation of an isothermal gaseous disk. The critical density is exceeded in the spiral arms, but not in the interarm regions. The locations of Giant Molecular Associations (GMAs) and \\ion{H}{ii} regions are consistent with this scenario of dynamically induced star formation. ",
        "introduction": "\\input{H0230_S1.tex} ",
        "conclusions": ""
    },
    "0404/astro-ph0404483_arXiv.txt": {
        "abstract": "s{ The observed velocities of radio pulsars, which range in the hundreds kilometers per second, and many of which exceed 1000~km/s, are not explained by the standard physics of the supernova explosion.  However, if a sterile neutrino with mass in the 1--20~keV range exists, it would be emitted asymmetrically from a cooling neutron star, which could give it a sufficient recoil to explain the pulsar motions.  The same particle can be the cosmological dark mater.  Future observations of X-ray telescopes and gravitational wave detectors can confirm or rule out this explanation.  } ",
        "introduction": "Pulsar velocities range from 100 to 1600~km/s\\cite{astro,astro_1}. Their distribution leans toward the high-velocity end, with about 15\\% of all pulsars having speeds over (1000~km/s)\\cite{astro_1}. The origin of these velocities is a long-standing puzzle. Pulsars are born in supernova explosions, so it would be natural to look for an explanation in the dynamics of the supernova\\cite{explosion}. However, state-of-the-art 3-dimensional numerical calculations\\cite{fryer} show that even the most extreme asymmetric explosions do not produce pulsar velocities greater than 200~km/s.  Earlier 2-dimensional calculations\\cite{sn_2D} claimed a maximal pulsar velocity up to 500~km/s to be possible.  Of course, even this size of the kick was way too small to explain the large population of pulsars with speeds above 1000~km/s.  Given the absence of a ``standard'' explanation, one is compelled to consider alternatives, possibly involving new physics.  One of the reasons why the standard explanation fails is because most of the energy is carried away by neutrinos, which escape isotropically.  The remaining energy must be distributed with a substantial asymmetry to account for the large pulsar velocities.  In contrast, only a few per cent anisotropy in the distribution of neutrinos would give the pulsar a kick of required magnitude.  Neutrinos are {\\em produced } anisotropically, but they {\\em escape} isotropically.  The asymmetry in production comes from the asymmetry in the basic weak interactions in the presence of a strong magnetic field. Indeed, if the electrons and other fermions are polarized by the magnetic field, the cross section of the urca processes, such as $n+e^+ \\rightleftharpoons p+ \\bar \\nu_e$ and $p+e^-\\rightleftharpoons n+ \\nu_e $, depends on the orientation of the neutrino momentum.  Depending on the fraction of the electrons in the lowest Landau level, this asymmetry can be as large as 30\\%, much more than one needs to explain the pulsar kicks\\cite{drt}.  However, this asymmetry is completely washed out by scattering of neutrinos on their way out of the star\\cite{eq}.  If, however, the same interactions produced a particle which had even weaker interactions with nuclear matter than neutrinos, such a particle could escape the star with an asymmetry equal its production asymmetry.  It is intriguing that the same particle can the dark mater.  The simplest realization of this scenario is a model that adds only one singlet fermion to the Standard Model.  The SU(2)$\\times$U(1) singlet, a sterile neutrino, mixes with the usual neutrinos, for example, with the electron neutrino.  For a sufficiently small mixing angle between $\\nu_e$ and $\\nu_s$, only one of the two mass eigenstates, $\\nu_1$, is trapped.  The orthogonal state, \\begin{equation} | \\nu_2 \\rangle = \\cos \\theta_m | \\nu_s \\rangle + \\sin \\theta_m | \\nu_e \\rangle , \\end{equation} escapes from the star freely.  This state is produced in the same basic urca reactions ($\\nu_e+n\\rightleftharpoons p+e^-$ and $\\bar\\nu_e+p\\rightleftharpoons n+e^+$) with the effective Lagrangian coupling equal the weak coupling times $\\sin \\theta_m$.  We will consider two ranges of parameters, for which the $\\nu_e \\rightarrow \\nu_s$ oscillations occur on or off resonance. First, let us suppose that a resonant oscillation occurs somewhere in the core of the neutron star. Then the asymmetry in the neutrino emission comes from shift in the resonance point depending on the magnetic field\\cite{ks97}.  Second, we will consider the off-resonance case, in which the asymmetry comes directly from the weak processes, as described above\\cite{fkmp}. ",
        "conclusions": ""
    },
    "0404/astro-ph0404160_arXiv.txt": {
        "abstract": "Recent {\\it Chandra} observations of clusters of galaxies revealed the existence of a sharp ridge in the X-ray surface brightness where the temperature drops across the front. This front is called the cold front. We present the results of two-dimensional magnetohydrodynamic simulations of the time evolution of a dense subcluster plasma moving in a cluster of galaxies. Anisotropic heat conduction along the magnetic field lines is included. In the models without magnetic fields, the numerical results indicate that the heat conduction from the hot ambient plasma heats the cold dense plasma of the subcluster and diffuses out the cold front. When magnetic fields exist in a cluster of galaxies, however, cold fronts can be maintained because the heat conduction across the magnetic field lines is suppressed. We found that, even when the magnetic fields in a cluster of galaxies are disordered, heat conduction across the front is restricted because the magnetic field lines are stretched along the front. Numerical results reproduced the X-ray intensity distribution observed in the A3667 cluster of galaxies. ",
        "introduction": "High-resolution {\\it Chandra} observations have revealed various substructures in the hot X-ray emitting plasmas in clusters of galaxies. For example, sharp discontinuities of X-ray surface brightness distributions have been found around subclusters moving in a cluster of galaxies such as A2142 \\citep{mar00} and A3667 (Vikhlinin, Markevitch, \\& Murray 2001b). The temperature drops sharply across this boundary between the gravitationally confined, cold dense plasmas in the moving subcluster and the hot, less-dense ambient plasma. This sharp discontinuity is called the cold front. It is not a shock front, because {\\it Chandra} images of the X-ray surface brightness show that the temperature decreases in the denser side.  \\citet{ett00} suggested that the sharp temperature gradient observed by {\\it Chandra} in A2142 requires the heat conductivity near the front to be reduced by several orders of magnitude from the classical Spitzer value \\citep{spi62}, $\\kappa_{\\rm Sp}= 5 \\times 10^{-7} T^{5/2}$ $\\, {\\rm erg \\, s^{-1} \\, cm^{-1} \\, K^{-1}}$.  Based on the observations of A3667, Vikhlinin, Markevitch, \\& Murray (2001a) suggested that turbulent, intergalactic magnetic fields stretched along the cold front reduce the growth rate of the Kelvin-Helmholtz (K-H) instability, which possibly destructs the front. Furthermore, heat conduction across the magnetic fields can be suppressed. It has been recognized that heat conduction in a cluster of galaxies can be very efficient (e.g., \\citealt{tak77}) because the Coulomb mean free path, $l_{\\rm c} \\sim 4 \\, (kT / 5 \\, {\\rm keV})^{2}(n / 10^{-3}{\\rm cm^{-3}})^{-1} \\, {\\rm kpc}$, is large. Here, $n$ is the number density. The time required for heat to diffuse by conduction across a length $L$ is given by $\\tau_{\\rm Sp} \\sim \\rho L^{2}/ \\kappa_{\\rm Sp} \\sim 10^{7} (kT/ 5 \\, {\\rm keV})^{-5/2} (n / 10^{-3} \\, {\\rm cm^{-3}})(L/100 \\, {\\rm kpc})^{2} \\, {\\rm yr}$, where $\\rho$ is the density. However, if magnetic fields exist, the characteristic length of the heat exchange across magnetic field lines is reduced to the Larmor radius, $r_{\\rm L} = 2500 \\, (B/1 \\, {\\rm \\mu G})^{-1}$ $(T/5 \\, {\\rm keV})^{1/2} \\, {\\rm km}$. Since the typical magnetic field strength in a cluster of galaxies is $\\sim {\\rm \\mu G}$ (e.g., \\citealt{kro94,car02}), the heat conductivity across magnetic field lines is drastically reduced.  In this Letter, we present the results of numerical simulations of the formation of cold fronts. These simulations include the effects of heat conduction. We adapted the simulation code, originally developed for solar flare simulations \\citep{yok97,yok01}, which includes anisotropic heat conduction along magnetic field lines. ",
        "conclusions": "\\citet{hei03} carried out hydrodynamical simulations of the interaction of cold subcluster plasma and hot ambient matter. They suggested that a cold front can be formed after the passage of shock waves through the subcluster plasma, when the cold plasma, initially pushed forward by the shock, falls back. However, in their simulations heat conduction and magnetic fields were neglected. In this Letter, we presented the results of numerical simulations which included both magnetic fields and heat conduction. We confirmed that, since the Spitzer conductivity is very large in hot ambient plasma, cold fronts cannot be maintained longer than $0.3 \\, {\\rm Gyr}$, unless the heat conduction across the front is suppressed by the magnetic fields.  We also carried out numerical simulations which included heat conduction starting from the same initial condition as \\citet{hei03}. We found that, when we adopt the Spitzer conductivity, the subcluster evaporates rapidly and the cold front is not formed. When we included magnetic fields and the effects of anisotropic heat conduction, the numerical results resembled Heinz's results. As suggested by \\citet{hei03}, after the passage of shock waves, the cold plasma pulled back by gravity helps sharpen the cold front.  In \\S \\ref{sec:m}, we presented numerical results that showed the K-H instability does not grow near the cold front, because the velocity of inflowing plasmas is reduced after the plasmas pass through the bow shock. In the magnetohydrodynamical models, the magnetic fields, which stretched along the contact surface, also stabilized the K-H instability. Thus, the K-H instability does not threaten the existence of cold fronts.  In \\S \\ref{sec:mtc} we showed results of a model with initially disordered magnetic fields. Numerical results show that the magnetic fields near the front are gradually stretched. Thus, the initially tangled magnetic field lines become parallel to the front and suppress the heat conduction across the front. We confirmed the suggestion by \\citet{vik01a} that the stretching of turbulent intergalactic magnetic fields enables the maintenance of cold fronts. If the scale of magnetic turbulence is smaller than the mean free path, heat can be conducted almost isotropically. The heat conductivity in such a turbulent MHD medium is as large as $1/5$ of the Spitzer value \\citep{nar01} and still large enough to diffuse out the cold front. We expect that coherent fluid motions and the stretching of the magnetic fields along the cold front increase the coherent length of field lines and make that region less turbulent.  Let us now discuss the three-dimensional (3D) effects. Since magnetic fields compressed in front of the cold front can expand in the third direction, the strength of the magnetic fields in this region will be reduced in 3D. However, even such weakened magnetic fields suppress the heat conduction across them. Thus, we expect that cold fronts can be sustained even in 3D. We hope to report the results of such 3D MHD simulations in subsequent papers.  In this Letter, we assumed that the ambient plasma has magnetic fields. Recently, \\citet{oka03} pointed out the possibility that magnetic fields are spontaneously generated by the Ramani-Laval instability around the cold front. Such magnetic fields can also help to suppress the heat conduction across the front."
    },
    "0404/astro-ph0404497_arXiv.txt": {
        "abstract": "It is argued that the observed vacuum energy density and the small values of the neutrino masses could be due to anthropic selection effects. Until now, these two quantities have been treated separately from each other and, in particular,  anthropic predictions for the vacuum energy were made under the assumption of zero neutrino masses. Here we consider two cases. In the first, we calculate predictions for the vacuum energy for a fixed (generally non-zero) value of the neutrino mass. In the second we allow both quantities to vary from one part of the universe to another. We find that the anthropic predictions for the vacuum energy density are in a better agreement with observations when one allows for non-zero neutrino masses. We also find that the individual distributions for the vacuum energy and the neutrino masses are reasonably robust and do not change drastically when one adds the other variable. ",
        "introduction": "The smallness of the vacuum energy density $\\rL$ (the cosmological constant) is one of the most perplexing mysteries that we now face in theoretical physics.  The values of $\\rL$ indicated by particle physics models are far too large, at least by 60 orders of magnitude. The only plausible explanation that has so far been suggested is that $\\rL$ is a stochastic variable, taking different values in different parts of the Universe, and that the smallness of the observed value is due to anthropic selection effects. The proposed selection mechanism is very simple\\footnote{For earlier suggestions that the cosmological constant may be anthropically selected, see \\cite{Davies,BT,LindeLambda}.} \\cite{Weinberg87,Vilenkin95a,Efstathiou95,Vilenkin95d,Weinberg97,MSW,GV,GLV03}. The growth of density fluctuations leading to galaxy formation effectively stops when the vacuum energy dominates the Universe.  In regions where $\\rL$ is greater, it will dominate earlier, and thus there will be fewer galaxies (and therefore fewer observers).  Suppose ${\\tilde \\rL}$ is the value for which the vacuum energy dominates at about the epoch of galaxy formation, $z_\\L\\sim z_G$.  Then, values $\\rL\\gg{\\tilde\\rL}$ will be rarely observed, because the density of galaxies in the corresponding regions is very low.  On the other hand, values $\\rL\\ll{\\tilde\\rL}$ are also rather unlikely, simply because this range of values is very small. (This argument assumes that the prior distribution for $\\rL$ is not sharply peaked at $\\rL=0$; see below.)  Thus, a typical observer should expect to measure $\\rL\\sim{\\tilde\\rL}$: the vacuum energy should dominate roughly at the epoch of galaxy formation.  This anthropic prediction was subsequently confirmed by supernovae and CMB observations.  The observed value \\cite{supernovae,Spergel03} for the vacuum energy density, in units of the critical density, $\\Omega_\\L^{obs}\\approx 0.7$, is in a qualitative agreement with the anthropic probability distribution, which is peaked at \\cite{Vilenkin95a,Efstathiou95,MSW,GV} $\\Omega_\\L \\sim 0.9$.  The most recent analysis \\cite{GLV03} finds that anthropic predictions are in agreement with observations at a $2\\sigma$ level.  Here we show that the agreement can be further improved by allowing for a nonzero neutrino mass. The value $\\tilde \\rho_\\L$ at the peak of the probability distribution for $\\rho_\\L$ depends on the magnitude of density fluctuations on the galactic scale, $\\sigma_G$. The calculation in \\cite{GLV03} was performed assuming massless neutrinos and a scale-invariant spectrum of fluctuations, normalized using the large-scale CMB measurements. If $\\sigma_G$ is decreased, galaxies form later, and $\\tilde \\rho_\\L$ becomes smaller, thus improving the agreement with $\\rho_\\L^{obs}$.  This is precisely the effect of massive neutrinos, since they slow down the growth of density fluctuations.  If the sum of the three neutrino masses is $m_\\nu \\sim 1-2$ eV, this suppression is sufficiently large to bring the anthropic prediction for $\\rho_\\L$ into agreement with observations at the $1\\sigma$ level.  It is very intriguing that the value $m_\\nu\\sim 1$ eV for the sum of the neutrino masses was itself recently predicted from anthropic considerations.  It was suggested in \\cite{TVP03} that the neutrino masses might also be stochastic variables, changing from one part of the Universe to another, and that their small values might be due to anthropic selection.  In regions with a larger value of $m_\\nu$, the suppression of density fluctuations is stronger.  As a result, there are fewer galaxies in such regions, and the corresponding values of $m_\\nu$ are less likely to be observed.  The probability distribution for $m_\\nu$ calculated in \\cite{TVP03} is peaked at $m_\\nu\\approx 2$ eV, with a $1\\sigma$ range between 0.5 eV and 3 eV. This corresponds to a mild but non-negligible suppression of galaxy formation.\\footnote{It is conceivable that the suppression of galactic-scale density fluctuations relative to the largest observable scales is not due to massive neutrinos, but reflects some feature in the primordial fluctuation spectrum. Such suppressed small-scale fluctuations have been suggested by certain galaxy cluster observations (assuming that the matter density parameter is $\\Omega_M\\sim 0.3$, see \\cite{Neta}), and by the WMAP team \\cite{Spergel03} who pointed out that an initial spectrum with a decreasing running spectral index is somewhat preferred by the data. Whatever the origin of such a suppression, it improves the agreement of the anthropic predictions for $\\rL$ with observations.} Indeed, recent analysis of X-ray cluster data combined with CMB and 2dF has suggested that $m_\\nu=0.64^{+0.40}_{-0.28}$ eV \\cite{Allen03}, although this remains controversial given cluster-related systematic uncertainties \\cite{MaxSDSS03}. In this paper, we shall also consider models where both $m_\\nu$ and $\\rL$ are assumed to be stochastic variables.  The paper is organized as follows. In Section \\ref{sec:general} we describe our general approach and justify our choice of the theoretical prior distributions for $m_\\nu$ and $\\rho_\\L$. In Section \\ref{sec:pmurho} we calculate the probability distribution for $\\rL$ at nonzero values of $m_\\nu$, as well as the joint probability distribution for $m_\\nu$ and $\\rho_\\L$ in models where both are variable. We conclude with a discussion of our results in Section \\ref{sec:summary}. Some technical details are given in the Appendix. ",
        "conclusions": "\\label{sec:summary}  In conclusion, we have argued that both the small vacuum energy density and small values of the neutrino masses may be due to anthropic selection. We have considered a model where $m_\\nu$ and $\\rho_\\L$ vary independently, with flat priors for both variables. The predictions of this model are qualitatively similar to those derived by considering sub-ensembles with only one of these parameters variable and the other fixed. But there is an interesting quantitative difference. The probability distribution for the joint ensemble is shifted towards smaller values of both $m_\\nu$ and $\\rho_\\L$. As a consequence, the agreement between the predicted and observed values of $\\rho_\\L$ is improved.  In this work we have only considered the case of positive $\\rho_\\L$. Negative values of $\\rho_\\L$ were discussed in \\cite{GV} and in \\cite{KL03}. There, it was found that the probability of $\\rho_\\L < 0$ is comparable to that of $\\rho_\\L > 0$.  In our analysis we have assumed that the number of observers does not sensitively depend of $m_\\nu$ or $\\rho_\\L$ and is simply proportional to the mass of the galaxy, Eq.~(\\ref{Nobs}). This appears to be a reasonable assumption for small values of $m_\\nu$ and $\\rho_\\L$, such that they have little effect on structure formation. For larger values, the properties of the galaxies, including the number of habitable stars per unit mass, may be affected.  For example, if $\\rho_\\L$ gets large, structure formation ends early and galaxies have a higher density of matter. This may increase the danger of nearby supernova explosions and the rate of near encounters with stars, large molecular clouds, or dark matter clumps. Gravitational perturbations of planetary systems in such encounters could send a rain of comets from the Oort-type cloud towards the inner planets, causing mass extinctions. Further increase of density may lead to frequent disruptions of planetary orbits \\cite{Q}. The effect of all this is to suppress the probability density at large values of $\\rho_\\L$ and thus to shift the distribution towards smaller values. This may further improve the agreement of the anthropic predictions with the data. With better understanding of galaxy and star formation, we may be able to obtain a quantitative measure of these effects in not so distant future.   \\bigskip  We thank Neta Bahcall, Gia Dvali and Jaume Garriga for useful comments and discussions. This work was supported in part by the National Science Foundation. AV was supported in part by the National Science Foundation and the John Templeton Foundation. MT was supported by NSF grant AST-0134999, NASA grant NAG5-11099, Research Corporation and the David and Lucile Packard Foundation.  \\appendix"
    },
    "0404/astro-ph0404518_arXiv.txt": {
        "abstract": "I describe work on two aspects of magnetic field evolution relevant for the ``recycling'' scenario for making millisecond radio pulsars. First, many of the theoretical ideas for bringing about accretion-induced field decay rely on dissipation of currents in the neutron star crust. I discuss field evolution in the crust due to the Hall effect, and outline when it dominates Ohmic decay. This emphasises the importance of understanding the impurity level in the crust. Second, I briefly discuss the progress that has been made in understanding the magnetic fields of neutron stars currently accreting matter in low mass X-ray binaries. In particular, thermonuclear X-ray bursts offer a promising probe of the magnetic field of these neutron stars. ",
        "introduction": "While spin up is naturally expected from accretion of angular momentum, magnetic field decay in an accreting neutron star is much less well-understood (see Bhattacharya \\& Srinivasan 1995 for a review). The timescales for evolution of currents in the core of a neutron star is extremely long (Baym, Pethick, \\& Pines 1969; Goldreich \\& Reisenegger 1992, hereafter GR92). However, short decay times ($\\lesssim 10^7$ years) are possible in the neutron star crust (as emphasised very early on, e.g.~Ewart et al.~1975), and this underlies one class of models for accretion-induced field decay.  One idea is to place the currents supporting the stellar magnetic field entirely in the crust. This might be the case if the field is generated from thermomagnetic effects, for example (Blandford, Applegate, \\& Hernquist 1983). The evolution due to Ohmic decay is then straightforward to calculate, and has been followed in many papers (Urpin, Geppert, \\& Konenkov 1997, Konar \\& Bhattacharya 1997, and references therein). The role of accretion is to heat the crust, reducing the electrical conductivity. The amount by which the field decays depends on the accretion lifetime, depth of the currents, and how much current is advected and subsequently ``frozen'' into the superconducting core. Thermomagnetic processes may also destroy field in the reheated crust (Blondin \\& Freese 1986).  A second proposal is that as a radio pulsar spins down, outward moving vortices in the superfluid and superconducting core push magnetic fluxoids into the crust, where magnetic energy dissipates (Srinivasan et al.~1990; Konar \\& Bhattacharya 1999). This model predicts changes in alignment of the magnetic field and spin axis (perhaps involving ``plate tectonics'' of the crust, e.g.~Chen \\& Ruderman 1993).  The importance of understanding evolution of currents in the crust for these different models motivated our recent study (Cumming, Arras, \\& Zweibel 2004, hereafter CAZ04). Whereas calculation of Ohmic decay is straightforward, the evolution of currents in the crust is complicated by the non-linear Hall effect, first studied by Jones (1988) and GR92. This complex process in fact has very simple underlying physics: in the crust, where the ions are held fixed in the solid lattice, {\\em the magnetic field is frozen into the electron fluid}. To see this, combine the Hall electric field (familiar from the simple laboratory experiment) $\\vec{E}_H=\\vec{v}_e\\vcross\\vec{B}/c$, where $\\vec{v_e=J/n_ee}$ is the electron velocity, $\\vec{J}$ is the current density, and $n_e$ is the electron density, with Faraday's law, giving $\\partial\\vec{B}/\\partial t=-c\\curl{\\vec{E}_H}=-\\curl{\\vec{v_e}\\vcross\\vec{B}}$. Therefore the field $\\vec{B}$ evolves due to the currents $\\vec{J}\\propto\\curl\\vec{B}$. An initial dipole field will spontaneously ``twist'' and ``buckle'' as the currents distort the dipole field lines inside the star, generating higher order multipoles (see CAZ04 \\S 4.1).  How fast is this evolution, and when is the Hall effect important? The Hall effect always dominates for magnetar-strength fields $\\sim 10^{14}$--$10^{15}\\ \\mathrm{G}$. However, for $B<10^{13}\\ \\mathrm{G}$, the importance of the Hall effect depends sensitively on the composition of the crust. A simple timescale estimate is given by writing the Hall time across a pressure scale height at the base of the crust, $t_{\\mathrm{Hall}}\\sim 10\\ \\mathrm{Myrs}/B_{12}$ where $B_{12}$ is $B$ in units of $10^{12}\\ \\mathrm{G}$ (see also GR92). The Ohmic decay time is $t_{\\mathrm{Ohm}}\\approx 2\\ \\mathrm{Myrs}/T_8^2$ when phonons dominate the electrical conductivity (temperatures $T_8=T/10^8\\ \\mathrm{K}\\gtrsim 1$), and $t_{\\mathrm{Ohm}}\\approx 6\\ \\mathrm{Myrs}/Q$ when impurities dominate ($T_8\\lesssim 1$), where the impurity parameter $Q$ measures the level of impurities in the crystal lattice (e.g.~Itoh \\& Kohyama 1993). Figure 1 shows curves of $t_{\\rm Ohm}=t_{\\rm Hall}$ for different $Q$, and rough locations of different types of neutron star\\footnote{We choose a density $\\rho=10^{14}\\ \\mathrm{g\\ cm^{-3}}$ (near the base of the crust) in Figure 1. The ratio $t_{\\rm Ohm}/t_{\\rm Hall}\\propto \\rho^{1/6}$ for phonon scattering, and $\\propto 1/\\rho^{2/3}$ for impurity scattering.}.  Figure 1 shows that recent work on the crust composition in both isolated and accreting neutron stars has a direct impact on our understanding of field evolution in the crust. Accreting neutron stars are hot ($T_8\\gtrsim 1$ for accretion rates $\\dot M\\gtrsim 10^{-11}\\ M\\odot\\ \\mathrm{yr^{-1}}$), but also likely have impure crusts $Q\\gtrsim 1$, since their crusts are replaced by a mixture of heavy elements made by hydrogen and helium burning on the surface of the neutron star (Schatz et al.~1999 found $Q\\sim 100$). Therefore, the Hall effect is never important for an accreting star, and the Ohmic time remains short even after accretion switches off. The original estimate of $Q$ for an isolated neutron star by Flowers \\& Ruderman (1977) gave $Q\\sim 10^{-3}$, in which case Hall effects dominate in a radio pulsar which has cooled to $T_8\\lesssim 1$, and a cascade to small lengthscales is expected. However, recent calculations by Jones (2004, and references therein) give $Q\\gtrsim 1$ for an isolated neutron star, in which case an extensive Hall cascade is less likely, with Ohmic decay dominating for $B_{12}\\lesssim Q\\rho_{14}^{2/3}$.  When Hall effects dominate the evolution, the field is expected to grow more and more complex. GR92 suggested that a turbulent ``Hall cascade'' would transport magnetic energy to small scales where Ohmic dissipation is very efficient. For a constant electron density, this cascade has now been confirmed in numerical simulations (Biskamp et al.~1999). However, unfortunately it is difficult to draw firm conclusions about this evolution. Many fascinating physics issues remain to be resolved, including the nature of the cascade (perhaps strong and anisotropic), the behavior of Hall waves in the presence of a steep density gradient and at the boundaries of the crust, and the response of the crust to the wave for strong fields (Rheinhardt \\& Geppert 2000; Vainshtein, Chitre, \\& Olinto 2000; CAZ04). These issues are particularly important for magnetars, since the time for the Hall effect to operate is very short (and always dominates Ohmic decay, independent of $Q$), and perhaps drives the magnetic field decay believed to power these sources (Thompson \\& Duncan 1996; see Arras, Cumming, \\& Thompson 2004 for a recent model).  \\begin{figure} \\begin{center} \\epsfxsize=3.5in \\epsfbox{btproc_rev.ps} \\end{center} \\caption{Regimes of dissipation of crustal currents: Ohmic decay vs.~Hall effect for different neutron stars (adapted from CAZ04). The dotted lines show the magnetic field $B$ above which $t_{\\rm Hall}<t_{\\rm Ohm}$ as a function of temperature in the crust (for a density $\\rho=10^{14}\\ {\\rm g\\ cm^{-3}}$), and different impurity parameters $Q$. Boxes mark the approximate location of different types of neutron star. } \\end{figure} ",
        "conclusions": ""
    },
    "0404/astro-ph0404032_arXiv.txt": {
        "abstract": "The physical process leading to the acceleration of the expansion of the universe is unknown.  It may involve new high energy physics or extensions to gravitation.  Calling this generically dark energy, we examine the consistencies and relations between these two approaches, showing that an effective equation of state function $w(z)$ is broadly useful in describing the properties of the dark energy.  A variety of cosmological observations can provide important information on the dynamics of dark energy and the future looks bright for constraining dark energy, though both the measurements and the interpretation will be challenging. We also discuss a more direct relation between the spacetime geometry and acceleration, via ``geometric dark energy'' from the Ricci scalar, and superacceleration or phantom energy where the fate of the universe may be more gentle than the Big Rip. ",
        "introduction": "\\label{sec.intro}  The acceleration of the expansion of the universe poses a fundamental challenge to the standard models of both particle physics and cosmology.  There is no established framework for the new physics required, but we know that within Einstein's field theory we can treat modifications to either the right hand side -- energy-momentum components, e.g.\\ a scalar field called quintessence -- or the left hand side -- the geometry of spacetime.  Certainly the cosmological constant $\\Lambda$ is equally at home in either location.  Here we examine what future observations can teach us about dark energy and the differences between a new physical component and an extension of gravitation theory.  We give a very brief summary of dark energy theory, moving on to dark energy phenomenology -- how to describe it in a manner amenable to both theory and observations, then dark energy fantasy -- can we constrain complexities in the description, and finally dark energy reality -- what we might expect to learn in the next decade. ",
        "conclusions": ""
    },
    "0404/astro-ph0404562_arXiv.txt": {
        "abstract": "Among the many new irregular satellites that have been discovered in the last five years, at least six are in the so-called Kozai resonance.  Due to solar perturbations, the argument of pericenter $\\omega$ of a satellite usually precesses from 0 to $360^{\\circ}$.  However, at inclinations higher than $\\simeq 39.3^{\\circ}$ and lower than $\\simeq 140.7^{\\circ}$ a new kind of behavior occurs for which the argument of pericenter oscillates around $\\pm 90^{\\circ}$.  In this work we will concentrate on the orbital history of the saturnian satellite S/2000 S5 Kiviuq, one of the satellites currently known to be in such resonance  Kiviuq's orbit is very close to the separatrix of the Kozai resonance. Due to perturbations from the other jovian planets, it is expected that orbits near the Kozai separatrix may show significant chaotic behavior.  This is important because chaotic diffusion may transfer orbits from libration to circulation, and vice versa.  To identify chaotic orbits we used two well-known methods: the Frequency Analysis Method (Laskar 1990) and Maximum Lyapunov Exponents (Benettin {\\em et al.} 1980). Our results show that the Kozai resonance is crossed by a web of secondary resonances, whose arguments involve combinations of the argument of pericenter, the argument of the Great Inequality ($2{\\lambda}_{Jupiter}-5{\\lambda}_{Saturn}$), longitude of the node $\\Omega$, and other terms related to the secular frequencies $g_5$, $g_6,$ and $s_6$. Many test orbits whose precession period is close to the period of the Great Inequality (883 yrs), or some of its harmonics, are trapped by these secondary resonances, and show significant chaotic behavior.  Because the Great-Inequality's period is connected to the semimajor axes of Jupiter and Saturn, and because the positions of the jovian planets have likely changed since their formation (Malhotra 1995), the phase-space location of these secondary resonances should have been different in the past.  By simulating the effect of planetary migration, we show that a mechanism of sweeping secondary resonances, similar to the one studied by Ferraz-Mello {\\em et al.} (1998) for the asteroids in the 2:1 mean motion resonance with Jupiter, could significantly deplete a primordial population of Kozai resonators and push several circulators near the Kozai separatrix. This mechanism is not limited to Kiviuq's region, and could have worked to destabilize any initial population of satellites in the Kozai resonance around Saturn and Jupiter. ",
        "introduction": "Carruba {\\em et al.} (2002) studied an analytical model of the Kozai resonance (Kozai 1962) based on a secular development of the disturbing function, expanded in series of $a/a'$ (where the $'$ stands hereafter for quantities related to the perturber), truncated at second order in $a/a'$, and averaged over the mean anomalies of both perturber and perturbee (see also Innanen {\\em et al.} 1997). Two kinds of behavior are possible in this simplified model: for inclinations of less than 39.23$^{\\circ}$ (or higher than 140.77$^{\\circ}$ for retrograde satellites), the argument of pericenter freely circulates from 0$^{\\circ}$ to 360$^{\\circ}$, while at intermediate inclinations a new class of solutions, in which the argument of pericenter librates around $\\pm$90$^{\\circ}$, are possible.  Such behavior is called the Kozai resonance. The value of the critical inclination of 39.23$^{\\circ}$ is, in fact, an artifact of the averaging; the real boundary between the region in which libration is possible or not is actually a more complicated function of the satellite's semimajor axis and eccentricity (\\'{C}uk {\\em et al.} 2003, \\'{C}uk and Burns 2004). Currently, six irregular satellites of jovian planets are known numerically to be in the Kozai secular resonance: two around Saturn (S/2000 S5 Kiviuq, and S/2000 S6 Ijiraq), two around Jupiter (S/2001 J10 Euporie, and S/2003 J20), and two around Neptune (S/2002 N2 and S/2002 N4, Holman {\\em et al.} 2004).  A third neptunian satellite, S/2003 N1 is on a chaotic orbit that may also be in the Kozai resonance (Holman {\\em et al.} 2004).  A question that Carruba {\\em et al.} left unanswered regarded the possible presence of chaos and chaotic diffusion at the separatrix between the regions of circulation and libration. Since in that work the authors considered a three-body problem, but averaged over the mean anomalies of both perturber and perturbee, our analytical model was an integrable one-degree of freedom model. In the real system, however, perturbations from other jovian planets and short-period terms may produce chaos near the Kozai separatrix. Chaotic evolution might lead to the escape of objects whose orbits were originally inside the Kozai resonance. Here we test the hypothesis that the chaotic layer is due to the fact that the Kozai resonance overlaps with secondary resonances near the separatrix.  It is well known that resonance overlap is one of the major sources of chaos in dynamical systems (Chirikov 1979). Moreover, the positions of the planets may have changed after their formation, due to gravitational scattering of planetesimals (Malhotra 1995). As a consequence of the different planetary positions, the shape of the chaotic layer and the locations of the secondary resonances inside and outside the Kozai resonance should have been different in the past.  This might have had consequences on the stability of any primordial population of Kozai resonators.  This paper will concentrate on identifying the chaotic layer at the transition between circulation and libration for Kozai resonators, in particular for the saturnian satellite Kiviuq, whose orbit is very close to the separatrix of the Kozai resonance.  Our goal is to understand the origins of chaos and the effects of planetary migration.  Our investigation is divided in the following way:  section 2 describes the major features of our simplified model of the Kozai resonance and identifies the transition region between circulation and libration.  Section 3 recalls two methods to identify chaotic behavior in dynamical systems, the Frequency Analysis Method and the maximum Lyapunov exponents.  Section 4 investigates possible causes of the chaotic layer.  Section 5 shows how the Kozai resonance is crossed by a series of secondary resonances of different strengths.  In the final section we demonstrate how the effect of planetary migration, combined with the effect of the secondary resonances, might have operated to depopulate a possible primordial satellite group of Kozai resonators. ",
        "conclusions": "We have analyzed the orbital behavior of satellites inside and outside the Kozai resonance for the phase space around the Saturnian satellite Kiviuq.  Our goal was to identify areas of chaotic behavior and to understand if chaotic diffusion might have played an important role, now or in the past, for orbital stability inside the Kozai resonance. By applying the frequency analysis method (FAM) and computing the maximum Lyapunov exponents, we identified several secondary resonances.  The major source of chaos for orbits in the circulating region are secondary resonances involving the argument of pericenter $\\omega$ (= $\\varpi-\\Omega$), the Great-Inequality's argument ($2{\\lambda}_{J}-5{\\lambda}_{S}$), the longitude of the node ($\\Omega$), and the planetary secular frequencies ($g_5, g_6$, and $s_6$).  The region of transition between circulation and libration of $\\omega$ is strongly chaotic, when perturbations from the other jovian planets are considered. According to our secular model, variations in Saturn's semiminor axis (and, therefore in its eccentricity and semimajor axis) are what drive the chaotic behavior, especially the transfer from circulation to libration. While secular variations modulated by the $g_5, g_6, $ and $s_6$ frequencies are not negligible, once again our results indicate that the dominant effect is connected with perturbations having the Great-Inequality's period.  Also, contrary to what might be expected from the secular model, our results indicate that the chaotic layer at the boundary between circulation and libration is not symmetric, but shows an asymmetry between values of $\\omega$ larger and smaller than $90^{\\circ}$.  This is a consequence of other perturbations connected with secular pericentric resonance, evection inequality, and variations in Saturn's semimajor axis, which are strongly dependent on the initial conditions for particles' $\\varpi$ and the Sun's position.  Our simulations suggest that the dominant effect is connected with the secular resonance, but variations in Saturn's semimajor axis are not entirely negligible.  From our simulations we can introduce the following scenario for chaotic orbits near Kiviuq's Kozai resonance.  Starting from librating orbits close to the libration center, we possibly identify two weak secondary resonances of argument $2(\\varpi-\\Omega)-3(5\\lambda_S-2\\lambda_J)+8\\Omega_S+\\varpi_S$ and $3(\\varpi-\\Omega)-4(5\\lambda_S-2\\lambda_J)+11\\Omega_S+\\varpi_S$ that introduce chaotic behavior.  The first strong source of chaos is however connected with the transition region between circulation and libration: orbits start to switch back and forth from circulation to libration and the whole region is characterized by high values of $\\sigma$.  Once in the circulating regime, a strong resonance of argument $3(\\varpi-\\Omega-5\\lambda_S+2\\lambda_J+3\\Omega_S)+(\\Omega-\\Omega_S)$ appears at $x$ = 0.73. In addition to that, the dominant source of chaos, excluding the nearby pericentric secular resonance, is connected to a secondary resonance with the Great Inequality of argument $\\varpi-\\Omega-5\\lambda_{S}+2\\lambda_{J}+3\\Omega_{S}$ (these orbits have argument-of-pericenter periods of order 840 yrs).  Orbits in the region of this secondary resonance show some of the highest $\\sigma$ values among the test particles we studied. Other secondary resonances connected with the Great Inequality are also observed.  The fact that the dominant source of chaos for orbits in Kiviuq's region is connected with the Great-Inequality secondary resonance opens interesting perspectives.  It is believed that jovian planets formed in different locations than their current ones, and then migrated by gravitationally scattering the primordial population of planetesimals (Malhotra 1995).  The possibility that planets occupied different semimajor axes in the past has important repercussions, since the Great-Inequality's period is connected with the osculating semimajor axes of Jupiter and Saturn.  In particular, when the two planets were closer together, the period of the Great Inequality was shorter.  We simulated this effect first by changing the position of Jupiter so that the Great-Inequality's period was 640 yrs and 480 yrs, respectively (``static integrations'').  These simulations showed that, indeed the position of the Great-Inequality's secondary resonance shifted higher.  A twin resonance for librating particles appeared when the Great-Inequality's period equaled the average libration period (640 yrs), and the minimum (480 yrs) (this resonance is not visible today because the maximum period of libration inside the Kozai resonance is of $\\simeq 820$ yrs).  Then we fully integrated Jupiter and Saturn with an integrator that simulates planet migration according to the exponential law of Malhotra (1995).  Our results show that a mechanism of sweeping resonances, analogous to the one that operated in the Kuiper Belt (Malhotra 1995), or inside the 2:1 mean motion resonance with Jupiter (Ferraz-Mello {\\em et al.} 1998) must have operated as well inside the Kozai resonance. Our simulations show that 15\\% of the satellites originally inside the Kozai resonance became circulators or exhibited switching behavior when planets acquired their final orbits. Indeed several particles were captured into the secondary resonance with the Great Inequality (or other secondary resonances) and, as a consequence, their amplitudes of oscillation in $\\omega$ increased until the test particles reached the boundary between circulation and libration, where escape became possible.  Several particles on circulating orbits were also pushed toward the separatrix by an analogous mechanism of sweeping secondary resonances (but none crossed the boundary to become a librator). We believe that a similar mechanism might have acted in the past also for the region of phase space associated with the other jovian satellites currently inside or near the Kozai resonance. This process could explain why so many irregular satellites are currently found near the Kozai separatrix (\\'{C}uk and Burns 2004).  In this work we used several numerical tools.  In a few cases, like for example the median filter for our FAM results, we introduced a technique, that, while employed in other fields (for example in remote-sensing and in the analysis of data from martian probes), was not, to our knowledge, used previously in celestial mechanics for reducing FAM results. Still on the technical side, we created new integrators to account for the Bretagnon solutions of the giant planets and to simulate planetary migration (Malhotra 1995).  We believe these tools might be useful for other studies, not only for the irregular satellites of jovian planets.  Several questions remain that might be the subjects of future research. In this work we could not compute reliable capture probabilities for the Great-Inequality resonances when planetary migration is considered, since short-period variations in the semimajor axes of Jupiter and Saturn should also be taken into account (Hahn and Malhotra 2000, Beaug\\'{e} {\\em et al.} 2002). It would also be interesting to extend our studies to the neighborhoods of the other jovian satellites currently inside the Kozai resonance, and see if a web of secondary resonances (possibly involving the Lesser Inequality, the 2:1 quasi-resonance between Neptune and Uranus) is also present for the case of those neptunian satellites in Kozai resonance (Holman {\\em et al.} 2004).  We hope that this work might have opened interesting new prospects for the study of origins and stability of irregular satellites in Kozai resonance."
    },
    "0404/astro-ph0404612_arXiv.txt": {
        "abstract": "We show the effectiveness of using the procedure of the potential reconstruction (the V-method) in the estimation of the transition redshift corresponding to switch from deceleration to acceleration phase of the Universe. We investigate the FRW models with dark energy by applying the particle-like description of their dynamics. In this picture the evolution of the universe is represented by motion of a fictitious particle in the one-dimensional potential $V(a)$ or $V(z)$ where $a$ is the scale factor and $z$ is the redshift. The V-method solves the inverse problem where we reconstruct the potential function from empirical data. We use the Riess et al.'s gold and gold+silver samples of SN Ia data in this reconstruction. In the same framework we obtain both the estimates of the form of the equation of state for dark energy and the transition redshift. We obtain that transition redshift is $0.38_{-0.07}^{+0.10}$ (gold subset) and $0.37_{-0.07}^{+0.06}$ (gold+silver sample) when the linear model of dark energy ($w(z)=w_{0} + w_{1}z$) and $\\Omega_{m,0}=0.3$ are assumed. We compare the estimation of transition redshift with the Riess et al. results. The V-method also alows us to find the value of Hubble function at the moment of transition. ",
        "introduction": "The recent supernovae observations indicate that our universe is currently accelerating \\cite{riess98,perlmutter99}. However, from the standard cosmological model we know that earlier the universe decelerated because of domination of matter, and there should exist the moment when the switch from the deceleration phase to the acceleration phase takes place. The value of redshift for this transition $z_{T}$ has been recently found by Riess et al. \\cite{riess04}. They obtained $z_{T} = 0.46 \\pm 0.13$ using a kinematic model independent on the content of the universe. Their method does not allow to obtain both the equation of state parameters and the transition redshift in the same framework. The dark energy equation of state parameters are estimated in models with different ansatz on $w(z)$. However the constraint on the transition redshift was obtained using other relation for $d_{L}(z)$. Let us note that there is no reason to treat the value of $z_{T}$ estimated in a certain model as a real value in another model.  In this paper we propose the method which allows to calculate both equation of state parameters and the transition redshift in the same model. We estimate the value of $z_{T}$ through the potential function $V[a(z)]$ of the Hamiltonian system determining the evolution of the universe. In this approach the evolution of the universe is reduced to the motion of a fictitious particle with a unit mass in a one-dimensional potential. The information about the influence of matter content on the dynamics of the universe is derived from the shape of the potential function $V[a(z)]$. Namely the scale factor $a$ accelerates (decelerates) in the interval of $z$ if $V[a(z)]$ is an increasing (decreasing) function of $z$. Therefore we should expect the potential function has a maximum for the universe with deceleration during its earlier epoch and acceleration at present.  In our previous papers \\cite{szydlowski03,szydlowski03c} we developed the potential function method to probe the dark energy. The use of the potential $V[a(z)]$ instead of $w_{X}(z)$ has an advantage because the former suffers less from the smearing effect caused by the double integral which relates $w_{X}(z)$ and $d_{L}(z)$ \\cite{maor01}. In these papers we showed that the dynamics of the Friedmann-Robertson-Walker (FRW) model of the universe, filled with non-interacting, non-relativistic matter with equation of state $p_{m}=w_{m}(a)\\rho_{m}$ and dark energy with the equation of state $p_{X}=w_{X}(a)\\rho_{X}$, can be reduced to the Hamiltonian system \\begin{equation} \\label{eq:3} \\mathcal{H} = \\frac{\\dot{a}^{2}}{2} + V(a), \\end{equation} where an overdot denotes the differentiation with respect to the cosmic time $t$, and the potential \\begin{equation} \\label{eq:4} V(a) = - \\frac{\\rho_{\\rm eff} a^{2}}{6}, \\quad \\rho_{\\rm eff}(a) = \\rho_{m}(a) + \\rho_{X}(a), \\end{equation} where the units $8\\pi G = c =1$ are used and $\\rho_{\\rm eff}$ is the effective energy density of the mixture of non-interacting fluids which satisfy the conservation equation \\begin{equation} \\label{eq:5} \\dot{\\rho}_{i} = -3 \\frac{\\dot{a}}{a} (\\rho_{i} + p_{i}). \\end{equation} Then $\\rho_{\\rm eff}$ also satisfies equation (\\ref{eq:5}). On the other hand the potential function $V(a)$ can be reconstructed from the SN Ia data through the Hubble function $H(z)$ which in the spatially flat universe is related to the luminosity distance $d_{L}(z)$ as \\begin{equation} \\label{eq:2} H(z) = \\left[\\frac{d}{dz} \\left( \\frac{d_{L}(z)}{1+z} \\right)\\right]^{-1}. \\end{equation}  Because the trajectories of system (\\ref{eq:3}) always lie on the zero-energy level (as a consequence of the Hamiltonian constraint) we can rewrite potential (\\ref{eq:4}) to the form \\begin{equation} \\label{eq:6} V(z) = - \\frac{1}{2} \\left[ (1+z) \\frac{d}{dz} \\left( \\frac{d_{L}(z)}{1+z} \\right) \\right]^{-2}. \\end{equation}  The potential $V(z)$ contains all information necessary to determine the full dynamics of the universe on the phase plane $(a,\\dot{a})$. Thus both the effective energy density $\\rho_{\\rm eff}(a)$ and the equation of state coefficient $w_{\\rm eff}[a(z)] \\equiv p_{\\rm eff}/\\rho_{\\rm eff}$ can be calculated unambiguously from the potential function $V(a)$ \\begin{equation} \\label{eq:7} \\rho_{\\rm eff}(a) = - \\frac{6V(a)}{a^{2}}, \\quad w_{\\rm eff}(a) = - \\frac{1}{3} [1 + I_{V}(a)], \\end{equation} where $I_{V}(a) = (d \\ln{V})/(d \\ln{a})$ is the elasticity of the potential function $V$ with respect to the scale factor $a$.  As it is well known the behavior of the potential function in the neighborhood of a maximum can be approximated by a quadratic part of the expansion in Taylor's series. Therefore we have \\begin{equation} \\label{eq:8} V(a) = V(a_{T}) + \\frac{1}{2} \\left. \\frac{\\partial^{2} V}{\\partial a^{2}} \\right|_{a=a_{T}} (a-a_{T})^{2} \\end{equation} where $V(a_{T}) = -\\dot{a}_{T}^{2}/2$ from the Hamiltonian constraint, and \\begin{equation} \\label{eq:9} \\rho_{\\rm eff}(a) = - \\frac{6}{a^{2}} \\left[ -\\frac{\\dot{a}_{T}^{2}}{2} + \\frac{1}{2} \\left. \\frac{\\partial^{2} V}{\\partial a^{2}} \\right|_{a=a_{T}} (a-a_{T})^{2} \\right] \\end{equation} or \\begin{align} \\rho_{\\rm eff}(z) =& - 3 (1+ z_{T})^{4} \\left. \\frac{\\partial^{2} V}{\\partial z^{2}} \\right|_{z=z_{T}} \\nonumber \\\\ & + 3 (1+z)^{2} \\left[ \\frac{1}{(1+z_{T})^4} \\left( \\left. \\frac{d z}{d t} \\right|_{z=z_{T}} \\right)^{2} - (1+z_{T})^{2} \\left. \\frac{\\partial^{2} V}{\\partial z^{2}} \\right|_{z=z_{T}} \\right] \\nonumber \\\\ & - 6(1+z) (1+z_{T})^{3} \\left. \\frac{\\partial^{2} V}{\\partial z^{2}} \\right|_{z=z_{T}}. \\label{eq:10} \\end{align}  In the case of transition from deceleration to acceleration one can see from formula~(\\ref{eq:10}) that there are three terms which can be interpreted as the cosmological constant (positive), the curvature contribution (negative), and the topological deffects with positive energy ($p=-2/3 \\rho$). Note that as $z$ is close to $-1$ ($a \\gg a_{0}$) than the last two terms are negligible and only the ``cosmological constant'' dominates the dynamics of the late epoch of the universe.  The value of the potential function at the maximum is given by \\begin{equation} \\label{eq:11} V(a_{T}) = -\\frac{1}{2} H_{T}^{2}(1+z_{T})^{-2}. \\end{equation} The value of $z_{T}$ and $V(z_{T})$ inform us about the moment of transition and the Hubble function at this moment. The estimation of these parameters will be done in the next section. ",
        "conclusions": "We presented the V-method of reconstruction of the potential function. We applied it to the estimation of the state parameters of the cosmological model with dark energy. We concentrated on the estimation of the position of the maximum of the potential function which can be interpretated as the moment of switch from the deceleration to the acceleration phase during the Universe evolution.  The shape of the potential function is reconstructed from the distant SN Ia data. For our analysis we used the samples of SN Ia prepared by Riess et al.  The Hamiltonian formulation of the dynamics enables us to determine whole dynamics of the cosmological model from the potential function only. It is interesting that this function determines not only the qualitative structure of the phase space \\cite{szydlowski03,szydlowski03c} but also gives us some quantitative information about the model parameters.  We showed that the V-method is very effective in analysis of cosmological models with dark energy where we can find the moment of transition as the quantitave parameter. The other obtained parameter of state of the system is the velocity of the scale factor (or the Hubble parameter) at the moment of transition. We can treat these two estimated parameters as initial conditions and determine uniquely the evolution path of the model.  The main results of using this method to the cosmological model with dark energy are  1. The method allows to study quantitative and qualitative aspects of dynamics of models.  2. We showed that the state parameters at the moment of transition can be obtained in any models for which the reconstruction can be done.  3. We estimated the value of the transition redshift and the value of Hubble function at the transition moment moment.  Riess et al. estimated the model parameters assuming the different forms of the equation of state coefficient $w(z)$, but to find the transition redshift $z_{T}$ they used a certain kinematic model with a linear expansion for $q(z)$. It seems to be no a priori reasons to extrapolate the value of $z_{T}$ between two different models. The question is to which $w(z)$ model corresponds their value of $z_{T}$.  Our methods overcomes this obstacle. To obtain the model parameters and the transition redshift we use the same $d_{L}(z)$ relation."
    },
    "0404/astro-ph0404338_arXiv.txt": {
        "abstract": "Simulations in general relativity show that the outcome of collapse of a marginally unstable, uniformly rotating star spinning at the mass-shedding limit depends critically on the equation of state. For a very stiff equation of state, which is likely to characterize a neutron star, essentially all of the mass and angular momentum of the progenitor are swallowed by the Kerr black hole formed during the collapse, leaving nearly no residual gas to form a disk. For a soft equation of state with an adiabatic index $\\Gamma - 4/3 \\ll 1$, which characterizes a very massive or supermassive star supported predominantly by thermal radiation pressure, as much as $10\\%$ of the mass of the progenitor avoids capture and goes into a disk about the central hole. We present a semi-analytic calculation that corroborates these numerical findings and shows how the final outcome of such a collapse may be determined from simple physical considerations. In particular, we employ a simple energy variational principle with an approximate, post-Newtonian energy functional to determine the structure of a uniformly rotating, polytropic star at the onset of collapse as a function of polytropic index $n$, where $\\Gamma = 1+1/n$. We then use this data to calculate the mass and spin of the final black hole and ambient disk. We show that the fraction of the total mass that remains in the disk falls off sharply as $3-n$ (equivalently,  $\\Gamma - 4/3$) increases. ",
        "introduction": "Determining the final state of a rotating star undergoing gravitational collapse is an important issue in relativistic astrophysics. Such a collapse is the principle route by which a rotating black hole forms in nature. Several different scenarios involving the collape of rotating stars to black holes have been the focus of considerable attention recently. The collapse of massive stars in hypernovae explosions (``collapsars'') may be the origin of long-period gamma-ray bursts (MacFadyen \\& Woosley 1999; MacFadyen, Woosley \\& Heger 2001). The remnant of a binary neutron star, following inspiral and coalescence, is likely to be rapidly rotating and undergo collapse, either prompt or delayed, depending on its mass (Shibata \\& Uryu 2000, 2002; Baumgarte, Shapiro \\& Shibata 2000). Binary neutron star merger and collapse provides a plausible scenario for generating short-period gamma-ray bursts (Narayan, Paczynski \\& Piran 1992; Ruffert \\& Janka 1999) and a good candidate for the detection of gravitational waves by Advanced LIGO and other high-frequency, gravitational wave laser interferometers.  The collapse of a rotating supermassive star may be a promising route for forming the seeds of supermassive black holes (see Shapiro 2003 for a review and references), which are likely to reside at the centers of many, and perhaps most, bulge galaxies (Richstone et al. 1998; Ho 1999), including the Milky Way (Genzel et al. 1997; Ghez et al. 2000; Sch\\\"odel et al. 2002) and are believed to be the engines that power active galalactic nuclei (AGNs) and quasars (Rees 1998, 2001).  General relativity induces a radial instability in compact stars to catastrophic collapse. Rotation can support larger masses in stable equilibrium, but all stars become unstable to collapse when they are sufficiently compact. Uniformly rotating stars can have masses exceeding the maximum allowed mass of a nonrotating spherical star by about $\\lesssim 20\\%$. Cook, Shapiro \\& Teukolsky (1994a,b) have constructed numerical solutions of such ``supramassive'' stars, both for polytropes and for 14 realistic candidate nuclear equations of state. Differentially rotating stars can be ``hypermassive'', with masses exceeding the maximum value for supramassive stars by factors or two or more (Baumgarte, Shapiro \\& Shibata 2000; Lyford, Baumgarte \\& Shapiro 2003). Viscosity (molecular or turbulent) and magnetic fields drive differentially rotating configurations to uniform rotation on secular timescales. [See, e.g., Shapiro 2000 and Cook, Shapiro \\& Stephens 2003 for recent Newtonian calculations of magnetic braking and viscous damping of differential rotation, and  Liu \\& Shapiro 2004 for approximate, general relativistic calculations of these phenomena; see Duez et al. 2004 for detailed, general relativistic, numerical simulations of viscous damping of differential rotation.]  Here we focus on uniformly rotating stars and consider the collapse of supramassive stars. The first investigation of the gravitational collapse of supramassive stars was performed by Shibata, Baumgarte and Shapiro (2000), who employed a fully relativistic hydrodynamics code in three spatial dimensions plus time. They studied the collapse of a marginally unstable, relativistic polytrope with polytropic index $n=1$ spinning at the mass-shedding limit. The mass--shedding limit is the maximal spin rate for a uniformly rotating star; here matter on the equator moves in a Keplerian (geodesic) orbit about the star, supported against gravity entirely by centrifugal forces. Stars with $n=1$ have stiff equations of state, are not very centrally condensed and are fairly relativistic at the onset of collapse, with $R_e/M \\approx 5.6$, where $R_e$ is the equatorial (circumferential) radius and $M$ is the total mass-energy. Such stars provide crude models of rapidly rotating, massive neutron stars. The outcome of collapse is a Kerr black hole containing all the rest-mass $M_0$ of the initial configuration, and essentially all of the total mass-energy $M$ and angular momentum $J$ of the initial configuration (apart from the small amount radiated away by gravitational waves). The reason that no matter remains outside the hole is that at the equator the specific angular momentum of the gas, which is strictly conserved during axisymmetric collapse, is smaller than $j_{\\rm ISCO}$, the specific angular momentum of a particle at the innermost stable circular orbit {\\rm (ISCO)} about the final hole. Hence all the interior matter in the star is captured by the hole. This situation is similar to the collapse of a nonrotating spherical star, which forms a Schwarzschild black hole without any ambient disk.  By contrast, Shibata \\& Shapiro (2002) performed a fully relativistic hydrodynamics simulation in axisymmetry of the collapse of a marginally unstable $n=3$ polytrope at the mass-shedding limit. Such a configuration models the likely endpoint state of an evolved, radiation-dominated, very massive or supermassive star at the onset of collapse (Baumgarte \\& Shapiro 1999).  By contrast with an $n=1$ polytrope, an $n=3$ configuration has a soft equation of state, is very centrally condensed and is nearly Newtonian ($R_e/M \\approx 620$) at the onset of collapse.  A rotating black hole again forms during the collapse, but in this case the mass of the hole contains only about $90\\%$ of the total rest-mass of the system (with a spin parameter $J_h/M_h^2 \\sim 0.75$). The remaining gas forms a rotating disk about the nascent hole. Here the angular momentum of the infalling matter in the outermost layers of the collapsing star exceeds $j_{\\rm ISCO}$, so the gas in the outer regions escapes capture and forms a disk. This result has been corroborated by a semi-analytic calculation by Shapiro \\& Shibata (2002), which we generalize below for arbitrary $n$.  Recently, Shibata (2003) performed systematic numerical simulations in axisymmetry to investigate the effects of the stiffness of the equation of state on the outcome of the collapse of marginally unstable, supramassive stars spinning at the mass-shedding limit. He studied collapsing relativistic polytropes with indicies $2/3 \\leq n \\leq 2$, representing stiff and moderately stiff equations of state. He concluded that for these configurations the final state is a Kerr black hole and the disk mass, if it exists at all, is very small ($ < 10^{-3}$ of the initial mass).  Simulations performed to date clearly show that the outcome of the collapse of a maximally spinning, supramassive star depends critically on the equation of state. In this paper we employ a straightforward semi-analytic analysis to corroborate this conclusion and to show how the final outcome of these detailed simulations may be inferred from simple physical considerations. In particular, we employ an energy variational principle to determine the approximate structure of the marginally unstable, rotating star at the onset of collapse as a function of polytropic index $n$ and use this data to estimate the mass and spin of the final black hole and ambient disk. We show that the mass of the disk depends sensitively on the difference $3-n$. Our semi-analytic calculations are based on an approximate post-Newtonian energy functional valid for rotational energies that are much smaller than gravitational potential energies. They are most accurate for soft equations of state with $n$ close to 3, since the parameters of the progenitor star are nearly Newtonian in this case and the effect of rotation in distorting the shape of the interior region containing most of the mass is minimal. Interestingly, stars with $n$ close to 3 are the most difficult to simulate with a relativistic hydrodynamics code, since the large size of the progenitor star requires that any such code span an enormous dynamic range and integrate for a huge number of (Courant) timesteps in order to follow the collapse to completion and resolve the central black hole and disk reliably. In this sense, the two approaches -- the semi-analytic analysis presented below and full numerical simulations -- are complementary. ",
        "conclusions": "We have employed a simple analysis to determine the effect of the stiffness of the equation of state on the fate of the collapse of a marginally unstable, relativistic polytrope spinning uniformly at the mass-shedding limit. We have used a variational principle and an approximate, post-Newtonian energy functional to determine the key parameters defining the structure of the critical progenitor star. We compared these parameters with the results of detailed nummerical model calculations of stationary configurations for select cases. We then employed a Roche model to obain analytic expressions for the density and angular momentum profiles in the envelope of the marginally stable star. We substituted these profiles into quadratures that determine the fractions of the stellar mass and angular momentum which escape capture by the central black hole assumed to form during the collapse. The fraction of the progenitor mass and spin which go into the hole versus the ambient disk is then iterated until convergence is achieved.  We find that for the stars treated here, the mass fraction in the disk is about $10\\%$ for an $n=3$ polytrope and decreases rapidly as $n$ decreases and the equation of state stiffens. The results are in agreement with the numerical simulations in $3+1$ by Shibata, Baumgarte \\& Shapiro (2000) for $n=1$ and simulations by Shibata (2003) in axisymmetry for $2/3 \\leq n \\leq 2$, which show that the mass fraction in a disk, if present at all, is less than $0.1\\%$ of the total mass. For the special case of $n=3$, the results are also in agreement with the simulation performed by Shibata \\& Shapiro (2002) in axisymmetry, as has been discussed previously by Shapiro and Shibata (2002). The spin parameter of the black hole is less sensitive to the stiffness of the equation of state, decreasing from $J_h/M_h^2 \\approx 0.75$ for $n=3$ to $J_h/M_h^2 \\approx 0.39$ for $n=2.5$.  The approach outlined here is more reliable for soft equations of state. In this limit the marginally unstable configuration is nearly Newtonian and the post-Newtonian energy functional describing the bulk of the mass, the Newtonian Roche model describing the envelope and the quadrature determining the escaping mass-energy fraction (assumed to equal the escaping rest-mass fraction) all become better approximations. Interestingly, this is precisely the limit where fully relativistic numerical simulations become more taxing, due to the large dynamic range and very long time integrations required for a hydrodynamic calculation.  While the semi-analytic approach presented here to predict the final black hole and disk masses and spins was applied to treat the collapse of stars with {\\it uniform} rotation, the method can be used equally well to treat the collapse of unstable stars with {\\it differential} rotation. Recently, Duez et al. (2004) performed fully relativistic numerical simulations of hypermassive stars with appreciable differential rotation. Hypermassive stars may form from the merger of binary neutron stars or from rotating core collapse in supernovae (Baumgarte, Shapiro \\& Shibata 2000). Magnetic braking or viscous damping of differential rotation in such stars can drive them unstable to collapse on secular timescales. Duez et al. (2004) performed simulations to demonstrate this effect in the case of viscosity and showed that the final black hole and disk parameters are in agreement with the values predicted by the method outlined in Section \\ref{disk}. In particular, they showed that for rapid differential rotation characteristic of hypermassive stars, the disk mass fraction is typically large ($10\\% -- 20\\%$ of the total initial mass) for stiff equations of state with $n=1$. We therefore conclude that, in general, the parameters of the final black hole and the disk formed during the collapse of an unstable star depend both on the equation of state and on the degree of differential rotation."
    },
    "0404/astro-ph0404424_arXiv.txt": {
        "abstract": "\\noindent Low-resolution, ultraviolet-to-visual spectra of two candidate globular clusters in the merger remnant \\n3921 ($cz_{_{\\rm LG}}=6021$ \\kms) are presented.  These two clusters of apparent magnitude $V\\approx 22.2$ ($M_V\\approx -12.5$) lie at projected distances of about 5~kpc (0.9\\reff) from the center and move with halo-type radial velocities relative to the local galaxy background. Their spectra show strong Balmer absorption lines [EW(H$\\beta$\\,--\\,H$\\delta$)~= 11\\,--\\,13 \\AA] indicative of main-sequence turnoffs dominated by A-type stars.  Comparisons with model-cluster spectra computed by Bruzual \\& Charlot and others yield cluster ages in the range of 200\\,--\\,530~Myr, and metallicities about solar to within a factor of three.  Given their small half-light radii ($\\reff\\la 5$~pc) and ages corresponding to $\\sim$10$^2$ core-crossing times, these clusters are gravitationally bound and, hence, indeed young globulars. Assuming that they had Chabrier-type initial mass functions, their estimated current masses are $(2.3\\pm 0.1)\\times 10^6 \\msun$ and $(1.5\\pm 0.1)\\times 10^6 \\msun$, respectively, or roughly half the mass of $\\omega$ Cen. Since \\n3921 itself shows many signs of being a $0.7\\pm 0.3$~Gyr old protoelliptical, these two young globulars of roughly solar metallicity and their many counterparts observed with the {\\it Hubble Space Telescope} provide supporting evidence that, in the process of forming elliptical-like remnants, major mergers of gas-rich disks can also increase the number of metal-rich globular clusters. ",
        "introduction": "During the past decade it has become increasingly apparent that galactic mergers, bulge building, and the formation of large numbers of star clusters are closely intertwined subjects.  Evidence for systems of luminous blue clusters newly formed in mergers has accumulated both before and especially since the launch of the {\\it Hubble Space Telescope (HST)}. It now appears that of the thousands of clusters formed in ongoing mergers of gas-rich galaxies (e.g., \\n4038/39: \\citealt{ws95}, \\citealt{whit99}; \\n3256: \\citealt{zepf99}) only the densest and more massive will survive for $\\ga$500 Myr \\citep{fall01} and become the young globular clusters seen in merger remnants of recent vintage.  Observations of globular cluster (GC) systems in recent remnants such as \\n1275 \\citep{holt92,carl98}, \\n3597 \\citep{holt96,carl99}, \\n3921 (\\citealt{schw96b}, hereafter SMWF96), and \\n7252 \\citep{whit93,mill97} are especially valuable because in these remnants the starbursts have subsided, the dense gas and dust obscuring the view have diminished, and most of the freshly minted clusters have evolved sufficiently to reveal their true nature.  Over the several 100 Myr a major disk--disk merger takes to complete, loose star clusters and associations tend to disperse, while gravitationally-bound dense clusters survive as a fossil record of the merger's star-formation history.  Hence, determining the ages and metallicities of these surviving clusters is an important first step in trying to reconstruct that history.  This goal has taken added importance since the discovery of bimodal color and metallicity distributions in the GC systems of many giant ellipticals, predicted by \\citet{ashm92} from merger models of E formation. Since in such ellipticals the red globulars typically have metallicities of $Z\\approx 0.2$\\,--\\,2 \\zsun\\ (e.g., \\citealt{forb01}; \\citealt{peng04}), a question of interest is just how metal-rich the second-generation GCs formed during major mergers are.  There is growing evidence that such clusters have roughly solar metallicities, as measured to date for two young globulars in \\n7252 \\citep{ss93,ss98}, one such cluster in \\n1275 \\citep{zepf95,brod98}, three intermediate-age globulars in \\n1316 \\citep{goud01a}, and two in \\n3610 \\citep{stra03,stra04}.  The present paper describes exploratory spectroscopic observations aimed at verifying the nature of two candidate young globular clusters in \\n3921, and at determining their ages and metallicities.  Similar to, but slightly younger than the more well-known \\n7252, \\n3921 is a very recent remnant of two merged disk galaxies \\citep{tt72} and shows many signs of being a $0.7\\pm 0.3$ Gyr old protoelliptical (\\citealt{schw96a}, hereafter S96).  Predicted to fade by only $\\sim$1.5 mag from its present $M_V = -22.1$ over the next 12 Gyr, this remnant is likely to remain a luminous field elliptical.  Its young candidate GCs and looser associations were first discovered from the ground (S96) and then in larger numbers with \\hst\\ (SMWF96).  Our spectroscopic observations pertain to two of its brightest candidate globulars.  \\n3921 itself is located at $\\alpha_{\\rm J2000}=11^{\\rm h}51^{\\rm m}07\\fs0$, $\\delta_{{\\rm J}2000}=+55\\degr04\\arcmin43\\arcsec$ (SMWF96) and has a recession velocity relative to the Local Group of $cz_{_{\\rm LG}} = +6021\\pm 15$ \\kms\\ (S96), which places it at a distance of 86 Mpc for $H_0 = 70$ \\kms\\ Mpc$^{-1}$.  At that distance, adopted throughout the present paper, $1\\arcsec = 417$ pc.  The corresponding distance modulus is $(m-M)_0 = 34.67$. The Milky Way foreground extinction is small, with values in the literature ranging between $A_V=0.00$ (\\citealt{rc3}) and 0.047 \\citep{schl98}.  We adopt the mean value of $A_V = 0.023$, with which the absolute visual magnitude of \\n3921 becomes $M_V = -22.09$.  In the following, \\S 2 describes the observations and reductions, \\S 3 presents results concerning the kinematics, ages, and mean metallicity of the two young clusters, and \\S 4 discusses issues concerning the clusters' physical nature and origin.  Finally, \\S 5 summarizes our main conclusions. ",
        "conclusions": "The present discussion addresses some issues concerning the ages, metallicity, nature, and origin of the two young clusters in \\n3921. The single most interesting question is whether these clusters truly are young globulars.  \\subsection{Cluster Ages}  Given the ages derived in \\S 3.2 and their error bars, it seems likely that clusters S1 and S2 both have ages falling in the range 200\\,--\\,530 Myr.  Over the next 13 Gyr, solar-metallicity clusters in this age range will fade by $\\Delta M_V \\approx 3.9$\\,--\\,3.2 mag from stellar evolution alone (BC03, for \\citealt{chab03} initial mass function), and possibly more if loss of stars due to evaporation is significant (Fall \\& Zhang 2001).  Hence, when S1 and S2 will reach an age comparable to that of most present-day old globulars, they should have absolute magnitudes well within the range typical for such globulars.  Specifically, at age 13.5 Gyr and if evaporation is negligible, Cluster S1 will shine with $M_V\\approx -8.6$ and S2 with $-9.2$.  How secure are the spectroscopically determined cluster ages, and could---by any chance---the real ages lie well outside the above quoted range of 200\\,--\\,530 Myr?  To address these questions, we have performed three checks.  First, a visual comparison of the cluster spectra shown in Fig.~3 with sequences of Magellanic-Cloud cluster spectra arranged by age \\citep{bica86,leon00,leon03} suggests immediately that clusters S1 and S2 do, indeed, have ages within about the quoted range.  The spectra of younger Magellanic-Cloud clusters ($\\tau\\approx 10\\,$--\\,100 Myr) differ significantly from those of S1 and S2 by having much stronger UV continua and weaker K lines, while those of clusters older than $\\sim$600 Myr have stronger K lines, weaker Balmer lines, and redder continua.  Second, a more detailed comparison with an age sequence of model-cluster spectra of solar metallicity (BC03) strengthens this conclusion. Figure~6 shows the observed spectra of S1 and S2 with BC03 model spectra superposed.  In each case, we have simply selected from the age sequence the model spectrum nearest in time to the cluster age and overplotted it. Even without interpolating the model spectra in time or smoothing them, the matches between observed and model spectra are impressive.  These matches deteriorate rapidly when one superposes spectra of model clusters half or twice as old as those shown.  In addition, we have specifically checked that during the much earlier, red-supergiant phase of cluster evolution ($\\tau \\approx 8\\,$--\\,20 Myr) the model spectra remain significantly different from the observed ones, with both the Balmer and \\ion{Ca}{2} K lines weaker and He lines showing up that are not seen in S1 or S2.  Hence, these two spectroscopic checks appear to confirm that the true ages of S1 and S2 do, most likely, lie within the quoted range of 200\\,--\\,530 Myr.  Yet, a third, photometric check is less successful.  From the cluster colors \\vizero\\ given in Table~1 and the same BC03 (or BC96) models used for the spectroscopic analysis above, we derive photometric ages about 1.5\\,--\\,2 times as large as the spectroscopic ages based on Balmer-line equivalent widths: $\\tau_{\\rm phot} = 725\\pm 48$ Myr for Cluster S1, and $510\\pm 28$ Myr for S2 (BC03 models, $Z=\\zsun$, Chabrier IMF).  At least four different factors could cause this discrepancy: reddening, non-solar metallicity, an abnormal initial mass function (IMF), and problems with the models.  We briefly consider these four factors in turn: (1) To bring the photometric ages into agreement with the spectroscopic ages, reddening within \\n3921 would have to be significant, $E_{V\\!-\\!I}=0.15$ for S1 and 0.09 for S2 (corresponding to $A_V=0.36$ and 0.21, respectively), which seems possible given the apparent scatter in cluster colors seen in Fig.~2. (2) To achieve the same agreement through non-solar metallicity, the clusters would have to have $Z\\approx 3\\, \\zsun$, an unusually high value, but still on the edge of the range permitted by our observations (\\S 3.3). (3) A top-heavy IMF might yield abnormally strong Balmer absorption lines, as suspected for young clusters in \\n1275 \\citep{brod98}, thus falsifying the spectroscopic ages. And (4), there is always the possibility that the models used do not reproduce the cluster colors well during the Asymptotic-Giant-Branch (AGB) phase of evolution ($\\tau=10^8$\\,--\\,$10^9$ yr, see \\citealt{mara01}), when for clusters of mass $\\la${}$10^6 \\msun$ stochastic noise in the colors is significant (BC03, esp.\\ Fig.~8). Although we cannot adjudicate between these various possibilities at present, we suspect that a combination of reddening and stochastic color noise may be responsible for most of the discrepancy between the photometric and spectroscopic ages of S1 and S2.  How significant is the age difference between the two clusters, and could they both have the same age?  The formal age difference derived from the cluster ages given in \\S 3.2 is $\\Delta{\\rm Age} = 170\\pm 130$~Myr, which implies that this difference is significant only at the 1.3$\\,\\sigma$ level. Hence, equal or very similar ages for clusters S1 and S2 cannot be excluded with any certainty.  Yet, there is no reason to prefer a near-zero age difference over the most likely 170~Myr age difference, since in major disk mergers the period of enhanced star and cluster formation lasts typically about 200\\,--\\,500~Myr (e.g., \\citealt{miho93}, esp.\\ Figs.\\ 10, 13, and 24; \\citealt{barn04}, esp.\\ Fig.\\ 5).  Hence, even with an age difference of $170\\pm 130$ Myr clusters S1 and S2 are likely to have formed during the same starburst period of the merger.  How do the newly determined spectroscopic ages of S1 and S2 compare with what is known about the merger history of \\n3921? This remnant of a likely S0--Sc or Sa--Sc merger \\citep{hivg96} is a $0.7\\pm 0.3$ Gyr old protoelliptical (S96).  Its ``E\\,+\\,A'' spectrum near the center suggests that a major starburst occurred at some time during a period $\\sim\\,$0.5\\,--\\,1.0 Gyr ago, whereas the central morphological disequilibrium hints at a somewhat more recent coalescence.  Taken literally, the above starburst period and the spectroscopic ages of S1 and S2 would imply that the two clusters formed near the end or even in the aftermath of the main starburst.  However, the uncertainties are large enough that S1 and S2 may well have formed {\\it during} the main starburst. What seems certain is that S1 and S2 are unlikely to have formed before the main starburst, in agreement with the evidence from 102 candidate GCs observed with \\hst\\ (SMWF96).   \\subsection{Metallicity}  The metallicities of clusters S1 and S2 are about solar to within a factor of three (\\S 3.3).  This is not unexpected, given that the clusters formed from molecular gas within the past $\\sim$0.5 Gyr and the two galaxies that merged in \\n3921 were major disks.  The clusters' near-solar metallicities are also consistent with evidence that (1) \\n3921 itself is a protoelliptical (\\citealt{tt72}; S96) and (2) giant ellipticals tend to have globular-cluster systems with bimodal metallicity distributions, one peak of which contains ``metal-rich'' clusters of $[Fe/H]\\ga -1.0$ (e.g., \\citealt{az98,kund01,lars01}). Observations with \\hst\\ indicate that the recent merger in \\n3921 produced $\\ga$100 new globular clusters, increasing the total number of GCs by at least 40\\% (SMWF96). If most of these young GCs are as metal-rich as S1 and S2 are, then the ratio of metal-rich to metal-poor GCs must be $\\ga$0.4, in rough agreement with the average ratio and 1$\\sigma$-range of $0.8\\pm 0.4$ observed in normal elliptical galaxies \\citep{forb97,lars01}. Therefore, within a few gigayears \\n3921 may possess a GC system not unlike that of a giant field elliptical. The metallicity distribution will likely be bimodal, with a population of metal-poor clusters consisting of the GCs formerly belonging to the halos of the now-merged disk galaxies and a population of metal-rich clusters both stemming from these disks and formed during the merger itself.  If so, \\n3921 may be related to elliptical galaxies like \\n1316 \\citep{goud01a}, \\n3610 \\citep{whit97,stra03,stra04}, and \\n5128 \\citep{peng04}, which all appear to be descendants of more ancient mergers involving disk galaxies.  Their common trait is the presence of both fine structure indicative of past mergers and some metal-rich globular clusters of intermediate age ($\\sim$2\\,--\\,8 Gyr).  Though not showing any {\\it obvious} signs of tidal perturbation, the Virgo Cluster elliptical \\n4365 may belong to this category of merger descendants as well. It, too, features metal-rich globular clusters of intermediate age \\citep{lars03} and a kinematically decoupled, slightly bluish central disk that may stem from a past merger \\citep{surm95,caro97}.   \\subsection{Nature and Origin of Clusters}  What is the nature of clusters S1 and S2, and how did they form?  The spectroscopic observations presented above support the view that many of the bluish point sources observed in \\n3921 (S96; SMWF96) indeed are young globular clusters.  In essence, any dense cluster with $\\ga$10$^4$ stars, an effective (= half-light) radius \\reff\\ of the order of 10~pc or less, and an age exceeding 10\\,--\\,20 core-crossing times \\tcross\\ has to be gravitationally bound and is a globular cluster.  Assuming that S1 and S2 have \\citet{chab03} IMFs, their luminosities and ages indicate present-day masses of $(2.3\\pm 0.1)\\times 10^6 \\msun$ and $(1.5\\pm 0.1)\\times 10^6 \\msun$, respectively (BC03), whence they must contain well over 10$^6$ stars each. Both clusters have measured effective radii of $\\reff\\la 5$~pc, implying that---like Milky-Way globulars---they have core-crossing times of a few Myr \\citep{meyl97}.  Hence, each is of order $\\sim$100\\,\\tcross\\ old and clearly a gravitationally bound {\\it young globular cluster}.  According to the BC03 models, the two clusters have so far lost about 35\\% of their original mass due to stellar evolution and will loose another $\\sim$15\\% over the coming 13 Gyr due to the same cause.  At an age of 13.5 Gyr and if stellar evolution were the only source of mass loss, S1 and S2 would be 45\\% and 28\\% as massive as $\\omega$~Cen is at present ($\\mass_{\\omega {\\rm Cen}}\\approx 4.0\\times 10^6 \\msun$, \\citealt{meyl02}). Hence, these are both dense and massive GCs by any measure.  Although their dynamical evolution and evaporative mass loss will depend on their unknown orbital characteristics, it seems likely that both clusters will survive for a long time, of the order of a Hubble time or longer.  How did S1, S2, and the many other candidate young GCs in \\n3921 form? There is growing evidence that massive young clusters in present-day mergers form from Giant Molecular Clouds (GMC; e.g., \\citealt{schw03}; \\citealt{whit03}). As \\citet{jog92} first pointed out and \\citet{elme97} reasoned in more detail, the rapidly mounting pressure of interstellar gas heated in starbursts provides a natural mechanism for triggering the collapse of the embedded GMCs, which are cold and only marginally stable. In support of this hypothesis, the luminosity functions of young globular-cluster systems are remarkably similar to the mass functions of GMCs in present-day spiral galaxies \\citep{harr94}.  Both kinds of functions are power laws with similar exponents and upper-end cutoffs. Specifically, the luminosity function of young clusters in \\n3921 has the form $\\phi(L) dL \\propto L^{\\alpha} dL$, where $\\alpha = -2.1\\pm 0.3$ (SMWF96), in accord with the luminosity functions of many other systems of young clusters and GMCs.  A further piece of evidence in favor of the young globulars of \\n3921 having formed from GMCs in the merging disk galaxies is provided by the clusters' radial distribution.  This distribution closely follows the $r^{1/4}$-law of \\n3921's light distribution.  The agreement indicates that the progenitors of the young GCs experienced the same violent relaxation as did the disk stars, which in turn suggests that these progenitors were compact and reacted to the merger like point masses (SMWF96). The GMCs of the input galaxies are the only known gas-rich objects to be both compact and massive enough to fit the bill.  Hence, the radial cluster distribution in \\n3921 strongly points to GMCs having been the progenitors of the young globulars.  The fact that S1 and S2 have approximately solar metallicities (\\S 3.3) is consistent with this hypothesis."
    },
    "0404/astro-ph0404155_arXiv.txt": {
        "abstract": "{see next page ",
        "introduction": "Blazars are a class of radio-loud active galactic nuclei (AGNs) well known for their peculiar properties, such as their intense and extremely variable non-thermal emission across all the electromagnetic spectrum from the radio band to $\\gamma$-rays (sometimes up to TeV energies). Blazars are divided into two subclasses: flat-spectrum radio quasars and BL Lac objects, whose main difference is the lack or weakness of emission lines in BL Lac objects.  The commonly accepted paradigm capable of explaining most of the observed radio-loud AGN properties involves a supermassive black hole surrounded by an accretion disc, feeding a relativistic plasma jet emitting synchrotron and inverse-Compton radiation. In this scenario blazars would be the fraction of sources with the jet oriented at a small angle with respect to the line of sight.  To investigate the details of the physical processes and geometric conditions at the base of the extreme emission variability, a great observing effort is needed, possibly organizing multiwavelength campaigns to follow the source emission behaviour simultaneously at different frequencies. For example, the Whole Earth Blazar Telescope (WEBT; {\\tt http://www.to.astro.it/blazars/webt/}; e.g.\\ Villata et al.\\ \\cite{vil00}, \\cite{vil02}) is an international collaboration of optical and radio observers whose aim is the intensive and accurate monitoring of selected blazars during time-limited campaigns. The distribution in longitude of the WEBT members makes continuous (24 hours per day) optical monitoring possible, at least in principle. Since its inception in 1997, various BL Lac objects have been the target of the WEBT\\footnote{Targets of WEBT campaigns have been, besides BL Lacertae, \\object{AO 0235+16} (Raiteri et al.\\ \\cite{rai01}), \\object{Mkn 421}, \\object{S5 0716+71} (Villata et al.\\ \\cite{vil00}), \\object{Mkn 501}, \\object{3C 66A}.}; in particular, four campaigns have been devoted to \\object{BL Lacertae}, the prototype of the above mentioned subclass.  The first two BL Lac campaigns were organized in 1999 in conjunction with pointings of the X-ray satellites BeppoSAX (June 5--7) and ASCA (June 28--30). They lasted a few days only and involved a restricted participation of the WEBT members. The results of the BeppoSAX-WEBT campaign have been published in Ravasio et al.\\ (\\cite{rav02}): the shape of the X-ray spectrum was concave, with a very hard component above 5--6 keV; moreover, a very fast variability event was detected at the lowest X-ray energies, with no simultaneous optical counterpart.  The subsequent WEBT BL Lac 2000 campaign was an impressive observing effort with a long extension in time: more than 15000 optical observations were performed by 24 telescopes in 11 countries from May 2000 to January 2001. The  results of these observations are reported in Villata et al.\\ (\\cite{vil02}). The original motivation for the campaign was to provide the low-energy observing counterpart for a high-energy campaign organized in July--August 2000, involving X-ray satellites as well as TeV detectors (B\\\"ottcher et al.\\ \\cite{boe03}). In this period (the ``core campaign\") the densest sampling was achieved, with observing gaps limited to a few hours, mainly due to the lack of observers in the Pacific area.  After the BL Lac Campaign 2000, the WEBT collaboration launched a new BL Lac campaign in June 2001 (due to the detection of a fast brightening of the source), which lasted until the end of February 2002. In this paper we present the optical light curves of this BL Lac 2001 campaign and of its so-called ``extension\", which includes the two WEBT campaigns of 1999 and a collection of unpublished data back to 1994. We also include in our analysis previously published optical data from the BL Lac Campaign 2000 (Villata et al.\\ \\cite{vil02}) and literature data from 1994 to 2002. As a result, the optical light curves that we present in this work, based on new and previously published WEBT and literature data, cover a period of more than 8 years (9 observing seasons). Radio and optical light curves covering the period 1968--2003 will be presented in a forthcoming paper (Villata et al.\\ \\cite{vil04}), where timing and cross-correlation analyses will be performed on both optical and radio light curves.  In Sect.\\ 2 we report the optical observations of BL Lacertae carried out by the WEBT collaboration; the $UBVRI$ light curves from 1994 to 2002 are presented in Sect.\\ 3, and a detailed colour analysis is performed in Sect.\\ 4. Conclusions are drawn in Sect.\\ 5.  \\begin{table*} \\caption{List of participating observatories by longitude: $d$ is the telescope diameter; $N_{\\rm obs}$ is the total number of observations done, i.e.\\ the number of unbinned data; $N_U$, $N_B$, $N_V$, $N_R$, and $N_I$ are the numbers of data points in $UBVRI$ remained after discarding and binning some of the original data; the numbers in brackets refer to the WEBT campaign 2001 alone (May 2001 -- February 2002). } \\begin{tabular}{lccccccc} \\hline Observatory & $d$ (cm)& $N_{\\rm obs}$ & $N_U$ & $N_B$ & $N_V$ & $N_R$ & $N_I$ \\\\ \\hline Kyoto, Japan & 25 & 190 (0) & 0 & 0 & 0 & 19 (0) & 0 \\\\ Yunnan, China &100 & 25 (25) & 0 &2 (2)&2 (2) &7 (7) & 0 \\\\ Mt. Maidanak (T-60), Uzbekistan &60&50 (50) &15 (15) &16 (16) &16 (16) &0 &0 \\\\ Mt. Maidanak (AZT-14), Uzbekistan &50 &150 (150) &49 (49) &47 (47) &49 (49) &3 (3) &0\\\\ Mt. Maidanak (AZT-22), Uzbekistan &150 &418 (418) &50 (50) &87 (87) &65 (65) &108 (108) &61 (61)\\\\ Abastumani, Georgia (FSU) &70 &1823 (316) &0 &276 (36) &232 (0) &903 (213) &150 (0)\\\\ Crimean, Ukraine &70 &267 (267) &0 &67 (67) &67 (67) &67 (67) &66 (66) \\\\ Skinakas, Crete &130 &510 (239) &0 &186 (119) &69 (0) &184 (118) &67 (0)\\\\ Tuorla, Finland &103 &80 (0) &0 &0 &79 (0) &0 &0\\\\ MonteBoo, Czech Republic &62 &754 (754) &0 &0 &0 &282 (282) &0\\\\ Vallinfreda, Italy &50 &98 (0) &0 &0 &49 (0) &0 &0\\\\ Perugia, Italy &40 &524 (419) &0 &0 &112 (112) &208 (180) &166 (125)\\\\ Greve, Italy &32 &55 (0) &0 &0 &0 &0 &55 (0)\\\\ Torino, Italy &105 &487 (54) &0 &104 (1) &70 (2) &271 (24) &12 (2)\\\\ Sabadell, Spain &50 &32 (32) &0 &0 &3 (3) &12 (12) &0\\\\ Guadarrama, Spain &20 &3 (3) &0 &0 &1 (1) &1 (1) &0\\\\ Roque (KVA), La Palma &60 &322 (0) &0 &0 &0 &152 (0) &0\\\\ Roque (KVA), La Palma &35 &42 (42) &0 &0 &10 (10) &10 (10) &0\\\\ Roque (JKT), La Palma &100 &64 (64) &0 &9 (9) &10 (10) &32 (32) &6 (6)\\\\ Roque (NOT), La Palma &256 &44 (44) &4 (4) &7 (7) &7 (7) &7 (7) &3 (3)\\\\ Roque (TNG), La Palma &358 &138 (0) &3 (0) &62 (0) &5 (0) &62 (0) &2 (0)\\\\ Olin, Connecticut &51 &400 (361) &0 &0 &8 (0) &78 (64) &85 (71)\\\\ Hopkins, Massachusetts &60 &12 (0) &0 &0 &0 &0 &11 (0)\\\\ Foggy Bottom, New York &40 &248 (83) &0 &0 &0 &248 (83) &0\\\\ Boltwood, Canada &40 &96 (96) &0 &6 (6) &6 (6) &6 (6) &5 (5)\\\\ Bell, Kentucky &60 &23 (23) &0 &0 &0 &22 (22) &0\\\\ St. Louis, Missouri &36 &423 (340) &0 &10 (4) &40 (16) &161 (137) &39 (15)\\\\ Sommers-Bausch, Colorado &60 &9 (9) &0 &0 &4 (4) &4 (4) &0\\\\ Lowell, Arizona &180 &36 (0) &0 &0 &18 (0) &0 &18 (0)\\\\ San Pedro Martir, Mexico &210 &9 (9) &0 &3 (3) &3 (3) &3 (3) &0\\\\ San Pedro Martir, Mexico &150 &123 (123) &4 (4) &12 (12) &11 (11) &11 (11) &4 (4)\\\\ \\hline Total & &   7455 (3921) &125 (122) &894 (416) &936 (384) &2861 (1394) &750 (358)\\\\ \\hline \\end{tabular} \\end{table*}  \\begin{figure*} \\centering \\vspace{2cm} figure1.jpg \\vspace{2cm} \\caption{$UBVRI$ light curves of BL Lacertae from 1994 to 2002; the numbers on the left indicate the number of data points in the extended BL Lac Campaign 2001 (blue/dark), in the BL Lac Campaign 2000 (red/grey), and from the literature (light cyan/light grey); yellow (shaded) strips correspond to the periods of the 1999 and 2001 WEBT campaigns.} \\label{ubvri} \\end{figure*}  \\begin{figure*} \\centering \\vspace{2cm} figure2.jpg \\vspace{2cm} \\caption{Folded $R$-band light curve of BL Lacertae from 1997 to 2002; symbols as in Fig.\\ \\ref{ubvri}.} \\label{rtot_fold} \\end{figure*}  \\begin{figure*} \\centering \\includegraphics{figure3.eps} \\caption{Enlargements of the BL Lacertae $R$-band light curve corresponding to the WEBT campaigns of late June -- early July 1999 (top panel) and May 2001 -- February 2002 (middle and bottom panels); symbols as in Fig.\\ \\ref{ubvri}; in the top panel the yellow (grey) strip indicates the period of the ASCA pointing.} \\label{r_camp} \\end{figure*}  \\begin{figure*} \\centering \\includegraphics[width=13cm]{figure4.eps} \\caption{\\ An enlargement of the $R$-band light curve of BL Lacertae on June 29, 1999, corresponding to the green (grey) box in Fig.\\ \\ref{r_camp}; the yellow (grey) filled circles represent the simultaneous $B$-band light curve shifted by $-1.62$ mag.} \\label{rtot_pap2} \\end{figure*}  \\begin{figure*} \\centering \\vspace{2cm} figure5.jpg \\vspace{2cm} \\caption{$B-R$ (top panel) and $V-I$ (middle panel) colour indices versus time from 1997 to 2002 compared to the $R$-band light curve; the host-galaxy contribution has been removed from the data; symbols as in Fig.\\ \\ref{ubvri}; the horizontal lines divide each panel into four equal ranges (1 mag for the light curve and 0.25 mag for the colour indices) for clarity.} \\label{colori_jd} \\end{figure*}  \\begin{figure*} \\centering \\includegraphics{figure6.eps} \\caption{An enlargement of the $B-R$ versus time curve (top panel) compared with the $R$-band light curve (bottom panel) during the best-sampled period of the BL Lac Campaign 2001; orange (grey) dots represent the data used for calculating the colour indices.} \\label{colori_jd2} \\end{figure*}  \\begin{figure*} \\centering \\includegraphics[width=13cm]{figure7.eps} \\caption{\\ $B-R$ colour indices versus $R$ magnitudes after the host-galaxy removal (upper panel): BL Lac 2000 and literature data are represented by red (grey) and light cyan (light grey) dots, respectively, while BL Lac 2001 extended data are plotted with green (grey) circles (BL Lac 2001 extension), blue (dark) squares (best-sampled period of BL Lac 2001, the same of Fig.\\ \\ref{colori_jd2}), and cyan (grey) crosses (remaining part of BL Lac 2001); most of the data lie within a parallelogram whose sides have slopes of 0.1 and 0.4. In the bottom panel the same data have been corrected for the long-term trend as explained in the text.} \\label{b-r_vs_r} \\end{figure*}  \\begin{figure*} \\centering \\vspace{2cm} figure8.jpg \\vspace{2cm} \\caption{Galaxy-subtracted $R$-band flux light curve from 1997 to 2002 (folded, symbols as in Fig.\\ \\ref{ubvri}); the green (grey) line represents a cubic spline interpolation through the 4-day binned light curve; the blue (dark) line, obtained by normalizing the spline to its minimum value, represents the actual correction factor for removing the long-term trend.} \\label{fitta_fold} \\end{figure*} ",
        "conclusions": "We have presented $UBVRI$ light curves of BL Lacertae from 1994 to 2002 composed from data taken by the members of the WEBT collaboration as well as literature data.  The colour analysis we performed on the best-sampled period of these light curves (1997--2002) leads to the following results.  The variability observed in the optical light curves can be interpreted in terms of two components: a ``mildly-chromatic\" longer-term component and a ``strongly-chromatic\" shorter-term one. The longer-term component has a typical time scale of a few days, while the shorter-term one is responsible for the very fast, intraday flares, and determines a steep bluer-when-brighter slope of about 0.4 in the $B-R$ versus $R$ plot. Moreover, our analysis suggests that, in the long-term, the $B$-band flux level is related to the $R$-band one according to a power law with index 1.1. This means that long-term variations trace a 0.1 slope in the $B-R$ versus $R$ plot. In other words, all possible flux variations fill a parallelogram with 0.1 and 0.4 slopes in the $B-R$ versus $R$ plane.  The analysis presented in this paper shows that, in general, the BL Lac optical behaviour can be uniformly described in terms of these two components fairly well, with the caution that the presented decoupling method can work well only during well-sampled periods, where the light-curve binning can actually represent the long-term trend. However, peculiar events also occurred, as in the BL Lac 2000 period, and our scenario seems to be more a good compromise between slightly different behaviours rather than a definitive solution. Other factors appear to change in the source history, in one or both of the components.  Changes in the flaring component could explain the larger amplitudes of the 1997 outburst and the bluer BL Lac 2000 period, but the remaining blue residues of the 1997, 2001, and 2002 outbursts suggest other possibilities.  In Villata et al.\\ (\\cite{vil02}) the quasi-achromatism of the long-term variations was interpreted in terms of Doppler factor variations. If the intrinsic source spectrum around the infrared--optical bands is well described by a power law ($F_\\nu \\propto \\nu^{-\\alpha}$), a Doppler factor variation does not imply a colour change. What about a mildly-chromatic behaviour? It could be due to a Doppler factor variation on a spectrum slightly deviating from a power law (convex for a bluer-when-brighter trend). If the spectrum shape changes in time, the chromatism of the outbursts also varies, and the 0.1 $B-R$ slope adopted in this paper might represent only a mean value among different slopes. In other words, any period should be treated with a proper slope. A further investigation of this is beyond the scope of this paper.  The strongly-chromatic fast flares are likely due to intrinsic phenomena, such as particle acceleration from shock-in-jet events, widely described in the literature (see e.g.\\ Mastichiadis \\& Kirk \\cite{mas02} for a review). Other models are also discussed in the recent paper by B\\\"ottcher \\& Reimer (\\cite{boe04}).  Thus, the present extension/enrichment of the BL Lac optical dataset partially confirms previous results leading to a more refined scenario, but also gives rise to new intriguing questions."
    },
    "0404/astro-ph0404363_arXiv.txt": {
        "abstract": "The exact nature of $\\eta$ Carinae is still an open issue. Strict periodicity in the % light curves at several wavelengths seem to point out to a binary system, but the % observed radial velocities, measured from space with high spatial resolution are in % conflict with the ground based observations used to calculate the binary orbit. % Also, the observed 2-10 keV X-ray flux is much larger that what is expected from a % single star, and favors the wind-wind collision hypothesis, characteristic of high % mass binary systems. However, to explain the duration of the dip in the light curve % by wind collisions, it is necessary to postulate a very large increase in the $\\eta$ % Carinae mass loss rate. Finally,  the optical and UV light curves are better % explained by periodic shell-ejection events. In this paper we conciliate the two % hypothesis. We still assume a binary system to explain the strong X-ray emission, % but we also take into account that, near periastron and because of the highly % eccentric orbit, the wind emerging from $\\eta$ Carinae accumulates behind the shock % and can mimic a shell-like ejection event. For this process to be effective, at % periastron the secondary star should be located between $\\eta$ Carinae and the % observer, solving also the discrepancy between the orbital parameters derived from % ground and space based observations. We show that, as the secondary moves in its % orbit, the shell cools down and the number of available stellar ionizing photons is % not enough to maintain the shell temperature at its equilibrium value of about 7500 % K. The central part of the shell remains cold and under these conditions grain % formation and growth can take place in timescales of hours. We also calculated the % neutral gas column density intercepting the line of sight at each point of the orbit % near periastron, and were able to reproduce the form and duration of the X-ray light % curve without any change in the $\\eta$ Carinae mass loss rate. This same column % density can explain the observed H$\\alpha$ light curve observed during the 2003 % event. ",
        "introduction": "$\\eta$ Carinae is a peculiar star surrounded by a dense cloud of gas and dust, % formed in several episodes of mass ejection, typical of  LBV stars. Its mass, % calculated from the bolometric luminosity, corresponds to approximately 120 % M$_{\\odot}$, making $\\eta$ Carinae, if single, the most massive star in the Galaxy. During the last years strong evidences of the system binary nature  were reported: a % 5.52 year periodicity in the light curve at infrared wavelengths % \\citep{whitelock94}, in the HeI $\\lambda$10830 \\AA$\\;$  high excitation line % \\citep{dam96} and in the radial velocity of the Pa$\\gamma$ line, the latter pointing % out to a highly eccentric orbit \\citep*{damineli97, davidson97}. The 5.52-year periodicity was also found at centimeter and millimeter wavelengths % \\citep{coxa95,abra99,duncan03,abr03}. VLBI observations showed a very compact emitting region, coincident with the star % position at the minimum in the light curve, which grew to an extended and elongated % feature at maximum \\citep*{duncan97}, indicating the probable existence of an % ionized disk of variable size. Infrared observations also showed a compact disk, with a power law spectrum between % 2 and 8 $\\mu$m, attributed to thermal emission by dust at temperatures between 1500 % and 400 K \\citep{morris99}.  A draw back in the binary system hypothesis came from the high resolution $HST$ % spectroscopic observations of $\\eta$ Carinae \\citep{davidson00}, which showed that, % although the radial velocity of the emission lines formed close to the star varied % at the different phases of the binary orbit, they were incompatible with the orbital % parameters derived by \\citet{damineli97}. Without the assumption of a binary companion, the spectroscopic events could be % explained if they were the result of periodic shell ejections % \\citep*{zanella84,davidson99, martin04}. Under these conditions, the dense gas in the shell would screen the stellar UV % radiation, responsible for the formation of the ionized envelope surrounding $\\eta$ % Carinae, favoring also the formation of dust, which could explain the enhanced IR % emission.  On the other hand, a strong argument in favor of the binary hypothesis is based in  % the intensity of the observed X-ray emission, which cannot be explained by shell % ejection but can be easily accounted for by free-free emission from wind-wind % collisions \\citep{pittard98}. In fact,  this process was already  postulated in % other massive binary systems to explain the differences between the observed X-ray % emission ($\\sim 10^{35}$ erg s$^{-1}$) and what is expected from individual stars % ($\\sim 10^{32}-10^{33}$ erg s$^{-1}$) \\citep{prilutskii76, pollock87, zhekov00}.   The 1997 and 2003 events, characterized by  minima in the light curves at other % wavelengths, were also observed at X-rays, preceded by a strong increase in the flux % density. The enhanced  emission was  attributed to an increase in the wind-wind collision % strength as the stars approach each other during  periastron passage; the sharp % decrease towards the minima was assumed to be produced by an increase in the H % column density, perhaps as the result of an increase in the $\\eta$ Carinae mass loss % rate  \\citep{pittard98, corcoran01, pittard02}.  In this paper we conciliate the two hypothesis: a binary system and a shell % ejection-like event. We use the binary system to explain the X-ray emission by shock formation in the % wind-wind collision. What we call the shell ejection-like event is the appearence of  dense material % between the observer and the binary system, formed by the stellar wind matter % accumulated behind the cool expanding  shock remnant. For this process to be effective in absorbing the X-rays, the periastron must be in % the oposite direction of what was proposed by \\citet{damineli00} and % \\citet{corcoran01}, being compatible with the $HST$ observations reported by % \\citet{davidson00}. We show that in this event dust can be formed in the shock interface, where the gas % cools down by radiation and it is shielded from the UV stellar radiation by the % large amount of accumulated gas. The dense gas and dust shell will eventually expand and form a disk in the plane of % the orbit, its internal parts will be ionized by the UV radiation of the two stars % in the binary system and will be seen at radio wavelengths as free-free emission. The disk will also slow down the wind of $\\eta$ Carinae in the plane of the orbit, % which could contribute with other effects (e.g. non symmetrical wind from the % massive rotating primary star) to the  latitude-dependent behavior observed by % \\citet{smith03}.  In Section 2 of this paper we calculate the physical conditions in the shocked % region, taking into account radiative absorption and emission. In Section 3 we calculate the grain formation rate and in Section 4 we compare our % results with the optical and UV observations  of the 2003 low excitation event % reported by  \\citet{martin04}. Finnaly  in Section 5 we present the derived X-ray % light curves for different values of the orbital parameters and of the $\\eta$ % Carinae mass loss rate and in Section 6 we present our conclusions. ",
        "conclusions": "In this work we showed that the wind-wind collision model of a binary system in a % highly eccentric orbit can explain the periodic events in the 2-10 keV X-ray light % curve of $\\eta$ Carinae, as well as the shell-like events at optical and UV % wavelengths, without any increase in  the primary mass loss rate, as postulated by % \\citet{corcoran01}.  As in previous works \\citep{pittard98, corcoran01,pittard02}, we calculated the % physical conditions in the shock region close to periastron passage, for an % eccentricity $e = 0.9$,  mass loss rates of  $\\dot M_p = 2.5\\times 10^{-4}$ % M$_{\\odot}\\;$ yr$^{-1}$ and $\\dot M_s =  10^{-5}$ M$_{\\odot}\\;$ yr$^{-1}$, and wind % velocities $v_p = 700$ km s$^{-1}$ and $v_s =3000$ km s$^{-1}$. We  obtained for the % X-ray emitting region a density  $n_{2s}\\sim 5.6 \\times 10^{9}$ cm$^{-3}$ and % temperature $T\\sim 2\\times 10^8$ K, in agreement with recent $Chandra$ observations % and numerical wind collision models \\citep{pittard02}.  The difference with the cited papers is that in our work we took into account that % the wind from $\\eta$ Carinae accumulates behind the shock, forming a shell about % $10^{12}$ cm thick. We then showed that, as the secondary moves in its orbit, the shell cools  due to % free-free and line emission. Also, the number of ionizing photons from $\\eta$ % Carinae and its companion star are only enough to maintain the gas temperature at % its equilibrium value of about 7500 K in the external parts of the shell (10$^9$ cm % in depth), leaving the central part neutral and cold. We calculated that in these % conditions dust grains can form and grow in timescales of a few hours. These grains % would absorb the optical and UV radiation emitted by the shock material as it cools % down. We used the recent $HST$ observations of the 2003 low excitation event, % presented by \\citet{martin04}, and showed that the decrease in the continuum flux at % wavelengths of 2200, 2500, 3300 and 5500 \\AA$\\,$ can very well be explained in this % scenario. On the other hand, the decrease in the observed H$\\alpha$ emission, without any % similar effect in the adjacent continuum, was attributed by us to absorption by the % neutral and cool H gas accumulated in the central region of the shell. From the % amount of absorbed radiation we were able to calculate the number of hydrogen atoms % in the second excited energy level ($\\sim 10^{10}$ cm$^{-2}$). This is compatible % with the inferred H column densities and temperatures of the order of 1000 K % expected in the neutral region.   We also calculated the total amount of gas that intercepts the line of sight for % each point in the orbit near periastron passage, assuming it is close to % conjunction, and derived a profile similar to that found phenomenologically by % \\citet{corcoran01}, but without the need of an increase in mass loss rate from % $\\eta$ Carinae.  We constructed the expected 2-10 keV light curves for several values of the % eccentricity, phase of opposition, mass loss rates  and wind velocities for $\\eta$ % Carinae and we were able to fit very well the shape and duration of the emission % dip. We found that the mass  loss rate is the only critical parameter that affects % the shape of the X-ray light curve, since for high eccentricities $\\Phi$ angles lower than 30 degrees represent less than 1 day in the light curve timescale. The best fitting was obtained for $\\dot M_p = % 2.5\\times 10^{-4}$ M$_{\\odot}\\;$ yr$^{-1}$, which coincides with the value obtained % by \\citet{pittard02} in their numerical simulations.  An interesting result of our model is that it explains the difference in the % duration of the dip in the light curves at different wavelengths. At optical and UV % wavelengths the dip is due to absorption by dust. Since dust growth rates are very % sensitive to gas density, its formation rate, as well as the absorption it causes, % will decrease very fast as the distance between the stars increases after periastron % passage. X-rays, on the contrary, are absorbed by neutral gas, which will accumulate % across the line of sight when the gas shell formed behind the shock expands away % from the stars, and will remain there for a long time after grain formation had % stopped.  In the scenario presented here, the material accumulated in the shell will expand % into the orbital plane and form a thick disk,  which can be seen at millimeter % wavelengths through its free-free emission \\citep{coxa95,abra99,duncan03}. The dip % in the millimeter light curves can be explained by a decrease in the number of % available ionizing photons, which are absorbed by the neutral shell close to the % binary orbit.  Finally we must mention that what we presented here is a very simplified version of % a very complicated process and any refinement in the derived orbital and physical % parameters of the binary system should be obtained from detailed numerical % simulations."
    },
    "0404/astro-ph0404013_arXiv.txt": {
        "abstract": "In modeling strong gravitational lens systems, one often adopts simple models, such as singular isothermal elliptical density plus lowest-order external perturbation. However, such simple models may mislead us if the real mass distribution is more complicated than that in the assumed models. In particular, assumptions on mass models are crucial in studying flux ratio anomalies that have been suggested as evidence for a cold dark-matter substructure. We reinvestigated four quadruple lens systems using power-law Fourier models, which have advantages of clear physical meanings and the applicability of a linear method, as well as a simple singular isothermal elliptical density model. We also investigated the effect of external perturbations, including a singular isothermal sphere, lowest-order expansion, and next-order expansion. We have found that the $\\cos3\\theta$ terms of the primary galaxy and/or of external perturbation significantly reduce the $\\chi^2$ in PG 1115+080 and B1422+231. In particular, we could reproduce the flux ratios of B1422+231 with next-order external perturbation assuming 5\\% flux uncertainties, suggesting that external perturbation cannot be described by a simple singular isothermal sphere approximation. On the other hand, we could not fit B0712+472 and B2045+265 very well even with our models, although the $\\chi^2$ were reduced compared with the case of using simple models. Our results clearly demonstrate that both the primary lens galaxy and the external perturbation are often more complicated than we usually assume. ",
        "introduction": "Gravitational lenses of distant quasars are quite useful in cosmological and astrophysical studies, such as measuring the Hubble constant from time delays \\citep{refsdal64}, studies of galaxy formation and evolution, and direct probes of galaxy structure, including cold dark-matter substructures. The method used to measure the Hubble constant does not rely on the distance ladder, which inevitably suffers various systematic errors. After the discovery of the first lens system, Q 0957+561 \\citep{walsh79}, many lens systems have been studied in detail, adopting various lens models. However, the model uncertainties, including several intrinsic degeneracies \\citep{falco85,gorenstein88,saha00}, prevent us from accurately deriving the Hubble constant (e.g., \\cite{wambsganss94,keetonks97,keetonk97,courbin97,wucknitz99,witt00,chiba02,wucknitz02,oguri02,ogukawa03}).  Among several applications of gravitational lenses, flux-ratio anomalies have attracted much attention. There are mainly three explanations for them: Effects of the interstellar medium, problems in smooth models, and lensing by the substructure (e.g., \\cite{kocda03}). While substructure lensing is often regarded as being the best explanation for the flux anomalies (e.g., \\cite{mao98,chiba02b}), the former two explanations have not yet been sufficiently investigated to reach the conclusion that the flux-ratio anomalies are connected with the substructure. The limitation comes from the fact that it is difficult to apply all of the models, which are physical and complicated, to the observed lens systems. In addition, since galaxies are believed to consist of not only luminous matter, but also dark matter, the shapes of lens galaxies do not have to be elliptical, nor other simple shapes, as the isophotes indicate. Therefore, it is of great importance to investigate the contribution of many forms of the angular structure of the lens potential to find the true origin of the flux anomalies. Indeed, \\citet{evansw03} showed that some of lens systems that have been claimed to be lenses with flux anomalies  can be well fitted by a smooth lens model using a power-law Fourier potential (see also \\cite{ogukawa03}). Kochanek and Dalal (2003) studied the contribution of only the isothermal $m=3$ and $m=4$ multipoles to the flux ratio anomalies reported so far. We then reinvestigate some of the systems with not only multipoles, but also several types of external perturbations.  Our aim in this paper is to discuss the contribution of both multipoles of the primary lens galaxy and external perturbation of the group, in particular the $\\cos3\\theta$ terms that are not usually included in lens modeling. We consider whether some multipoles can produce an acceptable fit without substructure or not. We also discuss the quadruple lens system B0712+472 in detail because of the interesting observations that the ellipticity of the galaxy and the shear of the group seem to be significantly misaligned ($\\sim \\timeform{90D}$) to each other, and that previous authors estimated the shear to be small compared to observations using the one-shear model \\citep{jackson98,fass02}. Since more than three images are necessary to constrain the mass models strongly, we restrict our considerations to four quadruple lens systems (PG 1115+080, B1422+231, B0712+472, and B2045+265). There are three reasons for this restriction. First, the existence of a substructure of the lens systems has been claimed to account for the observed flux ratios (e.g., \\cite{kocda03}). Second, there are no satellite galaxies that complicate the lens model. Third, there are sufficient astrometric, photometric, and spectroscopic (redshift) data. In this paper, we consider only isothermal models; however our analysis can be easily extended to non-isothermal radial mass profiles.  This paper is organized as follows. In section \\ref{sec:modelmethod}, we describe the models that we used to investigate four systems. In section \\ref{sec:results}, we show how we applied the  models to four systems and give the results. We finally discuss the results and give conclusions in section \\ref{sec:conc}. The linear method that we used is presented in Appendix \\ref{app:linear}. ",
        "conclusions": "} We investigated four quadruple gravitational lens system (PG 1115+080, B1422+231, B0712+472, and B2045+265) using mass models of SIED or PLF as well as external perturbations described by SIS, or several truncations of the expansion of the external perturbation. We performed $\\chi^2$ fittings with all of the models presented in section \\ref{sec:modelmethod} in order to study the contribution of terms that are usually not included in mass modelings, such as $m=3$ multipoles of primary lens galaxy and next-order external perturbations.  We have found that the $m=3$ multipoles and/or X3 external perturbation significantly reduce the $\\chi^2$ in PG 1115+080 and B1422+231. In PG 1115+080, not only the external perturbation of the $r^3$ order (X3), but also an internal $m=3$ multipole, reduces the $\\chi^2$ to statistically acceptable values, but the ellipticity of the galaxy does not. Their amplitudes and directions seem to be physical, comparing the results with optical and X-ray observations of the galaxy and the group \\citep{impey98,grant03}; in a ($m=3$)+XS model the ratio $A_3/a_0\\approx 0.002$ is small enough to reproduce the observed very spherical galaxy, and in a SIS+(XS+X3) model the ratio $\\sigma/\\gamma \\approx 0.1$ and the angles $\\theta_{\\gamma}=\\timeform{-115D}$ and $\\theta_{\\sigma}=\\timeform{-160D}$ can be expected from the positions of the member galaxies and the X-ray brightness contours. A SIS group also provides an acceptable fit. A significant (but small) $m=3$ galaxy multipole might be induced by interactions with the member galaxies. Regarding B1422+231, several authors concluded that the ``anomalous'' flux ratios reflected the existence of substructure lensing (e.g., \\cite{mao98,chiba02b,bradac02,kocda03}). However, Keeton, Gaudi, and Petters (2003) found that the cusp relation does not reveal a significant anomaly in the system. In fact, our analysis has revealed that a SIED+(XS+X3) model without substructure reproduces the fluxes within 5\\% errors. In addition, the X-ray brightness contours (figure 2 of  \\cite{grant03}) imply that the direction of the gradient of the group mass is $\\approx \\timeform{90D}-\\timeform{180D}$, and that the shapes of the contours are significantly deformed, which are consistent with the results of SIED+(XS+X3) model. In the other systems, B0712+472 and B2045+265, the contribution of the $m=3$ multipoles and X3 external perturbations to the lens potentials do not significantly improve the model fitting. In B0712+472, however, we find that the ellipticity and shear are well misaligned, $\\theta_{\\mathrm{e}}\\approx \\timeform{56D}$ and $\\theta_{\\gamma}\\approx \\timeform{146D}$, and that the result is very consistent with the lens galaxy and the foreground group observed by Fassnacht and Lubin (2002). Therefore, the models including only ellipticity or shear may underestimate these (see subsection \\ref{subsec:impl}). The relatively large shear may arise from the elongated shape of the group.  In summary, the terms of an $m=3$ multipole of the primary lens galaxy and of an X3 external perturbation, which are often neglected in strong lens analyses, sometimes improve the fit significantly. The contribution of the group particularly needs to be studied in more detail. Simple models (in particular one-shear models) may often mislead us.   \\bigskip This work is supported by a Grant-in-Aid for JSPS Fellows.    \\appendix"
    },
    "0404/astro-ph0404539_arXiv.txt": {
        "abstract": "We have reanalyzed in a consistent way existing reverberation data for 35 AGNs for the purpose of refining the black hole masses derived from these data. We find that the precision (or random component of the error) of reverberation-based black hole mass measurements is typically around 30\\%, comparable to the precision attained in measurement of black hole masses in quiescent galaxies by gas or stellar dynamical methods. As discussed in this volume by Onken et al., we have established an absolute calibration for AGN reverberation-based masses by assuming that AGNs and quiescent galaxies follow an identical relationship between black hole mass and host-galaxy bulge velocity dispersion. The scatter around this relationship implies that the typical systematic uncertainties in reverberation-based black hole masses are smaller than a factor of three. We present a preliminary version of a mass--luminosity relationship that is much better defined than any previous attempt. Scatter about the mass--luminosity relationship for these AGNs appears to be real and could be correlated with either Eddington ratio or source inclination. ",
        "introduction": "While black holes have been invoked since the early days of quasar research as the principal ingredient in AGNs, it is only within the last several years that it has become possible to measure the masses of the central objects in galactic nuclei. At the present time, there has not yet been a definitive detection of the relativistic effects that would be required for unambiguous identification of a singularity, although studies of the Fe K$\\alpha$ emission line in the X-ray spectra of AGNs currently afford some promise (e.g., Reynolds \\& Nowak 2002). However, it seems to be true that the centers of both active and quiescent galaxies host supermassive objects that must be so compact that other alternatives are very unlikely.  Black hole masses are measured in a number of ways. In Type 1 active galaxies, reverberation mapping (hereafter RM; Blandford \\& McKee 1982; Peterson 1993) of the broad-line region (BLR) can be used to determine the central masses. RM is the only method of black hole mass measurement that does not depend on high angular resolution, so it is of special interest as it is extendable in principle to both very high and very low luminosities and to objects at great distances. Moreover, RM studies reveal the existence of simple scaling relationships that can be used to anchor secondary methods of mass measurement, thus making it possible to provide estimates of the masses of large samples of quasars, including even very distant quasars, based on  relatively simple spectral measurements (e.g., Vestergaard 2002, 2004; McLure \\& Jarvis 2002).  The evidence that RM-based black hole masses are valid is twofold: \\begin{enumerate} \\item There is an anticorrelation between emission-line time lag $\\tau$, which measures the size of the line-emitting region by light-travel time, and emission-line width $\\Delta V$ that is consistent with the virial prediction $\\tau \\propto \\Delta V^{-2}$. Relative to lower ionization lines, higher ionization lines are both broader and closer to the central source, consistent with ionization stratification of the BLR and dynamics that are dominated by a central mass (Peterson \\& Wandel 1999, 2000; Onken \\& Peterson 2002; Kollatschny 2003). \\item There is a relationship (Gebhardt et al. 2000b; Ferrarese et al.\\ 2001) between the RM-based black hole masses and the host galaxy stellar velocity dispersion $\\sigma_*$ similar to that seen in quiescent galaxies (Ferrarese \\& Merritt 2000; Gebhardt et al.\\ 2000a). \\end{enumerate}  A difficulty with the existing RM database is the lack of uniformity in how the measurements have been made. This leads to considerable scatter in many of the relationships among various AGN properties. We have therefore undertaken a complete reanalysis of existing RM data with two goals in mind: \\begin{enumerate} \\item In order to improve the {\\em precision} (i.e., reduce the random errors) of AGN black hole mass measurements, we reanalyzed all of the readily available RM data to determine the best measures of time lag and line width for these studies. \\item In order to improve the {\\em accuracy} (i.e., reduce the systematic errors) of AGN black hole mass measurements, we obtained high-precision measurements of $\\sigma_*$ to determine the $M_{\\rm BH}-\\sigma_*$ relationship for AGNs. This is the subject of a companion paper (Onken et al., this volume), which we will draw on for the absolute calibration of the RM-based black hole mass scale. \\end{enumerate} The results of this program are described in detail by Peterson et al.\\ (2004), Onken et al.\\ (2004), and Kaspi et al.\\ (2004) and are intended to supersede previous compliations of AGN black hole masses by Wandel, Peterson, \\& Malkan (1999) and Kaspi et al.\\ (2000). ",
        "conclusions": ""
    },
    "0404/astro-ph0404069_arXiv.txt": {
        "abstract": "Through time-series CCD photometry of the globular cluster NGC 5466, we have detected nine SX Phoenicis stars including three new ones. All the SX Phoenicis stars are located in the blue straggler region in the color-magnitude diagram of NGC 5466. Five of them show clearly double-radial mode features, the periods of  which are well consistent with the theoretical ratio of the first overtone mode to the fundamental mode (P$_{1H}$/P$_F$). Normally, it has not been easy to secure a P-L relation of the SX Phoenicis stars, because determination of the pulsational mode of the SX Phoenicis stars has been difficult. The existence of five SX Phoenicis stars in NGC5466 with double-radial modes allows us to derive reliably a P-L relation for the fundamental mode of  the SX Phoenicis stars. Using seven SX Phoenicis stars including five stars with double-radial modes, we derive a P-L relation for the fundamental mode in NGC 5466, $<V> = - 3.25 (\\pm 0.46) Log P + 14.70 (\\pm 0.06 ),~ (\\sigma = \\pm0.04) $, corresponding to $<M_V> = - 3.25  (\\pm0.46)Log P - 1.30 (\\pm0.06 ) $ for an adopted distance modulus of $(m-M)_0=16.00$ and zero reddening. ",
        "introduction": "NGC 5466 (R.A.\\,=\\,14$^h$ 05$^m$ 27$\\fs$3, Decl.\\,=\\,+28$\\arcdeg$ 32$\\arcmin$ 04$\\arcsec$, J2000.0) is a sparse globular cluster with an extremely low metallicity [Fe/H]=--2.22 \\citep{har96}, having a large number of blue stragglers. \\citet{nem87} found 48 blue straggler stars (BSSs) which are mostly concentrated within $r=3'$ from the center of the cluster. Several kinds of variables stars are known to exist in this cluster: six SX Phoenicis stars \\citep{nem90,che99}, three eclipsing binaries \\citep{mat90,kal92,mck98}, a Cepheid and more than twenty RR Lyrae stars \\citep{cle01}.  In this paper we present a photometric study of SX Phoenicis stars in NGC 5466. In the long-term search for the short-period variable stars in Galactic globular clusters, we have found three new SX Phoenicis stars named Cl* NGC 5466 SXP 1, Cl* NGC 5466 SXP 2 and Cl* NGC 5466 SXP 3 in addition to the previously known six named Cl* NGC 5466 NH 27, Cl* NGC 5466 NH 29 Cl* NGC 5466 NH 35, Cl* NGC 5466 NH 38, Cl* NGC 5466 NH 39 and Cl* NGC 5466 NH 49 by \\citet{nem90}; % hereafter referred to as SXP1, SXP2 and SXP3 for new ones, and NH27, NH29, NH35, NH38, NH39 and NH49 for known ones, respectively. Results on the eclipsing binaries, a Cepheid and RR Lyrae stars in NGC5466 will be presented in a separate paper \\citep{lee04}. We adopt the zero interstellar reddening and the distance modulus $(m-M)_0=16.00$ in this study \\citep{har96}.  This paper is composed as follows. Observations and data reduction are described in $\\S$ 2. $\\S$ 3 describes selection of variable stars in NGC 5466, and $\\S$  4 presents the light curves and frequency analysis of the SX Phoenicis stars. $\\S$  5 discusses the characteristics of these SX Phoenicis stars, including mode identification and the P-L relation. Finally primary results are summarized in $\\S$ 6. ",
        "conclusions": "\\subsection{ Characteristics of the SX Phoenicis Stars }  In Figure 7, we show the position of the nine SX Phoenicis stars in the color-magnitude diagram of NGC 5466, listing their mean magnitudes and color indices in Table 2. Figure 7 shows that all the  SX Phoenicis stars are located in the blue straggler region, brighter and bluer than the main sequence turnoff point. It is known that all the known SX Phoenicis stars in globular clusters are BSSs, but there are non-variable stars among the BSSs. \\placefigure{fig7}  In Figure 8 we have compared the $V$ amplitudes and periods of the SX Phoenicis stars in NGC 5466 (filled circles) with those of SX Phoenicis stars in other globular clusters, field SX Phoenicis stars and $\\delta$ Scuti stars. \\placefigure{fig8} It includes  the SX Phoenicis stars in M53 (star symbols) and M15 (a cross) discovered by our previous searches for variable blue stragglers in globular clusters \\citep{jeo01,jeo03}. The sources of the data in Figure 8 are \\citet{rol20}  for field SX Phoenicis stars and  $\\delta$ Scuti stars, and \\citet{rod20}  for SX Phoenicis stars in  Galactic globular clusters. Figure 8 shows that the $V$ amplitudes and periods of  the SX Phoenicis stars in NGC 5466 are consistent with those for SX Phoenicis stars in  other globular clusters; the $V$ amplitudes increases steeply with increasing period, and the $V$ amplitudes of SX Phoenicis stars are larger than those of  $\\delta$ Scuti stars with the same period.  \\subsection{Radial Mode Identification}  A method for mode identification  is to use a pulsation constant, $Q$, denoted as $Q = P \\rho^{1/2}$ where $P$ and $\\rho$ are the period and mean density of variable stars, respectively. $Q$ values of overtone mode are smaller than that of the fundamental mode \\citep{bre00}. $Q$ values can be derived using a photometric method such as Str\\\"{o}mgrem {\\it uvby H$_{\\beta}$} observation \\citep{rod03}. Another photometric method is to use SX Phoenicis stars with double-radial mode, but the SX Phoenicis stars with double-radial mode are rare. Fortunately, five (including two candidates) of the SX Phoenicis stars in NGC5466 are found to have double-radial modes, so their modes are identified clearly. All the double-radial mode stars (SXP2, SXP3, NH35, NH38 and NH39) show the fundamental and first overtone modes, as described in $\\S$ 3.1.  The mode of the rest of the SX Phoenicis stars can be identified using the amplitude, period and luminosity. NH27 is identified as the first overtone mode because of  its $V$ amplitude, 0.246 mag, being slightly lower, and especially the position in the period-magnitude diagram as shown in Figure 9. SXP1 has a high amplitude, and is located in the fundamental region in the period-magnitude diagram. So it % is a fundamental mode star. In Table 3, NH35, NH38, NH39 and NH49 show harmonic frequencies supporting asymmetric sinusoidal features. Moreover their total light amplitudes have a range of 0.282 $\\sim$ 0.730 mag, which are characteristics of high amplitude $\\delta$ Scuti (HADS) stars. So, their mode are clearly identified to be fundamental radial oscillations. Their primary and secondary radial periods correspond to the fundamental and first overtone radial modes, respectively. NH29 % is the first overtone mode star according to its low amplitude (0.142 mag), well defined sinusoidal feature in the light curve (see Figure 3) and the location in the period-magnitude diagram. Seven of nine SX Phoenicis stars in NGC 5466 show the fundamental mode or double-radial mode with the fundamental and first overtone modes. They  are displayed by open and filled circles in the color-magnitude diagram in Figure 7, respectively. Two filled triangles denote the first overtone mode stars. The identified radial modes are listed in the last column of Table 3.  \\citet{mcn20,mcn01} suggested that the light amplitude and the degree of asymmetry of the light curves % are useful parameters for identifying the pulsating modes. Generally this method is very useful for systems with a small number of pulsators and/or no double-radial mode stars therein. However, in many cases these criteria do not give a unique solution to identify pulsating modes, especially for SX Phoenicis stars with very low amplitudes and short periods such as SXP2 and SXP3 in NGC5466.  \\subsection{ Period-Luminosity Relation }  The P-L relation of SX Phoenicis stars in the globular clusters is very useful to obtain the distance moduli of the clusters and nearby galaxies \\citep{mcn95}. However, it is not easy to define well the P-L relation from the observations, because there are often a mixture of different pulsation modes \\citep{jeo03,mcn01}. But we can easily identify the pulsating mode using the double-radial mode stars as described in the previous section.  Figure 9 presents the P-L relation for the SX Phoenicis stars in NGC 5466. \\placefigure{fig9} It shows that the sample can be separated into two discrete groups; one is identified as a fundamental mode group (filled circles), and the other is a first overtone % mode group (filled triangles and open triangles). The solid line represents the fundamental P-L relation, and the line is shifted to the dashed line corresponding to the first overtone mode stars by the ratio $P_{1H}/P_F = 0.783 $. The P-L relation for fundamental mode in NGC 5466 is derived to be \\begin{equation} <V> = - 3.25 (\\pm 0.46) Log P + 14.70 (\\pm 0.06 ),~ (\\sigma = \\pm0.04), \\end{equation} which corresponds to \\begin{equation} <M_V> = - 3.25  (\\pm0.46)Log P - 1.30 (\\pm0.06 ) \\end{equation} for an adopted distance modulus of $(m-M)_V=16.00$ and $E(B-V)=0.00$ \\citep{har96},  The empirical P-L relations have been obtained using field HADS stars and/or cluster SX Phoenicis stars identified by their pulsating modes. The slopes derived by field HADS stars and SX Phoenicis stars in $\\omega$ Cen show on the whole steeper ones up to --4.66 \\citep{mcn20}, whereas $\\delta$ Scuti stars, SX Phoenicis stars in other globular clusters and theoretical results show flatter slopes up to --2.88 \\citep{pyc01}.  The slope of $-3.25$ for NGC 5466 is in a good agreement with the results for M53 \\citep{jeo03} and M55  \\citep{pyc01}. \\citet{jeo03} obtained a slope of $-3.01$ for the fundamental mode stars in M53 and \\citet{pyc01} derived a slope of $-2.88$ for the fundamental mode and a slope of $-3.1$ for the first overtone mode stars in M55. The slope for NGC 5466 derived in this study agrees also well with the theoretical values of --3.04 by \\citet{san01} and --3.05 by \\citet{tem02}."
    },
    "0404/astro-ph0404543_arXiv.txt": {
        "abstract": "We present the results of \\chandra\\ observations of six radio-loud quasars (RLQs) and one optically bright radio-quiet quasar (RQQ) at \\hbox{\\(z\\approx\\)~4.1--4.4}.  These observations cover a representative sample of RLQs with moderate radio-loudness (\\hbox{\\(R\\approx\\)~40--400}), filling the \\xray\\ observational gap between optically selected radio-quiet quasars (predominantly \\hbox{$R \\lesssim 2$--10}) and the five known blazars at $z>4$ (\\hbox{$R\\approx$~800--27000}), where \\hbox{$R$=\\(f_{\\rm 5~GHz}/f_{\\rm 4400~\\mbox{\\scriptsize\\AA}}\\)} (rest frame). We study the relationship between \\xray\\ luminosity and radio-loudness for quasars at high redshift and constrain RLQ \\xray\\ continuum emission and absorption.  From a joint spectral fit of nine moderate-$R$ RLQs observed by \\chandra, we find tentative evidence for absorption above the Galactic $N_{H}$, with a best-fit neutral intrinsic column density of \\hbox{$N_{H}=2.4^{+2.0}_{-1.8}\\times 10^{22}$~cm$^{-2}$}, consistent with earlier claims of increased absorption toward high-redshift RLQs.  We also search for evidence of an enhanced jet-linked component in the \\xray\\ emission due to the increased energy density of the cosmic microwave background (CMB) at high redshift, but we find neither spatial detections of \\xray\\ jets nor a significant enhancement in the \\xray\\ emission relative to comparable RLQs at low-to-moderate redshifts.  Overall, the \\hbox{$z\\approx$~4--5} RLQs have basic \\xray\\ properties consistent with comparable RLQs in the local universe, suggesting that the accretion/jet mechanisms of these objects are similar as well. ",
        "introduction": "The study of the $z>4$ universe has become a major theme in astronomy over the past decade. In particular, \\xray\\ observations of quasars at $z > 4$ reveal the physical conditions in the vicinities of the first massive black holes to form in the Universe. Prior to 2000, there were only six published \\xray\\ detections of quasars at $z\\ge4$ (see Kaspi, Brandt, \\& Schneider 2000; hereafter KBS00).  In the last four years, through the unprecedented sensitivities of \\chandra\\ and \\xmm\\ as well as optical wide-area surveys such as the Sloan Digital Sky Survey (SDSS; York et al.\\ 2000), the number has risen to nearly one hundred.  Equally important are wide-area radio surveys, such as Faint Images of the Radio Sky at Twenty Centimeters (FIRST; e.g., Becker et al.\\ 1995) and the NRAO VLA Sky Survey (NVSS; Condon et al.\\ 1998), to constrain the properties of radio cores, jets, and lobes. Radio-selected samples of radio-loud quasars (RLQs), although sampling a minority of the total quasar population, are less prone to selection effects than optical surveys, as radio emission is not affected by absorption due to dust.  The radio-loudness parameter, as defined by Kellermann et al.\\ (1989), is given by \\hbox{$R$=\\(f_{\\rm 5~GHz}/f_{\\rm 4400~\\mbox{\\scriptsize\\AA}}\\)} (rest frame), where quasars with $R\\ge10$ are RLQs and those with $R<10$ are radio-quiet quasars (RQQs).  Some of the first $z>4$ quasars studied in X-rays were a small group of highly radio-loud blazars\\footnote{ The term ``blazar'' is admittedly vague. As defined in \\S 1.3 of Krolik (1999), blazars include BL Lacs and optically violently variable (OVV) quasars.  Previous studies have suggested that four of the $z>4$ quasars discussed in this paper be granted blazar status (PMN~0525$-$3343, Worsley et al.\\ 2004a; RX~J1028$-$0844, Grupe et al.\\ 2004; GB~1428$+$4217, Worsley et al.\\ 2004b; GB~1508$+$5714, Yuan et al.\\ 2003), and we tentatively place GB~1713$+$2148 (Vignali et al.\\ 2003a) in the same category based on its extreme radio-loudness. } (\\hbox{$R\\approx$~800--27000}) in which the \\xray\\ radiation is probably dominated by a jet-linked component. These blazars represent a minority of the RLQ population, however, and so are not suitable for representative statistical studies. Recently, {\\it Chandra's\\/} Advanced CCD Imaging Spectrometer (ACIS) has obtained \\xray\\ detections of many RQQs (predominantly with \\hbox{$R \\lesssim 2$--10}), which represent the majority population of quasars in the local universe and so too, apparently, at high redshift (e.g., Schmidt et al.\\ 1995; Stern et al.\\ 2000). Through these studies (e.g., Brandt et al. 2002; Vignali et al. 2001, 2003a,b), the basic \\xray\\ properties of \\hbox{$z\\approx$~4--5} RQQs have become reasonably well understood. Perhaps surprisingly, given the strong evolution of the quasar population (e.g., Boyle et al.\\ 2000; Fan et al.\\ 2001b), high-redshift RQQs have remarkably similar \\xray\\ and broad-band properties to RQQs at low-to-moderate redshift (after controlling for luminosity effects; e.g., Vignali, Brandt, \\& Schneider 2003, hereafter VBS03), suggesting that the accretion mechanism does not dramatically evolve with time. Prior to the current paper, however, there were only four published \\xray\\ detections of $z>4$ RLQs with \\hbox{$R\\approx$~10--1000} corresponding to typical RLQs in the local universe (excluding the blazar PMN~0525$-$3343 which has $R\\approx 800$).  Here we report \\chandra\\ detections of six \\hbox{$z\\approx$~4.1--4.4} radio-selected RLQs and one RQQ (FIRST 0747+2739) which is radio detected and optically bright. These objects have been selected as prime targets for \\xray\\ observations based on their bright optical and radio fluxes, and they have been chosen to represent the majority population of RLQs.  Four are FIRST-selected objects published in Benn et al.\\ (2002), two are from the Cosmic Lens All Sky Survey (CLASS; Myers et al.\\ 2003) published in Snellen et al.\\ (2001), and one is from the SDSS, first presented in Anderson et al.\\ (2001). Five of the RLQs have published detections at both 1.4~GHz and 5~GHz and have flat radio spectra with $\\alpha_{\\rm r} \\approx -0.4$~to $+$0.3 ($f_{\\nu} \\propto \\nu^{\\alpha}$). In the case of FIRST~1309$+$5733, the upper limit on the 5~GHz flux density gives $\\alpha_{\\rm r} < 0.06$, so it is possible that this object is also a flat-spectrum RLQ.  None of the quasars shows any significant structure in the FIRST radio images (these are the best radio images available, with an angular resolution of $\\approx 5$\\arcsec). While the RQQ FIRST~0747$+$2739 is relatively radio-weak compared to our other objects, it is one of the optically brightest quasars at $z>4$, and an overabundance of C~{\\sc iv} absorption in the optical spectrum suggests that it may be a transition Broad Absorption Line (BAL) quasar (Richards et al.\\ 2002).  Given that RLQs are often substantially \\xray\\ brighter than RQQs with comparable optical luminosity (e.g., Zamorani et al.\\ 1981; Worrall et al.\\ 1987), these objects are prime targets for \\chandra\\ snapshot (e.g., \\hbox{$\\approx$~4--10 ks}) observations. Indeed, all seven of our sources are detected with a minimum of $\\approx$~20 counts.  \\figurenum{1} \\centerline{\\includegraphics[angle=0,width=8.5cm]{bassett.fig1.ps}} \\figcaption[bassett.fig1.ps]{\\footnotesize Radio loudness versus redshift for \\xray\\ detected quasars at $z\\geq 4$. The large symbols represent the moderate-$R$ RLQs (triangles), blazars (stars), and RQQs (circles).  The filled symbols represent the quasars in the current sample.  Arrows show upper limits for radio loudness at the $\\approx 3\\sigma$~level based on the 1.4~GHz radio flux density. Note that, excluding upper limits, the present sample more than doubles the number of objects with moderate radio-loudness (\\hbox{$R \\approx$~10--1000}), filling the \\xray\\ observational gap between the RQQs and the blazars. \\label{rlvz}} \\centerline{} \\centerline{}  This brings the total sample of $z>4$ \\xray\\ detected RLQs to 15, more than doubling the number of moderately radio-loud quasars and filling the gap (as shown in Figure~1) between the many \\xray\\ observations of $z>4$ RQQs and the extreme blazars.\\footnote{ A regularly updated list of $z>4$ quasars with X-ray detections can be found at http://www.astro.psu.edu/users/niel/papers/highz-xray-detected.dat} To date there are $\\approx 50$ known $z>4$ RLQs; thus the objects in our sample comprise $\\approx 10$\\% of those known.  Our targets have been drawn from surveys covering a significant fraction of the sky (6400~deg$^{2}$, Snellen et al.\\ 2001; 4030~deg$^{2}$, Benn et al. 2002) and so are likely to be representative of $z>4$ RLQs as a whole, constituting the first set large enough to serve as a representative sample for reliable statistical evaluations of $z>4$ RLQ \\xray\\ properties.  Recent work (e.g., Cirasuolo et al.\\ 2003; Ivezi\\'{c} et al.\\ 2002, 2004) has sparked new controversy over the existence of bimodality in the radio-loudness distribution of quasars, suggested first by Strittmatter et al.\\ (1980).  Current results suggest that, although quasars cover the full range of radio-loudness, the underlying distribution does consist of two regimes with a local minimum in between.  As reported by Ivezi\\'{c} et al.\\ (2002) for a sample of 3225 SDSS quasars matched to FIRST sources, $8\\% \\pm 1\\%$ of all quasars with \\hbox{$i^{\\ast}<$~18.5} are radio-loud, apparently independent of optical luminosity or redshift, in the range $z<2$. Figure~2 shows the radio-loudness distribution for $z>4$ quasars with radio detections or upper limits in FIRST, with the best-fit distribution from Ivezi\\'{c} et al. (2004) overlaid (dashed curve). While the blazars lie on the extreme radio-loud tail of the distribution, the RLQs from this work fall near the peak of the radio-loud regime, likely representing the majority population of $z>4$ RLQs.  By integrating the best-fit curve directly, we find that our sample of RLQs with $R=40$--400 is representative of $\\simgt 50$\\% of the total RLQ population, while the blazars with $R>1000$ represent $\\simlt 5$\\% of all RLQs.  \\figurenum{2} \\centerline{\\includegraphics[angle=0,width=8.5cm]{bassett.fig2.ps}} \\figcaption[bassett.fig2.ps]{Radio-loudness distribution for known $z \\geq 4$ quasars with detections or upper limits in the FIRST survey.  The solid histogram shows all $z \\geq 4$ quasars with radio detections, and the dotted histogram shows all those with radio upper limits.  The dashed curve is adapted from the radio-loudness distribution in Ivezi\\'{c} et al. (2004), determined from a sample of $\\sim 10,000$ objects detected by FIRST and SDSS (with arbitrary normalization). The objects in the current sample (hatched shading) represent the majority population of RLQs, while the blazars (solid shading) lie on the extreme radio-loud tail. \\label{rlhist}} \\centerline{} \\centerline{}  Throughout this paper we adopt $H_{0}$=70 km s$^{-1}$ Mpc$^{-1}$ in a $\\Lambda$-cosmology with $\\Omega_{\\rm M}$=0.3 and $\\Omega_{\\Lambda}$=0.7 (e.g., Spergel et al.\\ 2003).  \\begin{table*} \\footnotesize \\tabletypesize{\\footnotesize} \\caption{\\chandra\\ Observation Log} \\begin{center} \\begin{tabular}{lccccccc} \\tableline \\tableline \\multicolumn{1}{c}{Object} &  & FIRST & FIRST & $\\Delta_{\\rm Rad-X}$\\tablenotemark{a} & \\xray & Exp.~Time$^{\\ \\rm b}$ & \\\\ \\multicolumn{1}{c}{Name} & $z$ & $\\alpha_{2000}$ & $\\delta_{2000}$ & (arcsec) & Obs.~Date & (ks) & Ref. \\\\ \\tableline FIRST~0725$+$3705 & 4.33 & 07 25 18.26 & $+$37 05 18.3 & 0.2 & 2002 Dec 17 & 4.89 & (1) \\\\ FIRST~0747$+$2739 & 4.11 & 07 47 11.20 & $+$27 39 04.0 & 0.9 & 2002 Dec 03 & 4.96 & (1,3) \\\\ SDSS~0839$+$5112 & 4.39 & 08 39 46.20 & $+$51 12 02.9 & 0.3 & 2003 Jan 23 & 4.90 & (2,4) \\\\ CLASS~J0918$+$0637 & 4.19 & 09 18 24.39 & $+$06 36 53.3 & 0.2 & 2002 Dec 10 & 4.90 & (2) \\\\ FIRST~1309$+$5733 & 4.27 & 13 09 40.69 & $+$57 33 10.0 & 0.4 & 2003 Jul 07 & 4.64 & (1) \\\\ CLASS~J1325$+$1123 & 4.40 & 13 25 12.48 & $+$11 23 30.0 & 0.5 & 2003 Mar 02 & 4.70 & (2) \\\\ FIRST~1423$+$2241 & 4.32 & 14 23 08.24 & $+$22 41 58.2 & 0.2 & 2002 Dec 02 & 4.70  & (1) \\\\ \\tableline \\end{tabular} \\vskip 2pt \\parbox{5.3in} {\\small\\baselineskip 9pt \\footnotesize \\indent {\\sc Note. ---} Radio and optical identifications of the quasars presented here can be found in the papers cited in the reference column.  Radio positions are from the FIRST catalog (Becker et al.\\ 1995). Units of right ascension are hours, minutes, and seconds, and units of declination are degrees, arcminutes, and arcseconds. \\\\ $^{\\rm a}$ Distance between the radio and \\xray\\ positions. $^{\\rm b}$ The \\chandra\\ exposure time has been corrected for detector dead time. \\\\ {\\sc References. ---} (1) Benn et al.\\ 2002; (2) Snellen et al.\\ 2001; (3) Richards et al.\\ 2002; (4) Anderson et al.\\ 2001. } % \\end{center} \\vglue-0.9cm \\end{table*} \\normalsize ",
        "conclusions": "We have presented new \\chandra\\ detections of six RLQs and one RQQ at \\hbox{$z\\approx$~4.1--4.4}, thereby significantly enlarging the number of $z>4$ RLQs with X-ray detections. Our RLQ targets have moderate $R$ parameters of \\hbox{$\\approx$~40--400}; they are more representative of the overall RLQ population than many of the \\hbox{high-$R$} ($R>500$) RLQs at $z>4$ previously studied in X-rays. Our main results are the following: \\begin{enumerate} \\item Moderate-$R$ RLQs at $z>4$ have enhanced X-ray emission relative to RQQs of the same redshift and optical luminosity. The degree of enhancement, a factor of $\\approx 2$ in X-ray luminosity, is consistent with that observed at low redshift for comparable objects and is presumably associated with a jet-linked emission component. As in the local universe, the X-ray-to-optical flux ratios of \\hbox{moderate-$R$} RLQs at $z>4$ are intermediate between those of RQQs and those of \\hbox{high-$R$} RLQs. \\item Joint X-ray spectral analyses of moderate-$R$ RLQs show that, on average, their spectra can be fitted with a power-law model and Galactic absorption; the average best-fit power-law photon index of \\hbox{$\\Gamma = 1.65\\pm0.15$} is consistent with that for comparable objects at low redshift. Adding an absorption component to the fit, we find tentative evidence (90\\% significance) for excess intrinsic X-ray absorption on average, and our best-fit estimate of the cold absorption column of \\hbox{$N_{H}=2.4^{+2.0}_{-1.8} \\times 10^{22}$~cm$^{-2}$} is consistent with previous claims of increased absorption toward high-redshift RLQs.  The somewhat higher best-fit photon index of \\hbox{$\\Gamma = 1.80^{+0.23}_{-0.25}$} in this model is still consistent with previous estimates for RLQs at low redshift. \\item We do not detect extended X-ray emission associated with jets from any of our RLQs, which, if dominated by the IC/CMB process as suggested, should be significantly enhanced at $z>4$.  If the jets in these objects are similar to the more modest \\xray\\ jet recently discovered from the $z=4.3$ blazar GB~1508$+$5714 (which accounts for only $\\approx 3$\\% of the quasar \\xray\\ luminosity), we expect that the jets are below the sensitivity limit of our snapshot observations. \\item Overall, our data are consistent with the idea that the dominant X-ray production mechanisms of $z>4$ and $z\\approx 0$ RLQs are similar. \\item  Richards et al.\\ (2002) suggested that an abundance of C {\\sc iv} absorption systems in the optical spectrum of FIRST~0747$+$2739 are intrinsic to the quasar and that it represents a class of objects in a transition phase between normal and BAL quasars. From our observations, FIRST~0747$+$2739 does not appear to be notably \\xray\\ weak, but we cannot rule out neutral intrinsic absorption columns with \\hbox{$N_{H} < 5.5 \\times 10^{22}$~cm$^{-2}$}. \\end{enumerate} Finally, several of the RLQs we have detected with \\chandra\\ are among the X-ray brightest non-blazars known at $z>4$ (see Figure~4a). These are prime targets for follow-up X-ray spectroscopy with \\xmm\\ (with the caution that in some cases nearby bright sources may contaminate the results). Such spectroscopy can reveal if the majority population of RLQs at high redshift shows intrinsic X-ray absorption."
    },
    "0404/astro-ph0404129.txt": {
        "abstract": "{We present results on number counts of Extremely Red Objects (EROs) in a 2850 arcmin$^2$ near-infrared survey performed in European Large Area ISO Survey (ELAIS) fields at $K<17.5$.  Counts of EROs are extended to brighter levels than available previously, giving $0.002 \\pm 0.001 \\rm arcmin^{-2}$ at $K<16.5$ and consistent numbers with literature values at fainter magnitudes. Photometric redshifts from HYPERZ as well as GRASIL model SEDs of galaxies imply that our EROs are located in the range $z=0.7-1.5$, with the bulk of the population at $z\\sim1$. Taking advantage of the ISO data in the fields, we use mid-IR detections to constrain the number of dusty EROs, and also discuss the superior capabilities of Spitzer Space Telescope to detect dusty EROs. Both the mid-IR data and the use of colour-colour diagrammes indicate that at most 10-20\\% of the EROs in this bright regime are dusty starbursting systems.  The space density of our EROs, interpreted to be counterparts of local $>2-3L^{\\star}$ massive galaxies at around $z\\sim1$, is estimated to be $\\approx2\\times10^{-5} \\ {\\rm  Mpc^{-3}}$, which is consistent with local values.  Furthermore, the cumulative number counts at our bright magnitudes are remarkably well fitted by pure luminosity evolution models. ",
        "introduction": "% (Elston, Rieke \\& Rieke \\cite{El88}; Elston, %Rieke \\& Rieke \\cite{El89} ; Mcarthy, Persson \\& West %\\cite{Mc92}). % (eg. Daddi et al. \\cite{Da00}; McCarthy et %al. \\cite{McC01})  Extremely Red Objects (EROs, selected for example by $R-K>5, I-K>4$ colours) have received much focus recently by virtue of their potential as a powerful window into the galaxy formation era.  By and large, the majority of EROs have been incorporated into a bimodal population, where the extreme red colours are attributed either to old passively evolving distant $(z>1)$ elliptical galaxies or to extremely dust reddened starburst galaxies (see e.g.\\ recent papers by Cimatti et al. \\cite{Ci03}, Wold et al. \\cite{Wo03}, Yan \\& Thompson \\cite{Ya03}, Takata et al. \\cite{Ta03}, Daddi et al. \\cite{Da02}, Smail et al. \\cite{Sm02}, Roche et al. \\cite{Ro02}, and references therein for earlier pioneering ERO surveys).  It is the class of aged early type galaxies whose number densities provide the strongest constraints on models of galaxy evolution. On the other hand, the dusty EROs could be related to the (ultra) luminous IR-galaxies producing the bulk of the total energy in the Universe since the recombination era (see e.g.\\ Elbaz \\& Cesarsky \\cite{El03} for a review).  %(e.g. Cohen et al. \\cite{Co99}) %(e.g. Smail et al. \\cite{Sm99}).  The dusty ERO population contaminates the otherwise, in principle, clean samples of high redshift massive ellipticals. This is the reason why the separation of elliptical EROs from dusty EROs, using photometric, spectroscopic, and morphological analyses, has received the main attention of ERO studies recently.  In general, a large number of massive galaxies at redshifts of unity and over would support a traditional monolithic collapse scenario where galaxies form at high redshift in a single collapse and then evolve passively (e.g. Eggen et al. \\cite{Eg62}; Larson \\cite{La75}).  Significantly smaller numbers of massive high-$z$ ellipticals on the other hand fit predictions from the hierarchical assembly scenarios, where massive galaxies are formed only relatively recently from the merging of smaller units (e.g.\\ White \\& Rees \\cite{Wh78}; White \\& Frenk \\cite{Wh91}; Somerville \\& Primack \\cite{So99}; Cole et al. \\cite{Co00}).  Since EROs are extremely faint optically $(R\\ge 24)$, their redshifts and spectral and morphological properties have remained largely unknown until very recently. Cimatti et al. (\\cite{Ci02}) carried out a VLT spectroscopic survey of 45 EROs with $R-K>5$ and $K_{s}<19.2$. They identified approximately $1/3$ of the EROs as old elliptical systems, $1/3$ as dusty starburst galaxies, while $1/3$\u0081\u00a0\\ remained unidentified. The mean redshift was found to be $z\\sim 1.0$ for both populations.  Yan, Thompson, \\& Soifer (\\cite{Ya04}) on the other hand find a small 10-15\\% fraction of isolated passive systems. Yan \\& Thompson (\\cite{Ya03}), Cimatti et al. (\\cite{Ci03}), Moustakas et al. (\\cite{Mo04}), and Gilbank et al. (\\cite{Gi03}) have used HST morphologies to differentiate between the classes: the results have further complicated the picture, as a large fraction (25-65\\%) of EROs seem to be disks (though see also Moriondo et al.\\ \\cite{Mo00}).  The expectation originally had been a clearer distinction into spheroids on one side, and irregular and interacting types on the other, expected for extremely dusty star forming galaxies. The relative fractions of different types of EROs thus still remain uncertain (due to for example different selection criteria used) and the results of testing surface densities against galaxy formation scenarios are inconclusive.  In addition to spectroscopic and morphological methods mentioned above, one may separate the old elliptical and dusty EROs by means of longer wavelength observations: any ERO detected at mid-IR to radio wavelengths should belong to the dusty population. Systematically this has been attempted by Mohan et al. (\\cite{Mo02}) in the sub-mm to radio and Smail et al.\\ (\\cite{Sm02}) in the radio. Similar fractions (albeit with wide spread) as referred to above of dusty EROs were found in the latter study, while Mohan et al. (\\cite{Mo02}) find a much lower fraction.  In this paper we attempt to separate EROs, for the first time, based on their mid-IR properties.  As with far-IR to radio methods, this is a clear-cut definition, since the difference of ellipticals and starbursting galaxies is very large - distant evolved ellipticals would not be detected with ISO whereas dusty EROs should have strong mid-IR flux. Also, mid-IR avoids the identification problems of e.g.\\ large sub-mm beams. We have surveyed an ISO/ELAIS (European Large Area \\textit{ISO} Survey; Oliver et al. \\cite{Ol00}) field in the near-IR and matched the results with an optical dataset.  We make use of the newly available Final ELAIS Catalogue (Rowan-Robinson et al. \\cite{Row04}). In this paper we consider the surface densities of EROs resulting from a wide-field relatively shallow $J$ and $K$-band survey.  We shall present the results from a deeper NIR survey around faint mid-IR detections in subsequent papers (V\\\"ais\\\"anen \\& Johansson \\cite{Va04}; Johansson, V\\\"ais\\\"anen \\& Vaccari 2004, in prep.).  The structure of this paper is as follows. The observational data are presented in \\S 2, and in \\S 3 we define and extract the ERO sample from the photometric catalogues.  In \\S4 we discuss the dusty vs.\\ evolved ERO separation and in \\S5 the SEDs of the detected EROs.  The resulting surface densities of EROs are presented in \\S6. We assume throughout this paper a flat ($\\Omega_{0}=1$) cosmology with $\\Omega_{m}=0.3$, $\\Omega_{\\Lambda}=0.7$ and $H_{0}=70 \\ \\rm km s^{-1} Mpc^{-1}$. ",
        "conclusions": "\\subsection{ERO counts and surface densities}  The numbers of detected EROs are given in Table \\ref{Tab2}, and the corresponding cumulative number counts are plotted in Fig.\\ref{allcounts}.  The numbers plotted are the total ERO counts: based on the previous discussion, the fraction of dusty EROs is low ($\\sim15$\\%) and does not have significant effect even if subtracted from the counts.   \\begin{table*} \\begin{center} \\begin{tabular}{lcccccccccc} \\hline \\hline \\noalign{\\smallskip} &   \\multicolumn{3}{c}{$r-K>5.5$} & \\multicolumn{3}{c}{$r-K>5.8$} & \\multicolumn{3}{c}{$i-K>4.4$} \\\\ \\noalign{\\smallskip} $K$ limit &  & Frac. & $\\Sigma_{K}$\u0081\u00a0& & Frac. & $\\Sigma_{K}$\u0081\u00a0& & Frac. & $\\Sigma_{K}$\u0081\u00a0 \\\\ (mag)     &  N & \\% & $\\rm arcmin^{-2}$ & N & \\% & $\\rm arcmin^{-2}$ & N & \\% & $\\rm arcmin^{-2}$ & \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip}  $K \\leq 16.5$  & 7 & 0.004 & 0.002 & 2 & 0.001 & 0.0007 & 3 & 0.002 &  0.001 \\\\ $K \\leq 17.0$  & 19 & 0.008 & 0.007 & 7 & 0.003 & 0.002 & 11 & 0.005 &  0.004 \\\\ $K \\leq 17.5$  & 51 (35) & 0.012 & 0.018 & 15 (11) & 0.004 & 0.005 & 19 (15) & 0.005 & 0.007 \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{center} \\caption[]{The sample of Mt.Hopkins EROs.  Cumulative counts are given for different ERO colour selection criteria.  Total numbers (N), Fraction of EROs compared to the total galaxy count (``Frac.'') is calculated using galaxy counts in these same fields (V\\\"ais\\\"anen et al.\\cite{Va00}).  At $K\\leq17.5$ N shows the completeness corrected value, and the raw count is given in parentheses.} \\label{Tab2} \\end{table*} %---------------------------------------------------------------    Our wide survey from Mt.Hopkins extends the ERO number counts to brighter magnitudes than observed before.  We report the first ERO counts at $K<16.5$, where the $r-K>5.5$ (ie.\\ $R-K>5$) selection yields a surface density of $0.002 \\pm 0.001 \\rm arcmin^{2}$, and the $i-K>4.4$ (ie.\\ $I-K>4$)  selection $\\approx50$\\% of this value. At $K<17$ magnitudes we arrive at $0.007 \\pm 0.002 \\rm arcmin^{2}$ of $r-K$ EROs. Beyond $K>17$ mag the Mt.Hopkins survey is starting to be affected by incompleteness, and based on simulations performed in V00 for the same dataset, we estimate a factor of 1.5 correction to the $K<17.5$ cumulative ERO density.  The resulting $r-K$ selected surface density at $K<17.5$ is $0.018 \\pm 0.003 \\rm arcmin^{2}$.  We compare these numbers to the only two other field ERO surveys with enough sky coverage to reach these bright magnitude levels: The largest ERO survey to date Daddi et al.\\ (\\cite{Da00}; with 701 arcmin$^{2}$), yielded $0.003 \\pm 0.002 \\rm arcmin^{2}$ at $K<17$ and $0.02 \\pm 0.006 \\rm arcmin^{2}$ at $K<17.5$.  These are totally consistent with our result.  The Yan \\& Thompson (\\cite{Ya03}; 409 arcmin$^{2}$) HST ERO counts were selected using $I-K>4.0$.  At $K<17$ they agree with our corresponding $i-K>4.4$ selection within the errors.  At $K<17.5$ there is a clear discrepancy, however. It is our $r-K$ selected EROs which are in closer agreement with the Yan \\& Thompson  ERO counts -- we systematically find a factor of $\\sim2$ less $i-K$ selected EROs than $r-K$ selected ones.  Especially at the faintest bin there appears to either be more incompleteness than we expect, or the $i-K>4.4$ selection does not pick up as many bright $K\\sim17$ galaxies at $z\\sim1$ than the $r-K$ selection does.  Indeed, the discussion on ``excess'' $r-K$ EROs in the next Section~\\ref{evolved} and the photometric redshifts derived for the EROs hint at the latter possibility.   \\begin{figure*} \\resizebox{18.0cm}{!}{\\includegraphics{allcnts_eros.ps}} \\caption[]{Cumulative ERO counts are shown for $R-K>5$ and $R-K>5.3$ selections (ie. $r-K>5.5$ and $r-K>5.8$), in the left and right panels, respectively, with solid triangles. Additionally, the $i-K>4.4$ selected EROs are plotted in the left panel with circles. Error bars are Poissonian. Our other field ERO counts from a deeper survey (IRTF, V\\\"aisanen \\& Johansson 2004) are shown as stars. Squares show the EROs of Daddi et al.\\ (\\cite{Da00}) and the asterisks the Roche et al.\\ (\\cite{Ro02}) EROs in the left panel.  The left panel shows $R-K>5$ selected EROs of Daddi et al. (\\cite{Da00}) and Roche et al.\\ (\\cite{Ro02}).  In addition to the Daddi et al.\\ $R-K>5.3$ selected EROs the right panel plots the corresponding Smith et al.\\ (\\cite{Smi02}) and Smail et al.\\ (\\cite{Sm02}) EROs. We plot PLE models of Daddi et al.\\ calculated using the indicated $R-K$ cuts. The curves are different by their galaxy formation redshifts, as indicated in the figure.  The two dashed-lines employ a different LF -- the higher curve results from a 2MASS normalized LF, and the lower from that of Marzke et al.\\ (\\cite{Ma94}). } \\label{allcounts} \\end{figure*}    \\subsection{Evolved ellipticals at $z\\sim1$} \\label{evolved}   The strength, and motivation, of ERO searches has been to look for the most massive galaxies at redshifts of unity and over (see e.g.\\ Saracco et al.\\ \\cite{Sa03}), since it is exactly these which place tight constraints on galaxy formation scenarios. The dusty vs.\\ early type ERO discriminators showed that most of our ERO sample is consistent with being early type galaxies.  However, we argue that the clearest population of ellipticals comes from the class of $r-K>5.5$ EROs which have $r-K<4.4$ -- these constitute 60\\% of all our EROs.  Fig.~\\ref{ikrk} shows an $i-K$ vs.\\ $r-K$ plot, where these ``excess'' EROs populate the upper right quadrant (see also Figs.~\\ref{jk_sep} and~\\ref{eroseds}).  Intuitively, the upper left quadrant galaxies should be sources at a redshift where a significant 4000 \\AA \\ break falls right in between the $r$ and $i$ bands: i.e.\\ ellipticals or very early type spirals at a redshift $z \\sim 0.8$. We calculated numerous GRASIL models, and it turns out to be fairly difficult to obtain colours in this ERO regime.  Any significant on-going star-formation drops the $r-K$ colour down, as noted also by e.g.\\ Yan \\& Thompson (\\cite{Ya03}; see their Fig.9.).  The effect is seen in the solid curve of the right panel of Fig.~\\ref{ikrk}: the elliptical model has a formation redshift of $z_{\\star} = 3$ and at redshifts of $z>1.5$ the effects of the starburst turn the $r-K$ colour sharply down.  Significant amounts of extinction, including edge-on spirals, do not bring models to the $r-K>5.5, i-K<4.4$ region either: the dotted line depicts the extremely dusty HR10 model as an example.  One gets closest to the $r-K>5.5, i-K\\sim4$ region by merely redshifting an old present-day elliptical to Sa galaxy to the appropriate redshift.  This is shown as the dashed curve, a 13 Gyr old Sa spiral; any $>10$ Gyr old elliptical produces a similar result.  While this is an unphysical model, it serves to point out that while pure reddening is unlikely to reach this part of the colour-colour diagramme, an old stellar population does that.  We thus conclude that EROs which are selected by $r-K>5.5$ (i.e.\\ $R-K>5$) which are not EROs by $i-K>4.4$ (i.e.\\ $I-K>4$) are mostly early type galaxies.  In fact, at $i-K<4.2$ there should not be any contamination from dusty galaxies (compare the vertical dotted line to the dusty model curves).  We also plot a sample of averaged SEDs of our EROs in Fig.~\\ref{eroseds}.  The ``excess'' EROs with $i-K<4.2$ are plotted with squares and are well fit by very evolved galaxy SEDs, as just discussed.  The critical observed band is the $i$ band, where the difference between elliptical and dusty starbursts (M82 plotted) is the largest at the redshifts of our targets. Note that the $g$-band is biased to the bluest sources, since the average is calculated from those EROs detected in the corresponding band, and only 25\\% of the sources have a $g$-band detection.  The $r-K$ and $i-K$ selections alone are quite different (see e.g.\\ Scodeggio \\& Silva \\cite{Sc00}).  At least with bright ERO surveys such as ours, it seems that $r-K$ selects more passive systems. This is consistent with spectroscopic followup of $R-K>5.3$ surveys finding significant amounts of ellipticals (e.g.\\ Cimatti et al.\\ \\cite{Ci02}) and those following up $I-K>4$ selections finding much more star formation (e.g.\\ Yan et al.\\ \\cite{Ya04}).   \\begin{figure*} \\resizebox{18.5cm}{!}{\\includegraphics{ikrk_eros.ps}} \\caption[]{$i-K$ vs.\\ $r-K$ colour-colour magnitude diagramme of the verified EROs is plotted in the left panel. The mid-IR detected EROs is overplotted as a large solid symbol. Several GRASIL models are plotted in both panels:  the models are drawn until redshift $z=2$, and the $z=1$ locations are indicated by small solid circles along the curves. The reddening direction is indicated in the left panel by the diagonal arrow which correcponds to $A_{V}=1$ extinction. } \\label{ikrk} \\end{figure*}    \\subsubsection{Number densities}  There are 33 ``excess'' $r-K$ EROs with $i-K<4.4$.  They appear to be at redshifts $z=0.7-1.2$: BC SED fitting favours redshifts below unity (Fig.~\\ref{zphot}), while the GRASIL models suggest redshifts at $z\\approx1-1.2$ (eg.\\ Fig.~\\ref{ikrk}). These EROs have typical magnitudes of $K\\approx17.2$, which would make them at least $\\sim2L^{\\star}$ galaxies, taking into account passive evolution since that redshift. Very simply, assuming that this group of EROs consists of ellipticals and is constrained in the mentioned redshift range we derive a co-moving volume of $2\\times10^{6}$ Mpc$^{3}$ for the Mt.Hopkins survey area, and thus a space density of $\\sim2\\times10^{-5}$ Mpc$^{-3}$. Taking $K$-band LF parameters from Kochanek et al.\\ (\\cite{Ko01}), this is a significant part of the expected density of $\\approx8\\times10^{-5}$ Mpc$^{-3}$ for local $>2L^{\\star}$ early type galaxies.  Given the incompleteness in our survey at $K>17$, and that we did not include ellipticals at $i-K>4.4$ in this calculation these ERO counts are a very conservative lower limit, and thus indicate that a major part of massive present day $\\sim 2-3L^{\\star}$ ellipticals were in place at $z\\sim1$.  The space density just derived is intriguingly close to those estimated for the higher redshift sub-mm population, which are suspected to be the most luminous starbursts, perhaps the progenitors of present day massive ellipticals (see e.g.\\ discussion in Scott et al.\\ \\cite{Sc02}).     \\subsection{Constraints on galaxy formation scenarios}  Figure~\\ref{allcounts} plots our ERO counts along with several other recent surveys and pure luminosity evolution (PLE) models of Daddi et al. (\\cite{Da00}).  All curves depict ellipticals with a $\\tau = 0.1$ Gyr initial starburst, and passive evolution afterwards.  The different lines have formation redshifts of $z_{\\rm form} = 2.4, 3.0,$ and 10.  All the models use a Marzke et al.\\ (\\cite{Ma94}) LF, except the uppermost dashed line, which instead uses a 2MASS based LF (Kochanek et al. \\cite{Ko01}) resulting in a higher normalization by approximately a factor of two.  We stress that {\\em all EROs are counted} in the figure, i.e.\\ no corrections for the dusty population have been made. As was seen previously, this is perfectly reasonable for our very bright sample of EROs.  However, even a very large correction (not supported by our observations) of 50 \\% of the EROs being in fact dusty, would not change the conclusion significantly.  At somewhat fainter magnitudes, recent studies have shown (Cimatti et al. \\cite{Ci02}, Smail et al.\\ \\cite{Sm02}) that the percentage of dusty EROs is likely somewhere around 30--60\\%.  The PLE model using a formation redshift of $z_{\\rm form}=3$ fits all the data remarkably well.  Formation redshifts slightly lower than this also fit well our bright ERO counts, but start underpredicting the numbers at fainter magnitudes. On the other hand, corrections for dusty galaxies might be larger at the fainter levels (see eg.\\ Smith et al.\\ \\cite{Smi02}). However, models with $z_{\\rm form}<2.4$ start underpredicting the counts at all levels, especially when using the $R-K>5.3$ selection criteria.  The very highest formation redshifts of $z_{\\rm form}=10$ and over, predict steeper number counts and are more difficult to fit to both the bright and faint end of ERO counts with the same normalization.  The figure shows two different LFs giving a factor of $\\sim2$ difference.  Note also that eventhough the $R-K>5.3$ criterion eliminates nearly all the (possibly) lower redshift $z\\sim0.8$ EROs, there is no significant difference in the fits of PLE models in the right and left panels of Fig.~\\ref{allcounts}.  Using the 2MASS normalized LF it is in fact slightly more difficult to fit both $R-K>5$ and $R-K>5.3$ counts simultaneously.  The $R-K>5$ counts are overpredicted by a factor of $1.5-3$ with all except the very lowest $z_{\\rm form}=2.2-2.4$ formation redshifts.  Good fits in the range $K=16.5-18.5$ are again acquired with $z_{\\rm form}\\ge4$, but the faintest counts are overpredicted.  Roche et al.\\ (\\cite{Ro02},\\cite{Ro03}) find their 2MASS-LF based PLE models significantly overpredicting $I-K>4$ ERO counts, and show that certain amount of merging and density evolution fits the counts much better in particular in the fainter $K>20$ regime.  That a range of formation redshifts are necessary to model EROs is seen for example by comparing Figs.~\\ref{ikrk} and~\\ref{allcounts} (see also Cimatti et al.\\ \\cite{Ci03}): though models with formation redshifts of $z\\approx3$ fit well the {\\em counts}, the {\\em colours} produced by such a scenario are not red enough for many EROs found in this survey, and others.  Very red colours of $r-K>6.5$ or $i-K>5$ can not be produced without pushing $z_{\\rm form}$ closer to 10 (unless all such extreme EROs are heavily reddened starbursts).  We do not investigate hierarchical formation scenarios in more detail here. We merely point out that predictions from some recent models (Cole et al.\\ \\cite{Co00} as presented in Smith et al.\\ \\cite{Smi02}), fall short an order of magnitude in the numbers of EROs in the range $K=17-20$.  See also e.g.\\ discussion by Martini (\\cite{Ma01}) and Firth et al.\\ (\\cite{Fi02}). It is important to realize that a drastic increase in the fraction of dusty EROs does not make the hierarchical models fit the ERO counts any better: they predict too few EROs of all kinds. On the other hand large fractions of dusty EROs would have to be subtracted from PLE models including only ellipticals, making all the Daddi et al.\\ PLE models used above to overpredict the counts by factors of $1.5-3$.  It is concluded that the PLE models do give remarkably good fits to the brightest ERO counts.  Formation redshifts around $z\\sim3$ are favoured -- however, by altering the details of the models, formation redshifts between $z=2-10$ are also consistent.  Moreover, a range in the formation era of ellipticals is suggested by the range in ERO colours."
    },
    "0404/astro-ph0404296_arXiv.txt": {
        "abstract": "We investigate the relationship between emission-line strength and continuum luminosity in the best-studied nearby Seyfert~1 galaxy NGC~5548. Our analysis of 13~years of ground-based optical monitoring data reveals significant year-to-year variations in the observed H$\\beta$ emission-line response in this source. More specifically, we confirm the result of Gilbert and Peterson (2003) of a non-linear relationship between the continuum and H$\\beta$ emission-line fluxes. Furthermore, we show that the slope of this relation is not constant, but rather decreases as the continuum flux increases. Both effects are consistent with photoionisation model predictions of a luminosity-dependent response in this line. ",
        "introduction": "Observations of correlated continuum and broad emission-line variations in Active Galactic Nuclei (AGN) provide powerful diagnostics of the structure and physical conditions within the spatially unresolved broad emission-line region (BLR) and of the origin and shape of the unobservable ionising continuum.  A well-established correlation is the observed decrease in broad emission-line equivalent width with increasing continuum level for AGN spanning a broad range in continuum luminosity; this is the so-called global ``Baldwin effect'' (Baldwin 1977; Osmer, Porter and Green 1994).  Originally observed in the strongest UV resonance lines (e.g. CIV$\\lambda 1549$ and Ly$\\alpha$), a global Baldwin effect has now been found for most of the strong UV emission-lines. Interestingly, evidence for a similar effect in the broad optical hydrogen recombination lines remains weak.  In a recent study of composite quasar spectra spanning a broad range in continuum luminosity and redshift, Dietrich et~al.\\ (2003) found no evidence for evolution of the Baldwin effect with cosmic time. Instead they suggested that the effect is a consequence of luminosity-dependent spectral variations, in the sense that the ionising continuum becomes softer for higher continuum luminosities. This agrees well with the theoretical work by Korista et~al.\\ (1998), who showed that if the gas metallicity and continuum spectral energy distribution are related statistically to the quasar luminosity, then so will be the emission-line equivalent widths, in a manner described by the Baldwin effect.  Formally, the relationship between the continuum luminosity ($L_{cont}$) and broad emission-line luminosity ($L_{line}$) may be represented by a single powerlaw, such that  \\begin{equation} L_{line} \\propto L_{cont}^{\\alpha} \\, . \\end{equation}  \\noindent In terms of the line equivalent width, the Baldwin relation is then given by  \\begin{equation} EW_{line} \\propto L_{cont}^{\\beta} \\, , \\end{equation}  \\noindent where $\\beta=\\alpha-1$. The measured values of $\\alpha$ are typically $<1$, for example, $\\alpha\\approx0.83$ for CIV and $\\alpha\\approx0.88$ for Ly$\\alpha$ (see e.g. Kinney, Rivolo \\& Koratkar 1990, Pogge \\& Peterson 1992), with corresponding slopes for the Baldwin relation of $\\beta=-0.17$, and $\\beta=-0.12$ respectively.  \\begin{figure*} \\epsfig{width=6.8in,angle=0,file=mgoad_fig1.ps} \\caption{(Upper panel)$F(\\lambda5100\\AA)$ continuum light curve (observed frame) for the full 13~yr AGNWatch monitoring campaign of NGC~5548. (Lower panel) The corresponding H$\\beta$ emission-line light curve.} \\end{figure*}  \\begin{figure*} \\epsfig{width=3.4in,angle=0,file=mgoad_fig2.ps} \\caption{Computed structure functions for the continuum (upper curve) and H$\\beta$ light curves. Dashed lines indicate the noise level in each of the light curves.} \\end{figure*}  Superposed on the global Baldwin relation, is a second relation whose slope reflects the emission-line response to variations in the ionising continuum within an individual source.  Now commonly referred to as the ``intrinsic Baldwin effect'', this phenomenon accounts for at least some of the scatter seen in the global Baldwin relation. However, the intrinsic Baldwin relation also shows considerable scatter, whose origin is largely attributable to continuum--emission-line time-delay effects (reverberation) within the physically extended broad emission-line region (see e.g. Krolik et~al.~1991, Pogge and Peterson 1992, Peterson et~al.~2002).  Formally, the relationship between the continuum flux $F_{\\rm cont}$ and emission-line flux $F_{line}$ within a single source can also be represented by a single powerlaw, such that  \\begin{equation} F_{line} \\propto F_{cont}^{\\alpha} \\, , \\end{equation}  \\noindent where in this case $\\alpha$ is a measure of the instantaneous emission-line response to changes in the ionising continuum flux, the {\\em responsivity\\/} of the gas. In terms of the line equivalent width, the intrinsic Baldwin relation is then given by  \\begin{equation} EW_{line} \\propto F_{cont}^{\\beta} \\, . \\end{equation}  \\noindent where again $\\beta=\\alpha-1$.  While a global Baldwin relation has yet to be seen in the optical recombination lines, there is no physical reason why an individual source should not display an intrinsic Baldwin effect in these lines. That is to say, the broad range in physical conditions extant within the BLR are such that case~B recombination, for which $\\alpha\\equiv 1$, does not apply.  Indeed, Gilbert \\& Peterson (2003) recently reported just such an effect for the broad H$\\beta$ emission-line in a comprehensive study of 10 years of ground-based optical monitoring data on the nearby Seyfert~1 galaxy NGC~5548, taken as part of the AGNWatch collaboration. Interestingly, Gilbert \\& Peterson noted (their Figure~4) that during a low continuum flux state (year 4 of the monitoring campaign), the relationship between the continuum and H$\\beta$ emission-line flux, appeared to change.  However, since there was only {\\em one\\/} low flux state, they were unable to determine whether this change was a result of flux-dependent effects (e.g. photoionisation) or time-dependent effects (e.g. a renormalization of the intrinsic relation due to a change in composition, and/or distribution of the line-emitting gas).  While variations in the emission-line response with continuum state have long been predicted by photoionisation calculations (e.g. O'Brien et~al. 1995; Goad 1995; Korista \\& Goad 2004), they have so far never been reported in the observations. Here, we take a detailed look at all 13~years of optical continuum and emission-line data available for NGC~5548, including two new low-flux states. Using this unique dataset we have found the first observational evidence for luminosity-dependent variations in the emission-line response.  These luminosity-dependent variations are not only apparent in the whole dataset (reinforced by the addition of new low continuum flux data from years 12 and 13), but are even visible within individual observing seasons, being particularly prominent during those seasons displaying both high and low continuum flux states (e.g. year 4, see \\S3).  Since the observed variations can occur on timescales of less than one year, they most likely represent a response to luminosity variations rather than to structural changes in the emission-line region. ",
        "conclusions": "Our analysis of $\\sim$13 years of optical continuum and H$\\beta$ emission-line data for NGC~5548 shows that the H$\\beta$ emission-line responsivity and hence slope of the intrinsic Baldwin effect $\\beta$ displays significant variation on timescales $\\sim1$~yr. The line responsivity is generally anti-correlated with continuum level such that in high continuum states the emission-line response is weaker on average. Conversely in low continuum states the responsivity is stronger than on average. This behaviour is consistent with predictions of photoionisation models."
    },
    "0404/astro-ph0404405_arXiv.txt": {
        "abstract": "The cluster soft excess emission indicates the presence of large amounts of warm gas (T$\\sim10^6$ K) in the neighborhood of galaxy clusters. Among the clusters that display this phenomenon is the Coma cluster, the nearest rich galaxy cluster. The excess emission is more prominent at the cluster's outskirts than at its center. Detailed studies of its large-scale emission -- up to $\\sim$2.6 Mpc from the cluster's center -- reveal that these warm baryons are as massive as, or possibly more massive than, the well-known hot intra-cluster medium (T$\\sim10^8$ K). A possible interpretation of the excess emission from the Coma cluster is radiation from low-density filaments located in the neighborhood of the cluster. In this case, the filaments would extend for much larger distances, or feature higher density, than predicted by current cosmological simulations. ",
        "introduction": "Clusters of galaxies are strong emitters of X-rays, which originate from a hot and diffuse intra-cluster medium (ICM) at temperatures of a few $\\times 10^{7}$ K. The soft X-ray band below $\\sim$ 1 keV often feature the `soft excess' emission phenomenon, i.e., radiation in excess of that expected from the hot ICM. The excess emission could originate from Inverse-Compton scattering of cosmic microwave background (CMB) photons against a population of relativistic electrons in the intra-cluster medium (Hwang 1997, Sarazin and Lieu 1998, Ensslin and Biermann 1998, Lieu et al. 1999). Alternatively, warm gas at T$\\sim 10^6$ K could be responsible for the soft emission (e.g., Lieu et al. 1996a,b, Nevalainen et al. 2003, Kaastra et al. 2003, Bonamente, Joy and Lieu 2003). Warm gas may reside inside the cluster, or in very diffuse filamentary structures outside the cluster, as seen in large scale hydrodynamic simulations (e.g., Cen and Ostriker 1999, Dav\\'{e} et al. 2001, Cen et al 2001). The warm gas scenario appears to be favored by the current X-ray data (e.g., Bonamente, Lieu and Mittaz 2001a, Buote 2001, Kaastra et al. 2003) ",
        "conclusions": "Detection of soft X-ray excess emission from galaxy clusters is commonplace. The radiation is more prominent at the cluster outskirts, and it indicates that clusters may be a significant reservoir of `warm' baryons. Precise mass estimates for the `soft' component require knowledge of the exact location of the emitter. If the gas is located in filaments outside the clusters, the detection of soft excess in the Coma cluster indicates that the `warm' gas is more massive than the hot ICM. Filaments of considerably higher density than predicted by the current cosmological simulations are needed to explain the detected excess emission from the Coma cluster."
    },
    "0404/astro-ph0404319_arXiv.txt": {
        "abstract": "By way of a complete integration of the Friedmann equations, in terms of observables, it is shown that for the cosmological constant  $\\Lambda > 0$ there exist non-flat FLRW models for which the total density parameter $\\Omega$ remains $\\sim 1$ throughout the entire history of the universe. Further, it is shown that in a precise quantitative sense these models are not finely tuned. When observations are brought to bear on the theory, and in particular the WMAP observations, they confirm that we live in just such a universe. The conclusion holds when the classical notion of $\\Lambda$ is extended to dark energy. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404411_arXiv.txt": {
        "abstract": "Recently, the High Resolution Fly's Eye (HiRes) experiment claims that there is no small scale anisotropy in the arrival distribution of ultra-high energy cosmic rays (UHECRs) above $E>10^{19}$ eV contrary to the Akeno Giant Air Shower Array (AGASA) observation. In this paper, we discuss the statistical significance of this discrepancy between the two experiments. We calculate arrival distribution of UHECRs above $10^{19}$ eV predicted by the source models constructed using the Optical Redshift Survey galaxy sample. We apply the new method developed by us for calculating arrival distribution in the presence of the galactic magnetic field. The great advantage of this method is that it enables us to calculate UHECR arrival distribution with lower energy ($\\sim 10^{19}$ eV) than previous studies within reasonable time by following only the trajectories of UHECRs actually reaching the earth. It has been realized that the small scale anisotropy observed by the AGASA can be explained with the source number density $\\sim 10^{-5 \\sim -6}$ Mpc$^{-3}$ assuming weak extragalactic magnetic field ($B \\le 1$ nG). We find that the predicted small scale anisotropy for this source number density is also consistent with the current HiRes data. We thus conclude that the statement by the HiRes experiment that they do not find small scale anisotropy in UHECR arrival distribution is not statistically significant at present. We also show future prospect of determining the source number density with increasing amount of observed data. ",
        "introduction": "\\label{intro}   The small scale anisotropy in the observed arrival distribution of ultra-high energy cosmic rays (UHECRs) is key in identifying their sources which will tell us a great deal about acceleration mechanisms, composition of UHECRs, and so on. The Akeno Giant Air Shower Array (AGASA) observation reveals the existence of the small scale clusterings in the isotropic arrival distribution of UHECRs \\citep{takeda99,takeda01}. The current AGASA data set of 57 events above 4 $\\times 10^{19}$ eV contains four doublets and one triplet within a separation angle of 2.5$^\\circ$. Chance probability to observe such clusters under an isotropic distribution is only about 1 $\\%$ \\citep*{hayashida00,takeda01}.    On the other hand, the recent claim by the High Resolution Fly's Eye \\citep*[HiRes; ][]{wilkinson99} experiment is that there is no small scale anisotropy in the UHECR arrival distribution above $E>10^{19}$ eV (162 events) which is observed by a stereo air fluorescence detector \\citep{finley03}. The discrepancy between the AGASA and HiRes involves a number of consideration. Among them are the different numbers of the events, the possibility of the measured energies being shifted between the two experiments, \\citep*{marco03} and the different angular resolutions. The purpose of this paper is to discuss the first possibility. That is, we consider whether the discrepancy between the two experiments is statistically significant or not at present.    There is furthermore the contradiction on the presence or absence of the GZK cutoff \\citep*{greisen66,zatsepin66} in the cosmic-ray energy spectrum due to photopion production with the photons of the cosmic microwave background (CMB). The HiRes spectrum shows the GZK cutoff \\citep*{abu02}, but the AGASA spectrum does not \\citep*{takeda98}. The statistical significance of this discrepancy is already discussed in \\cite{marco03} and shown to be low at present. This issue is left for future investigation by new large-aperture detectors under development, such as South and North Auger project \\citep*{capelle98}, the EUSO \\citep*{euso92}, and the OWL \\citep*{owl00} experiments.    When we consider the statistical significance of the small scale anisotropy, we need to compare the observations with numerical calculations predicted by some kind of source distributions. In this paper, we use the Optical Redshift Survey \\citep*[ORS; ][]{santiago95} galaxy sample to construct realistic source models of UHECRs. This sample has been adopted in our previous studies \\citep*{yoshiguchi03a,yoshiguchi03c,yoshiguchi03d}. It has been realized that the small scale anisotropy well reflect the number density of UHECR sources \\citep*{yoshiguchi03a,yoshiguchi03c,blasi04}. And also, it is shown that the small scale anisotropy above $E=4\\times 10^{19}$ eV observed by the AGASA can be explained with the source number density $\\sim 10^{-5 \\sim -6}$ assuming weak extragalactic magnetic field ($B \\le 1$ nG, EGMF) and that observed UHE particles are protons \\citep*{yoshiguchi03a,yoshiguchi03c,blasi04}. We thus take the source number density as a parameter of our source model, and discuss the statistical significance of the small scale anisotropy by considering to what extent we can determine the source number density from the comparison of our model predictions with the HiRes observation.    As mentioned above, the claim by the HiRes experiment on the small scale anisotropy is based on the data of UHECR arrival directions above $10^{19}$ eV. UHE protons at such energies are significantly deflected by the galactic magnetic field (GMF) by a few $- \\sim 10$ degrees. In order to accurately calculate the expected UHECR arrival distribution and compare with the observations, this effect should be taken into account.    In our recent work \\citep*{yoshiguchi03d}, we presented a new method for calculating UHECR arrival distribution which can be applied to several source location scenarios including modifications by the GMF. Brief explanation of our method is as follows. We numerically calculate the propagation of anti-protons from the earth toward the outside of the Galaxy (we set a sphere centered around the Galactic center with radius $r_{\\rm src}=$40 kpc as the boundary condition), considering the deflections due to the GMF. The anti-protons are ejected isotropically from the earth. By this calculation, we can obtain the trajectories and the sky map of position of anti-protons that have reached to the boundary at radius $r_{\\rm src}=$ 40 kpc.    Next, we regard the obtained trajectories as the ones of PROTONs from the outside of the galaxy toward the earth. Also, we regard the obtained sky map of position of anti-proton at the boundary as relative probability distribution (per steradian) for PROTONs to be able to reach to the earth for the case in which the flux of the UHE protons from the extra-galactic region is isotropic (in this study, this flux corresponds to the one at the boundary $r_{\\rm src}=40$ kpc). The validity of this treatment is supported by the Liouville's theorem. When the flux of the UHE protons at the boundary is anisotropic (e.g., the source distribution is not isotropic), this effect can be included by multiplying this effect (that is, by multiplying the probability density of arrival direction of UHE protons from the extra-galactic region at the boundary) to the obtained relative probability density distribution mentioned above.    By adopting this new method, we can consider only the trajectories of protons which arrive to the earth, which, of course, helps us to save the CPU time efficiently and makes calculation of propagation of CRs even with low energies ($\\sim 10^{19}$ eV) possible within a reasonable time.    With this method, we calculate the arrival distribution of UHE protons above $10^{19}$ eV for our source models. We also consider the energy loss processes when UHE protons propagate in the intergalactic space. Using the two point correlation function as statistical quantities for the small scale anisotropy, we compare our model prediction with the HiRes observation. We then discuss to what extent we can determine the source number density of UHECRs by the current HiRes data. We also briefly discuss future prospect of determining the source number density with the event number expected by future experiments.    In section~\\ref{gmf}, we introduce the GMF model. We explain the method for calculating UHECR arrival distribution in section~\\ref{method}. Results are shown in section~\\ref{result}. In section~\\ref{summary}, we summarize the main results. ",
        "conclusions": "\\label{summary}   In this paper, we considered the statistical significance of the discrepancy between the HiRes and the AGASA experiment on the small scale anisotropy in arrival distribution of UHECRs. We calculated arrival distribution of UHECRs above $10^{19}$ eV predicted by our source models constructed using the ORS galaxy sample. We applied the new method developed by us \\citep{yoshiguchi03d} for calculating arrival distribution with the modifications by the GMF. The great advantage of this method is that it enables us to calculate UHECR arrival distribution with lower energy ($\\sim 10^{19}$ eV) than previous studies by following only the trajectories actually reaching the earth.    It has been realized that the small scale anisotropy in the UHECR arrival distribution well reflects the number density of UHECR sources \\citep*{yoshiguchi03a,yoshiguchi03c,blasi04}. And also, it is shown that the small scale anisotropy above $E=4\\times 10^{19}$ eV observed by the AGASA can be explained with the source number density $\\sim 10^{-5 \\sim -6}$ assuming weak extragalactic magnetic field ($B \\le 1$ nG, EGMF) and that observed UHE particles are protons \\citep*{yoshiguchi03a,yoshiguchi03c,blasi04}. We thus took the source number density as a parameter of our source model, and discussed the statistical significance of the small scale anisotropy by considering to what extent we can determine the source number density from the comparison of the model predictions with the HiRes observation.   When we consider to what extent we can determine the source number density of UHECRs, we have to quantify deviations of predictions by our source models from the observation. And also, even if we specify the source number density of our model, source distribution itself can not be determined because of randomness when we select galaxies from the ORS sample. Therefore, we first evaluated the statistical significance of deviation ($\\chi^2$) of the two point correlation function predicted by a source distribution for a given number density from that expected for uniform source distribution. We then repeated such calculations for a number of realizations of the source selection. This gave the probability distributions of $\\chi^2$ for various number densities of UHECR sources.    The results for $N_{\\rm bin} = 1$ and $5$ show that differences between probability distributions of $\\chi^2$ for different source number densities are clearer when $N_{\\rm bin} = 1$. This is related to the fact that the two point correlation function is most sensitive to the source number density at the smallest angular bin. We thus focused our attention to the result of $N_{\\rm bin} = 1$. As mentioned above, we know that the small scale anisotropy observed by the AGASA can be explained with the number density $\\sim 10^{-6\\sim -5}$ Mpc$^{-3}$. It is found that the small scale anisotropy predicted by this source number density is also consistent with the prediction of isotropic source distribution when the event number is equal to the HiRes data ($\\sim 150$). Note that the data number of AGASA is about $1000$ above $10^{19}$ eV. The possibility that the number density is about $\\sim 10^{-6\\sim -5}$ Mpc$^{-3}$ can not be ruled out, although the HiRes result seems to be in agreement with the isotropic source distribution. We thus concluded that the statement by the HiRes experiment that they do not find small scale anisotropy in UHECR arrival distribution is not statistically significant at present.    We also discussed the future prospects of determining the UHECR source number density. As is evident from figure \\ref{histo}, difference between the probability distribution of $\\chi^2$ for different source number densities become clear when the number of data increases. We will be able to determine the source number density by the future observations of the small scale anisotropy. In order to discuss to what accuracy we can know the source number density, we need to calculate the probability distribution of $\\chi^2$ more precisely with a large number of simulations of the source selection. This issue deserves further investigation.   Finally, we mention the assumptions made in this paper. We neglected the effects of the extragalactic magnetic field and the random component of the GMF. If we take these effect into account, the small scale anisotropy obtained by numerical calculations become less obvious. In this case, our statement that the small scale anisotropy for $\\sim 10^{-5\\sim -6}$ Mpc$^{-3}$ is also consistent with the prediction of isotropic sources when the event number is equal to the HiRes data ($\\sim 150$) remains to be valid. Hence, the assumptions are not so important for our conclusion.  \\vspace{0.3cm} {\\bf{Note added:}} While we are finishing this study, a paper by the HiRes collaboration \\citep{abbasi04} appears on the net, where they present the results of a search for small scale anisotropy in the observed arrival distribution of UHECRs. They find no small scale anisotropy in 271 events above $10^{19}$ eV, which is slightly larger than the event number considered in the present study (164). However, this increase of the event number will not affect our conclusion very much, as seen from figure~\\ref{histo}."
    },
    "0404/astro-ph0404282_arXiv.txt": {
        "abstract": "We examine cosmic microwave background constraints on variations in the recombination history of the universe. We use a very general extension to the standard model of recombination characterized by two parameters, $a$ and $b$, which multiply the overall rates of recombination and ionization and the binding energies of hydrogen respectively.  We utilize WMAP temperature and TE cross-correlation data, along with temperature data at smaller scales from ACBAR and CBI, to place constraints on these parameters.  The range of recombination histories which gives the best fit to the data is $-0.6 < \\loga < 0.5$ and $0.85 < b < 1.15$, forming a diagonal region from low-$a$/high-$b$ to high-$a$/low-$b$.  We find $z_{rec} = 1055 \\pm 25$ and $\\Delta z = 198 \\pm 5$ at $68\\percent$ confidence for the range of recombination models which are best able to reproduce the observed CMB power spectra.  Standard recombination provides an acceptable fit to the data; while there is still room for non-standard recombination, the allowed variation is small. ",
        "introduction": "% \\label{sec:Introduction}  The cosmic microwave background (CMB) provides a snapshot of the early universe.  The CMB radiation last interacted directly (i.e., non-gravitationally) with matter during the epoch of recombination, the period when the universe became cool enough for atomic nuclei to capture and hold electrons for the first time.  Therefore, the CMB provides a sensitive test to any variations in the standard model of recombination.  A number of authors have looked at CMB constraints on variations in standard recombination \\cite{Peebles:2000, Landau:2001, DoroshkevichNaselsky:2002, Bean:2003}. Here we reexamine this question in the light of recent, high-quality data from the WMAP satellite \\citep{Bennett:2003}, ACBAR \\citep{Pearson:2002}, and CBI \\citep{Kuo:2002}, following the parametrization presented in Ref. \\cite{Hannestad:2000}.  The WMAP data represents a significant improvement over data used in most previous calculations, and significantly tightens the resulting constraints.  In addition, we include the analysis of polarization data in the search for constraints on recombination.  The only previous analysis to use the recent WMAP data has been given by Bean, et al. \\cite{Bean:2003}, who examined changes in the ionization history induced by Lyman-$\\alpha$ photons and ionizing photons.  Our approach is somewhat different; we have attempted a more general, model-independent investigation, based on the discussion in Ref. \\cite{Hannestad:2000}.  Hence, the constraints we derive are not tied to any particular model for altering recombination, but are potentially more general.  Further, we have used additional data from ACBAR and CBI. In addition to using the TE cross-correlation data, we investigate the ability of the E-mode polarization power spectrum to further constrain recombination.  In Section \\ref{sec:Methods} we describe a parameterization of the epoch of recombination and the details of how we obtain constraints on the parameters. In Section \\ref{sec:Results} we display the resulting constraints on recombination and discuss how the polarization power spectrum is affected by recombination. In Section \\ref{sec:Conclusions} we present our conclusions. ",
        "conclusions": "% \\label{sec:Conclusions}  We have examined the new data from the WMAP mission, along with data from higher $\\ell$ experiments, ACBAR and CBI, in an attempt to constrain recombination. We have included analyses of both the temperature power spectrum $\\Cl^{TT}$ and the TE cross-correlation power spectrum $\\Cl^{TE}$.   In the general parameterization that we used, the $1\\sigma$ allowed region in \\hbox{$a$-$b$} space was significantly contracted compared to earlier attempts with less precise data (\\cite{Hannestad:2000}).   We find that the allowed region falls within $-0.6 < \\loga < 0.5$ and $0.85 < b < 1.15$. The average width and position of the surface of last scattering for recombination histories falling within the $1\\sigma$ region are $\\zrec = 1055 \\pm 25$  and $\\Delta z = 198 \\pm 5$. Note that the WMAP team estimated values for $\\zrec$ and $\\Delta z$, namely \\cite{Spergel:2003} $\\zrec^\\text{WMAP} = 1088^{+1}_{-2}$ and $\\Delta z^\\text{WMAP} = 194 \\pm 2$. Our values for $\\zrec$ and $\\Delta z$ do not mean the same thing as those reported by the WMAP team.  In the latter case, standard recombination was assumed, and these values simply give the best-fit recombination history obtained by allowing the cosmological parameters to vary.  In our case, we allow variations in the recombination history itself, so we expect a larger range of values for $\\zrec$ and $\\Delta z$.  In fact, our values for $\\zrec$  and $\\Delta z$ are quite similar to the WMAP values; this is a reflection of the fact that standard recombination is an excellent fit to the data and is approximately in the center of the distribution of the allowed variant recombination histories.  The changes in $a$ and $b$ are largely compensated by changes in other parameters.  The largest effects are that an increase in $a$ is matched by corresponding increases in $\\omega_b$ and $h$ and a corresponding decrease in $n_s$, while an increase in $b$ is compensated by increases in $\\omega_b$ and $h$. Tighter constraints on these parameters, especially $h$, would allow us to narrow the range of allowed recombination histories.  The result of these competing effects --- changes in $a$ and $b$ being partially compensated for by changes in the free cosmological parameters --- is that the anisotropy power spectra from the $1\\sigma$ region all share very similar features.  In general, the positions of the peaks are strongly constrained, while the amplitudes of the peaks are less so.  The range in E-mode polarization power spectra corresponding to recombination histories from this region is very similar to the range in temperature and TE cross-correlation power spectra, and is unlikely to be useful in breaking degeneracies in \\hbox{$a$-$b$} parameter space.  We find that the range of alternate recombination histories allowed by the WMAP, ACBAR, and CBI temperature and WMAP TE cross-correlation data is very small compared to that allowed by earlier data, and is approximately centered on standard recombination.  While there is still some room for non-standard recombination histories, the constraints have narrowed considerably with this new data.  Standard recombination fits the data extremely well, and there is no compelling evidence from this study to suggest that a non-standard recombination is preferable."
    },
    "0404/astro-ph0404557_arXiv.txt": {
        "abstract": "We present mid-infrared photometry of three very young brown dwarfs located in the $\\rho$ Ophiuchi star-forming region -- GY5, GY11 and GY310 -- obtained with the Subaru 8-meter telescope. All three sources were detected at 8.6 and 11.7$\\mu$m, confirming the presence of significant mid-infrared excess arising from optically thick dusty disks.  The spectral energy distributions of both GY310 and GY11 exhibit strong evidence of flared disks; flat disks can be ruled out for these two brown dwarfs.  The data for GY5 show large scatter, and are marginally consistent with both flared and flat configurations.  Inner holes a few substellar radii in size are indicated in all three cases (and especially in GY11), in agreement with magnetospheric accretion models.  Finally, our 9.7$\\mu$m flux for GY310 implies silicate emission from small grains on the disk surface (though the data do not completely preclude larger grains with no silicate feature). Our results demonstrate that disks around young substellar objects are analogous to those girdling classical T Tauri stars, and exhibit a similar range of disk geometries and dust properties. ",
        "introduction": "Now that a large number of brown dwarfs are known, some with inferred masses close to those of extrasolar giant planets, their origin has become a subject of widespread interest and fundamental importance.  One key recent finding is that substellar objects harbor accretion disks and undergo a T Tauri phase similar to that of low-mass stars, possibly suggesting a common formation mechanism.  The evidence for disks around brown dwarfs includes near- and mid-infrared excess (e.g., Tamura et al. 1998; Natta \\& Testi 2001; Natta et al. 2002; Jayawardhana et al. 2003), spectroscopic signatures of accretion (e.g., Jayawardhana, Mohanty \\& Basri 2002, 2003; Muzerolle et al. 2003) and millimeter-wave emission (Klein et al. 2003). In some cases, these disks appear to last for up to 10 million years (Mohanty, Jayawardhana \\& Barrado y Navascu\\'es 2003), a timescale comparable to that of disks surrounding low-mass stars.  Moreover, there are now indications of jets/outflows from objects close to or below the substellar limit (Fern\\'andez and Comer\\'on 2001; Barrado y Navascu\\'es, Mohanty \\& Jayawardhana 2004), further strengthening the T Tauri analogy.  Given the much lower mass and intrinsic luminosity of brown dwarfs relative to T Tauri stars, the preceding discoveries raise intriguing questions about how the detailed properties of brown dwarf disks compare to those of T Tauri disks.  Observations in the mid-infrared, where optically thick disk emission dominates over the substellar photosphere, can provide important constraints on the disk models.  Here we report ground-based photometry in the 8.6, 9.7 and 11.7$\\mu$m filters, spanning the 10$\\mu$m silicate feature, of three young brown dwarfs in the nearby $\\rho$ Ophiuchi region (age$\\lesssim$1 Myr, d$\\approx$150 pc), and use these data to explore their disk characteristics in comparison with models. ",
        "conclusions": "SED model fits to our data for GY310, 11 and 5 are illustrated in Fig. 1. We plot both our new Subaru data at 8.6, 9.7 and 11.7$\\mu$m as well as additional flux measurements at shorter and longer wavelengths, mainly from ISO.  The flux-calibrated near-infrared (0.85--2.45$\\mu$m) stellar spectra from NTC02 are also shown.  All the observed fluxes have been corrected for reddening following NTC02.  For each object, we plot two flared disk models, one with a distribution of small grains and the other with larger grains; a flat disk (blackbody) model; and a model stellar spectrum (Allard et al. 2001; see NTC02) depicting the contribution of the underlying stellar photosphere.  The models incorporate inner holes of various sizes.  The effect of an inner hole in a flared disk is twofold.  On the one hand, it depresses the near-IR emission, by getting rid of the innermost hot dust; on the other hand, it increases the fraction of the stellar surface ``seen'' by the disk (from 1/2 to 1, in the extreme cases $R_i=R_\\ast$ (no hole) and $R_i\\gg R_\\ast$ respectively), and therefore the emission at longer wavelengths.  All three targets exhibit considerable mid-IR excess above the stellar continuum: while their earlier ISO detections left open the possibility of source confusion due to ISO's coarse spatial resolution, the Subaru detections confirm that these objects harbor circum-(sub)stellar disks. The salient results from our present work are as follows: {\\it (1)} both GY310 and GY11 clearly possess flared disks - flat disks can be confidently ruled out in their case, {\\it (2)} the 9.7$\\mu$m flux for GY310 indicates possible silicate emission from small surface grains (though larger grains with no silicate emission cannot be completely excluded), and {\\it (3)} inner holes of size $\\sim$ 3--7 stellar radii seem necessary to fit the data, analogous to the situation in classical T Tauri stars.  We now briefly discuss each source.  {\\it GY 310}: Our best fit for this object requires a flared disk, tilted 60$^{\\circ}$ from face-on and with an inner-hole radius of $R_i$ $\\approx$ 7$R_{\\ast}$.  Inner holes significantly smaller than this produce too much emission at 6.7$\\mu$m, while smaller inclinations overestimate the mid-IR flux.  Notice that a flat disk model is completely inconsistent with the data longward of 6.7$\\mu$m; the predicted fluxes are simply too low. Moreover, all the fluxes from 6.7 to 14.3 $\\mu$m agree remarkably well with the flared disk SED that includes 10$\\mu$m silicate emission from small surface grains.  A silicate feature is routinely observed in disks around T Tauri stars (e.g., Cohen \\& Witteborn 1985; Natta, Meyer \\& Beckwith 2000). If our detection is confirmed, it would be the first evidence of such emission from a brown dwarf disk. Nevertheless, we caution that our present data cannot unequivocally rule out a flared disk with larger grains and no silicate feature: our 9.7$\\mu$m flux remains marginally consistent with the latter scenario.  Measurements with smaller uncertainties, such as Spitzer low-resolution spectra, can settle this issue conclusively.  {\\it GY 11}: Our best-fit model for this brown dwarf is a flared disk seen face-on, with an inner hole of size $R_i$ $\\approx$ 3$R_{\\ast}$.  Much smaller inner holes produce unacceptably low mid-IR fluxes.  Similarly, a flat disk produces emission wholly inadequate to match the steeply rising mid-IR SED.  However, in the absence of 9.7$\\mu$m data, we are not sensitive to the presence of 10$\\mu$m silicate emission, and thus cannot distinguish between flared disks with small surface grains and those with somewhat larger particles.  We note that the \\lbol we adopt for GY11 (0.012 \\lsun), to obtain an adequate SED fit, is moderately higher than derived by NTC02 (0.008 \\lsun).  The NTC02 value produces mid-IR fluxes somewhat lower than observed (fit not shown); since we already maximize the mid-IR emission by adopting a face-on disk and invoking an inner hole$\\sim$ 3$R_{\\ast}$, higher luminosity seems the only remaining option in the context of our models.  The \\lbol adopted here is admissible within NTC02's \\av errors for this heavily reddened object (our \\lbol amounts to \\av $\\approx$8.5 instead of NTC02's 7.5$\\pm 1$).  The \\teff from synthetic spectral fits assuming our new \\av is 2500K, very close to NTC02's 2400K; the new \\lbol and \\teff also remain consistent with an age $\\lesssim$ 1 Myr. The mass now inferred from the evolutionary tracks is $\\sim$ 20 \\mj; while higher than NTC02's estimate of 10 \\mj, it still implies that GY11 is a very low-mass brown dwarf.  {\\it GY 5}:   The large scatter in the GY 5 data make any fit provisional at best.  However, if we disregard the ISO 3.6$\\mu$m flux, which is 50\\% larger than the ground-based value in the comparable L$'$ filter (from Comeron et al. 1998 and Jayawardhana et al. 2003) and assume that the true fluxes at 6.7 and 14.3$\\mu$m are at the lower end permissible by the ISO error bars, then a flat disk inclined at $\\sim$ 65$^{\\circ}$ to face-on, with an inner hole $R_i$ $\\sim$ 5$R_{\\ast}$, provides a plausible fit (especially since our 9.7$\\mu$m Subaru estimate is only a flux upper limit). However, we cannot discount a flared disk seen at large inclination.  The disk parameters presented here for the three $\\rho$ Oph sources are not unique: there are plenty of trade-offs between factors such as disk inclination, inner-hole size, grain albedos and sizes, and degree of dust settling.  For example, our models adopt only the two extremes of grain settling: well-mixed gas and dust leading to full disk flaring and complete settling resulting in flat disks.  If there is some, but not complete, dust settling in GY310 for instance, its moderate mid-IR excess may be fit by a disk model with a lower inclination.  Nonetheless, the SEDs we observe strongly suggest certain qualitative features that are relatively independent of the particular disk parameters adopted.  First, the presence of flared disks in GY310 and 11 is indubitable: no flat disk, whatever its inner-hole radius or inclination, can reproduce the observed increase in their flux from 6.7 to 14.3 $\\mu$m. Apai et al. (2002) carried out observations similar to ours for a single young object at the stellar/substellar edge, ChaH$\\alpha$2, and found that a flat disk with no silicate feature provides the best fit to its SED.  Our results indicate that, while such disks might indeed be present in some cases (e.g., perhaps GY5), this does not signify a fundamental change in disk properties upon crossing from the stellar to the substellar regime: brown dwarfs can have flared disks too, as well as (tentatively) silicate emission, and simply present a range of disk geometries just as T Tauri stars do (e.g., Miyake \\& Nakagawa 1995).  Second, inner holes of a few substellar radii may be common in brown dwarf disks - our best fits to all three objects require such a configuration.  Especially for GY11, it is hard to find an alternative; even assuming a face-on disk and a (sub)stellar luminosity higher than previously estimated, some mechanism is required to further raise the mid-IR fluxes to match the data, and a central hole seems the only viable possibility (see also Liu et al. 2003).  Indeed, magnetospheric accretion models predict the formation of inner holes of just these sizes in both classical T Tauri stars and accreting brown dwarfs, due to disk truncation by interaction with the (sub)stellar magnetic field (Shu et al. 1994; Muzerolle et al. 2003).  In summary, therefore, our study suggests that brown dwarf disk properties are qualitatively similar to those of higher mass T Tauri stars.  A more detailed understanding of grain characteristics, dust settling and geometry in brown dwarf disks will be possible with mid-IR Spitzer observations."
    },
    "0404/astro-ph0404231_arXiv.txt": {
        "abstract": "We present quantitative studies of 8 late O and early B-type supergiants in the Magellanic Clouds using far-ultraviolet $FUSE$, ultraviolet $IUE$/$HST$ and optical VLT-UVES spectroscopy. Temperatures, mass-loss rates and CNO abundances are obtained using the non-LTE, spherical, line-blanketed model atmosphere code of Hillier \\& Miller (1998).  We support recent results for lower temperatures of OB-type supergiants as a result of stellar winds and blanketing, which amounts to $\\sim$2000\\,K at B0\\,Ia.  In general, H$\\alpha$ derived mass-loss rates are consistent with UV and far-UV spectroscopy, although from consideration of the S\\,\\4 $\\lambda\\lambda$1063-1073 doublet, clumped winds are preferred over homogenous models.  AV\\,235 (B0\\,Iaw) is a notable exception, which has an unusually strong H$\\alpha$ profile that is inconsistent with the other Balmer lines and UV wind diagnostics.  We also derive CNO abundances for our sample, revealing substantial nitrogen enrichment, with carbon and oxygen depletion. Our results are supported by comparison with the Galactic supergiant HD\\,2905 (BC0.7\\,Ia) for which near-solar CNO abundances are obtained.  This bolsters previous suggestions that ``normal'' OB-type supergiants exhibit atmospheric compositions indicative of partical CNO processing. ",
        "introduction": "\\label{intro} The {\\it Far~Ultraviolet~Spectroscopic~Explorer} ($FUSE$) satellite (Moos et al. 2000) has provided an opportunity to study populations of early-type stars in the Milky Way and Magellanic Clouds in the 900-1200\\AA\\ region. Spectral atlases of $FUSE$ observations of O and early B-type stars have recently been presented by Pellerin et al. (2002) and Walborn et al. (2002). The principal scientific motivation for such observations is the comparison of stellar winds from early-type stars in different metallicity ($Z$) environments. The predicted theoretical dependence of radiatively driven winds from massive stars is well documented, such that \\mdot(Z)~$\\propto Z^{0.5-0.7}$ (Kudritzki, Pauldrach \\& Puls, 1987; Vink, de Koter \\& Lamers, 2001); observationally this has yet to be firmly established. The precise sensitivity of mass-loss to metallicity for early-type stars is keenly sought as it is a necessary ingredient for stellar evolution models (and in turn evolutionary synthesis models).  Towards this goal, we initiated a program to obtain high-quality optical observations with the UV-Visual Echelle Spectograph (UVES) at the Very Large Telescope (VLT) for the Magellanic Cloud targets in our $FUSE$ Principal Investigator programmes. We have focused particularly on the Clouds to avoid contamination associated with the high column densities of H$_2$ typical of sightlines to Galactic OB-type stars (e.g., Pellerin et al. 2002) and to minimize the effect of uncertainties in the distances to the targets on the atmospheric analysis.  At the same time, so-called ``unified'' stellar atmosphere models are now available which allow for the spherical extension of supergiants and treat line blanketing by metal species in a reasonably thorough manner.  Crowther et al. (2002; hereafter Paper I) presented the first study of the present series for four Magellanic Cloud O-type supergiants, employing the model atmosphere code {\\sc cmfgen} (Hillier \\& Miller, 1998). These extreme supergiants (all were of luminosity class Ia$^+$) were specifically chosen to test previously adopted stellar temperatures for O-type stars, for which Fullerton et al. (2000) had questioned the previous plane-parallel, unblanketed results (e.g., Vacca et al. 1996). Indeed, substantially lower temperatures were derived for these extreme supergiants.  The combination of very strong stellar winds together with metal line blanketing led to temperatures 15--20$\\%$ lower than those from plane-parallel models composed solely of hydrogen and helium.  Downward revisions were also found for Galactic O-type dwarfs (Martins et al. 2002) and supergiants in the Galactic cluster Cyg OB2 (Herrero et al. 2002).  For our second paper we have therefore chosen to study O-type stars with luminosity classes in the range II to Ia$^+$, to better address the question of the stellar temperature scale.  We have also extended our sample to include early B-type supergiants, thereby providing an overlap between Paper I and the recent results of Trundle et al. (2004).  In Paper~I, CNO abundances of extreme O-type supergiants were found to differ greatly from those inferred from nebular and stellar studies in the Magellanic Clouds. Nitrogen was strongly enhanced, with carbon (and oxygen) moderately depleted, suggestive of mixing of unprocessed and CNO-processed material at their surfaces. These results were supported by a companion study by Hillier et al. (2003) which used identical techniques.  They found unprocessed CNO abundances for AV\\,69 [OC7.5\\,III((f))] and partially CNO-processed material in AV\\,83 (O7\\,Iaf$^+$).  Therefore, morphologically normal OB stars appear to show evidence for moderate levels of chemical processing. This new study allows us to examine CNO abundances in less extreme supergiants.  The issue of clumping in early-type stars was raised in Paper~I and by Hillier et al.  (2003).  In order to reconcile the optical H$\\alpha$ and far-UV P\\,{\\sc v} $\\lambda\\lambda$1118-28 wind diagnostics, either clumped winds or a reduced phosphorus abundance was required.  Massa et al. (2003) arrived at similar conclusions from a study of $FUSE$ observations of O-type stars in the LMC. The question of clumping, and hence reduced mass-loss rates for OB-type stars, is an important one, although the uncertain abundance of phosphorus in the interstellar medium (ISM) is a major limitation. In contrast, we shall suggest that S\\,{\\sc iv} $\\lambda\\lambda$1062-73 is an analogue of P\\,{\\sc v} $\\lambda\\lambda$1118-28 amongst late O and early B-type supergiants, for which the ISM elemental abundance is well known (e.g., Russell \\& Dopita 1992).  This permits firmer conclusions regarding clumping in the winds of early-type stars.  As in Paper I, the $FUSE$ data are complemented by UV spectra from the {\\it International\\,Ultraviolet\\,Explorer} ($IUE$) telescope and the {\\it Hubble Space Telescope} ($HST$) archives, together with optical data from UVES.  The observational data are presented in \\sref{obs}, followed by a description of our methods in \\sref{methods} and results for each of our targets stars in \\sref{analysis}, except AV~235 which is discussed separately in \\sref{av235}. In \\sref{abundances} we discuss the CNO abundances derived for our targets and also present an analysis of the Galactic BC-type supergiant HD\\,2905 to provide a test of our methods.  Finally, we discuss the implications of our results for published temperature calibrations and give a comparison of our observationally derived mass-loss rates with theoretical predictions. ",
        "conclusions": "\\label{end}  We have studied a sample of LMC and SMC late O and early B supergiants based on modern line-blanketed, spherical models plus extensive $FUSE$ (far-UV), $IUE$ (UV) and UVES (visual) observations. In general, we find excellent agreement between alternative optical temperature diagnostics (e.g., He and Si) with H$\\alpha$ derived mass-loss rates. Adopting homogeneous models, some UV wind features are systematically too strong.  The UV discrepancies are reduced if OB winds are clumped. P\\,{\\sc v} $\\lambda\\lambda$1118-1128 was identified as a useful probe of clumping in O supergiants in Paper~I, unless phosphorus is depleted relative to other elements in the Magellanic Clouds.  In the present study, S\\,\\4 $\\lambda\\lambda$1063-1072 appears to offer an equivalent probe in early B-type supergiants, with the additional benefit that the ISM sulphur abundance is well known (Russell \\& Dopita 1992). We conclude that winds in OB-type stars are at least moderately clumped.  This leads to lower derived mass-loss rates than otherwise, which scale with the adopted filling-factor, (with $f_\\infty = 0.1$, $\\dot{M}$ is reduced by a factor of $\\sim$3).  AV\\,235 (B0\\,Iaw) is peculiar in the current sample in that it has inconsistent optical (Balmer line), UV and far-UV wind features. All diagnostics indicate a relatively weak (and clumped) wind for AV\\,235 with the exception of H$\\alpha$.  Stellar temperatures for O and early B-type supergiants are generally 2--4\\,kK lower than recent calibrations based on unblanketed, plane-parallel models. Supergiants with extreme (typically Ia$^{+}$) winds are more greatly affected, such that there is expected to be different spectral type--temperature calibrations for OB supergiants, depending on whether H$\\alpha$ is in emission. Reduced temperatures consequently indicate lower luminosities, with impact upon ionizing fluxes (recall Figure~13 from Herrero et al. 2002), and wind-momentum luminosity calibrations.  We also investigate CNO abundances for OB supergiants. As in Paper~I, ``normal'' OB supergiants are found to have partially processed abundances, i.e., log(N/C)$\\sim$1 versus log(N/C)$\\sim-1$ for H\\,{\\sc ii} regions.  Although such large nitrogen enrichments were found in the fast-rotating evolutionary models from Maeder \\& Meynet (2001), it seems unlikely that rapid rotation is solely responsible and that mixing is perhaps more effective than previously thought at more moderate velocities.  A differential analysis of HDE\\,269896 (ON9.7\\,Ia$^{+}$) versus Sk$-66^{\\circ}$169 (O9.7\\,Ia$^{+}$) indicates even more extreme abundances -- nitrogen is further enriched at the expense of carbon and oxygen in the ON supergiant. In contrast, two OBC stars exhibit fairly normal CNO abundances; HD~2905 studied here for comparison to the Magellanic Cloud B-type supergiants, plus AV\\,69 investigated by Hillier et al. (2003) using identical techniques. Consequently, quantitative models now provide strong evidence for the sequence OBC (normal CNO) $\\longrightarrow$ OB (CNO partially processed) $\\longrightarrow$ OBN (fully CNO processed) originally suggested on from morphological arguments by Walborn (1976, 1988) and supported by studies of He-contents by Smith \\& Howarth (1994)."
    },
    "0404/astro-ph0404007_arXiv.txt": {
        "abstract": "The NRAO Very Long Baseline Array (VLBA) was used at 1.4~GHz to image 20 compact sources in the inner 1\\arcdeg\\, radius of the {\\em Spitzer\\/} VLA First-Look Survey.  Twelve sources were detected with the VLBA, with peak flux densities above 2~mJy at 9-mas resolution. The positions of these VLBA detections have been measured in the International Celestial Reference Frame (ICRF) to 1$\\sigma$ accuracies of 8~mas East-West and 4~mas North-South.  These positions can help anchor ancillary data for the First-Look Survey to the ICRF and are being reported quickly to the community. ",
        "introduction": "Wide fields observed across the electromagnetic spectrum are beginning to be surveyed at 10-mas resolution, using the techniques of Very Long Baseline Interferometry (VLBI) to study sources stronger than about 0.1~mJy at 1.4~GHz \\citep{gar01,gar04}.  The corresponding linear resolution is 90~pc or finer for a {\\em WMAP\\/} cosmology \\citep{spe03}.  This faint radio population contains sources that are dominantly energized by star-formation, plus sources -- the active galaxies -- that are dominantly energized by massive black holes. VLBI surveys of these faint sources have two primary science goals: (1) for the star-forming galaxies, seek evidence for embedded candidate active galactic nuclei or else limit contributions from such nuclei; and (2) for the active galaxies, seek evidence for the active nuclei or else limit contributions from such nuclei.  Being unaffected by obscuration, these VLBI surveys thus help constrain models for the growth of massive black holes over cosmic time.  Moreover, although these VLBI surveys must be deep to probe the faint population, their limiting sensitivities are still such that confusing a luminous type IIn supernova like SN 1988Z \\citep{wil02} for a candidate active galactic nucleus is a potential problem only for redshifts below one-tenth.  There are about 700 faint sources per square degree at 1.4~GHz in the VLA counterpart to the {\\em Spitzer\\/} First-Look Survey \\citep[FLSVLA;][]{con03}.  That VLA survey had a resolution of 5\\arcsec\\, FWHM and uniformly covered the circular area with radius 1\\arcdeg\\ that was guaranteed to be surveyed by {\\em Spitzer\\/} (Figure~1).  Given the areal densities of the FLSVLA sources, it becomes both feasible and efficient to survey the faint radio population using VLBI phase referencing in the in-beam style \\citep{wro00}.\\footnote{Available at http://www.aoc.nrao.edu/vlba/html/MEMOS/scimemos.html.}  Such VLBI surveys require substantial time, telescope, and correlator resources, and therefore must be carefully optimized for maximum scientific return.  Two key optimizing factors are the suitability of the in-beam phase calibrators used and the number of faint sources that appear adjacent to those in-beam phase calibrators.  As a step toward quantifying these factors within the 1\\arcdeg\\, radial region shown in Figure~1, the NRAO Very Long Baseline Array \\citep[VLBA;][]{nap94} was used to conduct phase-referenced observations at 1.4~GHz, in the nodding style \\citep{wro00}, of 20 candidate in-beam calibrators that are unresolved and stronger than 5~mJy in the FLSVLA.  Figure~1 shows the locations of these 20 compact FLSVLA sources, while Table~1 lists for each source the FLSVLA name, position, position error, and integrated flux density \\citep{con03}. Each compact FLSVLA source has an upper limit to its deconvolved diameter at FWHM of 2.2\\arcsec\\, or less, so given the quoted position errors it is prudent to ensure that the VLBA search region covers a diameter of about 1-2\\arcsec.  Figure~1 also shows the guaranteed region, of radius 0.25\\arcdeg, for the deeper {\\em Spitzer\\/} Verification Survey.  Although the rectangular observation regions are now known for the the First-Look and Verification Surveys, this work continues to emphasize these circular regions to optimize synergy with ancillary surveys conducted prior to those {\\em Spitzer\\/} surveys. Coincidentally, the region of diameter 0.5\\arcdeg\\, shown for the Verification Survey corresponds to the primary beam at FWHM of a VLBA antenna at 1.4~GHz, a relevant factor in planning in-beam VLBI surveys.  As the nodding calibrator for the VLBA observations was selected from the VLBA Calibrator Survey \\citep[VCS1;][]{bea02}, all compact FLSVLA sources detected with the VLBA will have their positions measured in the International Celestial Reference Frame (Extension 1) \\citep[ICRF-Ext.1;][]{ier99,ma98}.  The positions of these VLBA detections can, therefore, help anchor ancillary data for the First-Look Survey to the ICRF, and the purpose of this paper is to report these positions quickly to the community.  The VLBA observations and calibration are described in Section~2, while Section~3 describes the imaging strategies and the astrometric implications.  A multi-wavelength analysis of both the detections and non-detections from the VLBA observations will be deferred until more ancillary data are available for the First-Look Survey.  Also, future VLBI surveys and multi-wavelength studies will benefit from a new survey with the Westerbork Synthesis Radio Telescope that complements the FLSVLA survey by going deeper at 1.4~GHz over a narrower region with a resolution of 15\\arcsec\\, FWHM \\citep{mor04}. ",
        "conclusions": ""
    },
    "0404/astro-ph0404377_arXiv.txt": {
        "abstract": "The \\textit{Chandra} X-ray data of the central Virgo cluster are re-examined to reveal a triple-ring structure around the galaxy M87, reminiscent of the spectacular triple-ring pattern of the SN1987A in the Large Magellanic Cloud (LMC). In the sky plane, the two apparent smaller ellipses are roughly aligned along the M87 jets; the larger ring centers at the M87 nucleus and is likely a circle roughly perpendicular to the M87 jet. Certain similarities of these two triple-ring structures might hint at similar processes that operate in these two systems with entirely different sizes and mass scales. We suspect that a major merging event of two galaxies with nuclear supermassive black holes (SMBHs) might create such a triple-ring structure and drove acoustic and internal gravity waves far and near. The M87 jets are perhaps powered by a spinning SMBH resulting from this catastrophic merging event. ",
        "introduction": "Many similar astrophysical phenomena happen on totally different spatial and temporal scales or energy and mass scales. Examples include collimated jets from active galactic nuclei (AGNs) containing SMBHs \\citep{beg03}, from microquasars containing stellar mass black holes \\citep{mir99} and from young stellar objects \\citep{ray96}; dynamical roles of waves in spiral galaxies \\citep{lin64,fan96,lou98}, in planetary rings \\citep{gt78} and in galaxy clusters \\citep{fab03a}; similar high-temperature atmospheres of the Sun and of stellar mass black hole systems \\citep{zha00}; similar gamma-ray flashes and bursts from explosions of perhaps massive stars at cosmological distances \\citep{mes01}, from solar flares \\citep{hai91} and from the Earth's atmosphere \\citep{fis94,fen02}. We report here the detection of an X-ray triple-ring structure in the core of the Virgo galaxy cluster, most likely associated with the AGN of the galaxy M87 (NGC 4486) and its powerful jet. This triple-ring structure is reminiscent of the spectacular triple-ring pattern of the supernova 1987A in the LMC \\citep{bur95}.  M87 is an active galaxy with relativistic jets in the central Virgo cluster. The X-ray morphology and spectroscopy of the Virgo cluster have been studied previously using data from \\textit{Chandra} \\citep{you02} and \\textit{XMM-Newton} \\citep{boh01,bel01}. Combining \\textit{Chandra}, \\textit{XMM} and \\textit{ROSAT} data, \\citet[hereafter F04]{for04} studied the Virgo cluster on various scales to identify cooling flow quenching by an AGN energy input. They noticed two prolate features around the M87 jet and counter jet, and referred to them as ``cavities'' caused by plasma expansions.  In this Letter, we re-analyze the \\textit{Chandra} data and focus on the core region around M87. Two smaller ring structures are identified to encircle the ``cavities'' noticed by F04; a third larger ring centers at the M87 nucleus. We adopt the distance to M87 as $\\sim 16$ Mpc with 1\\arcsec\\ for $\\sim 78$ pc in the image \\citep{whi95}. ",
        "conclusions": "The triple-ring pattern revealed in the central Virgo cluster is the only one observed since the discovery of triple rings in SN1987A\\footnote{http://oposite.stsci.edu/pubinfo/jpeg/SN1987A\\_Rings.jpg} \\citep{bur95}. The ring sizes in the Virgo cluster are several thousand times larger than those of SN1987A. The triple-ring structure of SN1987A has been simulated in terms of interacting winds near the end of stellar evolution. According to \\citet{Tan02}, during the star's red supergiant phase, slower winds persist. Prior to the supernova explosion, the red supergiant somehow evolved into a blue supergiant that drove faster winds. Magnetohydrodynamic (MHD) interactions of fast and slow winds from a rotating star produce the triple-ring structure in SN1987A. We suspect that a triple-ring structure might be generic in a rotating system involving magnetized faster winds catching up slower winds \\citep{lou94,lou96}. Meanwhile a central source of radiation is required for the rings to shine in the X-ray band. For SN1987A, the strong UV radiation from the supernova provides the necessary illumination.  We suspect that a catastrophic merging event around the M87 nucleus might be responsible for the triple-ring structure revealed here. Merging of two SMBH systems has been proposed to drive the observed X-shaped radio lobes \\citep{mer02}. In this scenario, a `slower wind' was present before the merging begins, e.g., a mixture of galactic winds from two merging galaxies. A `faster wind' was then driven during the merging process. It has been suggested that an active AGN phase occurred in the core region of M87 at about $10^8$ years ago \\citep{kai03}; perhaps, this AGN was triggered by the merging event discussed above, when significantly higher accretion rate was provided by the merger. The merger may fuel both the radiation energy output of the AGN and the growth of the supermassive black hole in the center; huge amount of high energy radiation shines the triple rings in X-ray bands. It is also plausible that the resulting SMBH spins rapidly to power the highly collimated M87 jets \\citep{bla77}.  F04 revealed two expanding cavities around the M87 jet and counter jet, enclosed within our ring 1 and 2 respectively. By different external pressures, they inferred the far front of the cavities is mildly supersonic while the inner part, close to the nucleus, is collapsing. We suggest that these two rings are actual rings instead of projections of prolate spheres, because when two cavities collide into each other it is hard to maintain their spherical shapes in the spatially overlapping portions of the two spheroids, and consequently their projections cannot be viewed as two complete and nested rings. We further suggest that when the jets form, they push materials around and drive powerful blast waves of ring shapes.  Qualitatively similar central structures are present in the clusters Perseus \\citep{fab03a}, Hydra A \\citep{mcn01}, Centaurus A \\citep{kra02,kar02} and Virgo (F04). Physically, core activities can excite large-scale acoustic waves ($p-$modes) and internal gravity waves ($g-$modes) in a deep gravitational potential well of central galaxy cluster \\citep{lou95}. Cluster core activities can be either violent and sporadic or sustained at a level or a combination of both. Depending on the phase of acoustic wave propagation and on physical properties of core environment, we may observe various X-ray morphologies superposed onto a smooth luminous core. Beautiful acoustic rings or arcs have been revealed in Perseus \\citep{fab03a} and in Virgo on larger scales (F04). Intense core X-ray emissions tend to induce inflows of gas that perturb the dark matter via gravity. The violent relaxation proceeds in a sound-crossing timescale \\citep{lb67}. This could be a persistent source of $p-$modes and $g-$modes in a cluster core. By such wave and Laudau-damping processes, slowly inward drifting gas and dark matter are virialized to sustain X-ray losses from hot electron gas. As a cluster core may involve magnetic field of strengths up to $\\sim 30-40\\mu$G, MHD waves should also play significant roles. Based on the hardness ratio map and the energy spectra fitting, we infer only parts of ring 1 and 3 involve supersonic shock flows and are in processes of being thermalized with the surroundings. By this reasoning, other parts without apparent difference with surroundings might have been thermalized already. In summary, we hope that the triple-ring structure of M87 in the central Virgo cluster would stimulate more theoretical, simulation and observational studies of MHD flow, wave and shock interactions."
    },
    "0404/astro-ph0404141_arXiv.txt": {
        "abstract": "In order to determine the relationship between the faint X-ray and faint radio source populations, and hence to help understand the X-ray and radio emission mechanisms in those faint source populations, we have made a deep $1.4\\GHz$ Very Large Array radio survey of the $13^h$ $+38^{\\circ}$ {\\it XMM/ROSAT} X-ray Survey Area. From a combined data set of 10hours B configuration data and 14hours A configuration data, maps with $3.35''$ resolution and a noise limit of $7.5\\uJy$ were constructed.  A complete sample of 449 sources was detected within a $30'$ diameter region above a $4\\sigma$ detection limit of $30\\uJy$, at the map centre, making this one of the deepest radio surveys at this frequency. The differential source count shows a significant upturn at sub-milliJansky flux densities, similar to that seen in other deep surveys at $1.4\\GHz$ (\\eg\\space the Phoenix survey), but larger than that seen in the HDF which may have been selected to be under-dense. This upturn is well modelled by the emergence of a population of medium redshift starforming galaxies which dominate at faint flux densities. The brighter source counts are well modelled by AGNs. ",
        "introduction": "In recent years, deep radio surveys \\citep[\\eg][etc.] {windhorst90,condon89,prandoni01,hdf,pdfs,bondi03,hopkins03,gruppioni99,ciliegi99} have shown that there is an upturn in the differential radio source counts, normalised to a Euclidean value, below $\\sim 1\\mJy$. This upturn has generally been interpreted as the emergence of a new population of starforming galaxies \\citep[\\eg][]{condon89,rr93,garrett01} which does not show up at higher flux densities, where number counts are dominated by AGN. Given the strong correlation between $1.4\\GHz$ radio flux and the far infrared flux \\citep{hsrr85,carilliyun00}, which is itself a starformation rate indicator, radio luminosities can now be used as an independent measure of the starformation rate and, combined with other data, as a way of deconvolving the contribution of AGN and starbursts to the energy budget of the universe.  In order to investigate the nature of the objects that contribute to the X-ray Background radiation (XRB) a deep X-ray survey \\citep{imh98} with \\ROSAT was performed, using the \\PSPC detector, of a field at RA 13 34 37 Dec 37 54 44 (J2000). 96 sources were detected in a $30'$ diameter field above a flux limit of $2\\times 10^{-15}\\fluxerg$ in the $0.5-2.0\\keV$ band. The majority of the identifications were with broad line AGN, but some galaxies with narrow optical emission lines, known as narrow emission line galaxies (NELGs), were found. These galaxies had a range of X-ray luminosities from $\\sim10^{40}$ to $\\sim10^{43} \\ergps$, and a wide range of X-ray/optical ratios and optical emission line ratios. The highest X-ray luminosity yet found for a starburst galaxy is $\\sim10^{42} \\ergps$ \\citep{moran99,halpern95} and so the higher luminosity NELGs are most likely to be obscured AGN. However the lower luminosity NELGs, with low ($\\sim0.01 - 0.001$) X-ray/optical ratios could easily be explained by starburst emission \\citep[\\eg][]{kfg01}. Our conclusion was that the X-ray emission from NELGs, both as a sample and within individual galaxies, is due to a mixture of AGN and starbursts.  As a starburst is associated with radio emission \\citep[\\eg][]{normalgalaxies} then, if some of the NELG X-ray emission is produced by starbursts, one would expect detectable, associated radio emission; steep radio spectrum emission, detected and resolved on galaxy scales, would be a very strong starburst indicator \\citep[\\eg][]{normalgalaxies}. Preliminary $1.4\\GHz$ observations made with the Very Large Array (VLA) radio telescope of the US National Radio Astronomy Observatory\\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} were reported in \\cite{imh98}. We noted that 4 of the 10 brightest NELGs were detected as radio sources whereas a much lower fraction of X-ray QSOs was detected indicating that complementary radio observations would help determine the emission processes of these NELGs.  In order to improve on the significance of this early result we carried out deeper observations with the VLA at $1.4\\GHz$, and at $5\\GHz$, of the $30'$ diameter X-ray survey area. At $1.4\\GHz$ the primary beam diameter of the VLA is $30'$ and so the whole X-ray field can be covered in a single pointing. Thus the X-ray and $1.4\\GHz$ radio fields are very well matched. At $5\\GHz$, due to the smaller size of the primary beam, we have to make a number of pointings, over a grid of positions, in order to fully cover the X-ray field. We also performed higher resolution ($\\sim 0.2''$ beam) observations with MERLIN at $1.4\\GHz$. The ultimate aim is to resolve and detect any extended radio emission on galaxy scales out to redshifts $z>0.5$ and to determine the radio spectra of the various sources.  X-ray spectra are also a strong diagnostic of the emission mechanism, \\eg\\space soft thermal spectra would be associated with starburst emission. With \\XMMc, we therefore made a very deep (200ksec) observation (Page et al. 2003; Loaring et al. 2004, in preparation), providing high quality X-ray spectra in the $0.1-12 \\keV$ band, and detecting fainter X-ray sources than are seen in the \\ROSAT survey. In order to provide X-ray positions of sub arcsecond accuracy, and hence unambiguous identifications, we also observed the 13hr field for $4 \\times 30$ksec with the \\Chandra ACIS-I instrument \\citep{chandracat}. We detected 214 X-ray sources.  A preliminary analysis of the X-ray/radio relationship, based on these new data, is given by \\cite{gunn03}.   In this paper we present the new $1.4\\GHz$ VLA observations. The $5\\GHz$ observations are described in Seymour 2002\\nocite{seymour02}, and the MERLIN observations and a detailed correlation of the X-ray and radio sources will be presented elsewhere. The observations and data reduction are described in Section 2. We note that our observational set-up and analysis techniques are similar to those of \\cite{hdf}, to whom we refer readers for a more detailed description. The source extraction is described in Section 3. This $1.4\\GHz$ survey is comparable in depth to the deepest surveys previously made at this frequency \\citep[\\eg][and references therein]{hdf,hopkins03} and we detect sources at $4\\sigma$ significance down to $30\\uJy$ at the map centre. We are therefore able to define accurately the shape of the source counts below $1\\mJy$. In Section 4 we present these source counts, compare them with previous results and model them. The radio/optical/X-ray cross-correlation will be presented in future papers. ",
        "conclusions": "From a total of $\\sim 25$hours of data from the VLA at $1.4\\GHz$ in both A and B configuration, 449 sources are detected down to a peak flux density $4\\sigma$ limit of $30\\uJy$ in a circular field of $30'$ diameter. This survey is one of the most sensitive made at this frequency.  The source counts are similar to that of other deep surveys, showing an upturn of the Euclidean normalised source counts below $1\\mJy$ and are similar to those of the deepest Phoenix sample \\citep{hopkins03}, and both lie above the counts from the HDF. The HDF may be underdense on account of its selection criteria.  As in other deep surveys we assume that the upturn of the source counts results from a starforming population at medium redshifts. We base our modelling of this population on the local galaxy luminosity function of Sadler et al. (2002), together with mild luminosity evolution. Although our source counts are similar to those of \\cite{hopkins03}, our modelling results in a slightly different value of the luminosity evolution parameter, $Q$, probably due to our choice of the Sadler et al. local luminosity function (the choice of a different cosmology to that of previous authors does not strongly affect this value). Future guaranteed {\\it Spitzer} observations of this field will help distinguish the SFG and AGN contributions and refine the two component model of the source counts."
    },
    "0404/hep-th0404185_arXiv.txt": {
        "abstract": "We argue that the instability of Euclidean Einstein gravity is an indication that the vacuum is non perturbative and contains a condensate of the metric tensor in a manner reminiscent of Yang-Mills theories. As a simple step toward the characterization of such a vacuum the value of the one-loop effective action is computed for Euclidean de Sitter spaces as a function of the curvature when the unstable conformal modes are held fixed. Two phases are found, one where the curvature is large and gravitons should be confined and another one which appears to be weakly coupled and tends to be flat. The induced cosmological constant is positive or negative in the strongly or weakly curved phase, respectively. The relevance of the Casimir effect in understanding the UV sensitivity of gravity is pointed out. ",
        "introduction": "Quantum Gravity possesses a number of fundamental questions which are unsolved so far. Its apparent conflict with Quantum Mechanics, namely canonical commutation relations imposed on equal time hypersurface require the knowledge of the metric before the operator equations defining it is completed, questions the need of the quantum treatment of gravitation altogether \\cite{adler}. Let us assume that the metric corresponds either to an elementary or to a composite field and the usual quantum field theoretical path integral quantization applies. The next problem arises from the UV divergences. At this point one either seeks some superstring formalism which might be interpreted as a physical regulator or relies on a formal regulator. We choose the latter being a simpler possibility and do not attempt to remove the cutoff, i.e. we are not interested in the problem of physics at arbitrary high energies. Instead, our goal in this work is to make a small step toward the understanding of the nature of the radiative corrections up to a high but finite energy for a particular background geometry, those of the de Sitter space.  But there are still problems even after such a drastic reduction of the scope of the investigation. Regulators which preserve the local symmetries of the theory and allows analytical computations are available for the Euclidean section of the space-time only. When Euclidean Quantum Gravity is considered then some or all of the conformal modes become unstable \\cite{instab}, rendering the perturbative or saddle point approximation problematical. The usual approach to this complication is to stabilize Euclidean Quantum Gravity on the tree level by an appropriate modification of the Wick rotation \\cite{wickqg} or on the quantum level by altering the rules of quantization \\cite{quantst} or the functional measure for the path integral \\cite{geom}.  Similar instability is well known in Yang-Mills models. In fact, Euclidean Yang-Mills theory considered in the presence of a homogeneous magnetic field has unstable modes \\cite{savvidy}. Since the Euclidean Yang-Mills action is bounded from below the Euclidean vacuum is stable though the detailed mechanism of this stabilization is still unknown. When the time extent of the Euclidean space can be sufficiently large then the large amplitude modes resulting from the instability can be regarded as an indication of a condensate of the true, physical vacuum of the theory with real time. Can this lesson be relevant for Quantum Gravity and pointing toward the non triviality of its vacuum?  One can imagine three dynamical, consistent stabilization mechanisms for Euclidean Quantum Gravity. (i) It might happen that the true gravitational action contain higher order derivative terms which renders the classical action bounded from below. If these new terms do not involve new dimensional coupling constant then the stable vacuum will be dominated by fluctuations with scale close to the cutoff because the instability of the Euclidean Einstein action is in the UV regime. (ii) Another possibility is that the classical action is unbounded but the theory is stabilized by quantum fluctuations. This happens if the classical instability tries to squeeze the system into a small region of the configuration space and the kinetic energy generated by the uncertainty principle becomes strong. Certain numerical simulations seem to support this scenario for quantum gravity regularized by the Regge-calculus \\cite{lattice}. (iii) Yet another possibility is that pure gravity is indeed unstable but the stabilization is achieved by the coupling of gravity to matter \\cite{axanom,conform}.  But one should bear in mind that either possibility is realized the vacuum is non perturbative and contain a condensate. If the vacuum is stabilized on the $n$-th order of the loop expansion then it corresponds to the minimum of the effective action computed up to $\\ord{G^n_B}$. Here $G_B$ denotes the bare Newton constant which plays the role of $\\hbar$ in the absence of matter and counts the order of the loop expansion. The determination of the minimum of this effective action requires the knowledge of all 1PI vertex functions in the given order of the loop expansion. In the case of $n=0$, c.f. case (i), the solution of the classical equation of motion can be regarded as given by summing over infinitely many tree graphs. When $n>0$, c.f. cases (ii) and (iii), then an infinite sum of graphs up to $n$-loops gives the vacuum configuration.  We do not attempt to identify the true stabilization mechanism in this paper. Instead, we shall be contended with a small step along the line of reasoning outlined above and present a study of the dynamical environment in which the stabilization takes place by computing the effective potential for a homogeneous space-time curvature with vanishing unstable modes and by reading off from it the running of the Newton constant with the length scale of such a background geometry.  Before we continue the discussion of the computation of the effective potential let us bring another pertinent element, the Casimir effect, into the discussion. The vacuum energy of quantum field theories in flat space-time depends on a rather complicated manner on the choice of the space-time region $\\cal R$ populated by quantum fields \\cite{casres}. It is well known that the metric tensor controls all aspects of the geometry, among them the spatial extent of the system. In a space-time of homogeneous, positive curvature the free geodesics are deflected in such a manner that they stay in a region with linear size in the order of \\be\\label{circ} L\\approx\\frac1{\\sqrt{R}}. \\ee Such double role of the curvature, being an external field and an IR cutoff in the same time, provides a new, physical insight into the inherent sensitivity of quantum gravity on the UV sector. The argument is the following. The partition function computed on a homogeneous or at least slowly varying background field in a fixed space-time region $\\cal R$ develops UV finite dependence on the background field after having introduced the usual background field independent counterterms. One would expect the same conclusion when the partition function is considered on a background curvature. But this is not what happens because of the Casimir effect. The reason is that the Casimir effect is always UV divergent. To understand this property let us consider first the number of modes \\be\\label{modusok} {\\cal N}\\approx(\\Lambda L)^d \\ee in dimension $d$ where $\\Lambda$ denotes the UV cutoff of dimension energy and $L$ stands for the linear size of the system. The expression \\eq{modusok} diverges when the UV and IR cutoffs are removed because the length scale of a mode can be infinitely large or infinitely small. The result is that the UV and the IR dependences are mixed in $\\cal N$. In a similar manner the Casimir effect mixes the UV and the IR sectors, too. Therefore the dependence of the partition function on the curvature is always UV divergent. It is worthwhile noting that this point is nothing but the throughly studied issue of the coupling of the zero point fluctuations of the vacuum to the gravitation.  Our main goal in this work is to determine the renormalization of the effective potential for the curvature and some effective coupling constants in one-loop order. The effective potential, $\\gamma(R)$, is defined by the value of the effective action for homogeneous de Sitter spaces of curvature $R>0$ and will be computed in $\\ord{G_B}$ by restricting the unstable modes to zero, except the homogeneous conformal mode which is kept constant. We do not attempt to remove the UV cutoff and confine our investigation on the finite energy dynamics by keeping $\\Lambda$ fixed and ignoring the formal issue of renormalizability. We find a phase transition when the 'best' de Sitter space is selected by minimizing $\\gamma(R)$. For low enough value of the cutoff, $\\Lambda<<\\Lambda_{cr}$ the energetically preferred value of the curvature is $\\ord{\\Lambda^4G}$. The theory is weakly coupled at such or larger values of the curvature but the modes with length scale $\\ell>2/\\Lambda^2\\sqrt{G}$ are non perturbative. The other phase, realized by large cutoff $\\Lambda>>\\Lambda_{cr}$, prefers small curvature and is perturbative. This phase structure agrees qualitatively with the one found in lattice simulations \\cite{lattice} where a flat and a strongly curved phase were found for large or small values of the bare Newton constant, respectively. One should note in this respect that although the UV cutoff represents an upper limit for the curvature its relation to the Newton constant can be arbitrary, i.e. quantum gravity is a well defined and physically nontrivial theory for either $G\\Lambda^2<1$ or $G\\Lambda^2>1$.  The Casimir contribution to the effective action is proportional to the space-time volume which raises the possibility of generating a cosmological constant even in the absence of matter. In fact, an induced cosmological constant is found whose sign distinguishes the two phases mentioned above.  The UV cutoff, $\\Lambda$, appears as an unavoidable parameter in the main results of this paper, in the renormalization of the effective action and the phase structure. Such a state of affairs is certainly erroneous in the case of a renormalizable theory. But one should remember that whatever is the ultimate theory of quantum gravity the quantum effects are supposed be governed by the Einstein-Hilbert action within the energy and distance range explored up to today. The predictions of this theory, being nonrenormalizable, naturally do depend in an essential manner on the choice of the UV cutoff. But this dependence is restricted to scales comparable to the cutoff. At energy scales far below the cutoff quantum gravity is a useful effective theory because the nonrenormalizable, UV irrelevant terms of the action play no role here \\cite{effth}. It is worthwhile noting that this energy regime is determined by the cutoff only, the other dimensional parameter of the theory discussed in this work, the Newton constant, may take arbitrary values with respect to the UV cutoff in this effective gravity.  There have already been investigation of effective coupling constants carried out before. The one-loop correction to the Newtonian gravitational potential on flat space-time has been computed in the IR domain in Refs. \\cite{effnc} and a weak, $\\ord{R^{-2}}$, UV finite effect was found. The result is expected to be rather different when higher loop contributions are retained because those graphs contain UV divergences. The self interaction of gravitons, considered in our computation, is UV divergent already at the one-loop level and the comparison of the running Newton constant presented in these works with those arising from the computation of the gravitational forces between matter particles is not obvious. In fact, the difference of the dependence on the scale of a space-time variable or of a field amplitude is usually a manifestation of a nontrivial wave function renormalization constant \\cite{int}. The non triviality of the wave function renormalization constant can be expected on the simple dimensional ground that the metric tensor is a dimensionless field as opposed to the usual one. The conventional renormalization group has already been applied to determine the running of the coupling constants of Grand Unified Theories on curved geometry and a number of scaling laws were identified \\cite{convrg}. But one should bear in mind that the multiplicative renormalization which serves as the basis for the conventional renormalization group equation is valid for renormalizable models only. The renormalization group flow was computed for the Newton and the cosmological constants in Ref. \\cite{reuter} for de Sitter spaces. Two kinds of fixed points were found, a Gaussian and gauge-dependent non-Gaussian. The flow was extended to arbitrary dimension and value of the gauge fixing parameter in Ref. \\cite{litim}. But it is difficult to compare thess result with ours, partly because this computation involves a formal one-loop integral over the unstable modes, partly because of the presence of the cosmological constant. The gauge dependence of the renormalization group flow was investigated in Ref. \\cite{gaugdep}. Though our motivation is different, our work is closest in technics to Ref. \\cite{venez} where the dynamical suppression of the cosmological constant was studied in de Sitter spaces.  The organization of the paper is the following. The effective potential and the technical details such as the regularization, gauge fixing and the eigenfunctions of the Laplace operator in de Sitter spaces are introduced in Section II. The results of the numerical computation are presented in Section III. The problem of the stability of the conformal modes is discussed in Section IV. Finally, Section V is for the summary and conclusions. A very brief summary of the simplest appearance of the Casimir effect is presented in the Appendix. ",
        "conclusions": "We argue in this paper that the instability of Euclidean Einstein gravity is an indication of a non trivial vacuum structure whose description requires non perturbative methods. As a simple step in this direction the one-loop renormalization was considered on an Euclidean de Sitter background geometry by suppressing the unstable modes. The similarity of this problem with the Casimir effect was emphasized together with the role of the curvature as an IR cutoff which makes quantum gravity specially sensitive to the UV regime. The effective action was minimized within the family of homogeneous spaces in order to identify the vacuum.  Two phases were found and the effective strength of interaction behaves in a manner which is contrary to what the tree-level dynamics would suggest. In the small cutoff phase, $\\Lambda<<1/\\sqrt{\\ell_B}$, the conformal modes are stabilised by the one-loop effective action for the other tree-level stable modes and we find a close analogy with the Savvidy vacuum of Yang-Mills theories. The system prefers a non vanishing curvature, $R_{min}\\approx\\Lambda^4\\ell^2_B$ and $\\ell_{Pl}=\\sqrt{2}\\ell_B$. The tree-level theory is weakly coupled for arbitrary curvature but the fluctuations acquire large amplitude at length scales $\\ell>1/\\sqrt{R_{min}}$ and the theory is weakly coupled for $R_{min}\\le R$ only when these modes are excluded due to the smallness of the space-time volume. Deeply in the large cutoff phase when the cutoff is increased far beyond the transition point then the minimum of the effective action is reached at $R_{min}\\approx0$, the fluctuations are suppressed. But the stability of the conformal modes remain an open question.  We proposed two parametrizations for the effective action. The running Newton constant can be defined in the effective theory without cosmological term which amounts to assuming formally massless quantum fluctuations. This coupling constant determines the length scale below which the quantum fluctuations are strong. Another possibility is to allow a radiative correction induced cosmological term explicitly in the effective action. This choice is motivated by the proportionality of the Casimir effect with the space-time volume and is more physical because it produces slower dependence on $R$ in the effective parameters than the previous case. Furthermore, a higher order derivative term is recovered by means of this parametrization. The effective Newton constant agrees with the bare one in this scheme and the two phases are distinguished by the sign of the induced cosmological constant and the coefficient of the higher order derivative term.  The change of the UV cutoff while the other parameters are kept constant modifies the physical content of the theory and even drives the system through a phase transition. The cutoff dependence in general is not surprising in a nonrenormalizable theory. It might turn out that some so far unknown nonperturbative effects render the Einstein-Hilbert theory renormalizable but in lacking these elements one has to accept that the theory which is nonrenormalizable in the framework of the loop-expansion produces cutoff-dependent results at one-loop level. Our conclusions drawn correspond to energy scales well below the cutoff and are therefore independent of the details of the regularization procedure since this latter is always represented by irrelevant operators.  The cutoff dependence of nonrenormalizable models usually remains quantitative. The strong cutoff-effects reported in this work can be traced back to the different sign of the induced cosmological constant in the two phases. This parameter is renormalizable and its sign influences the physics in a more substantial, quantitative level.  We encountered two level of instabilities in this work. In the tree-level theory the conformal modes are unstable but they seem to be stabilized by the quantum fluctuations at least in the long distance region of the high cutoff phase. The one-loop improved effective theory remains unstable in the low cutoff phase. The cosmological constant appears formally as the parameter $-m^2$ in a scalar model. The instability of the low cutoff phase with  $\\lambda>0$ should therefore be similar to those of a scalar model with spontaneously broken symmetry. In the latter case the tree-level contributions to the renormalization group equation were found to be crucial \\cite{trrg} in order to establish a systematical approach to the mixed phase of first order phase transitions, in particular the spinodal phase separation \\cite{spinod} where the soft modes arising from the breakdown of space-time symmetries by the inhomogeneous saddle points generate strong nonperturbative effects, among others the Maxwell-cut.  We continue with few comments on the results of the previous Section from the point of view of the case (ii) of the stabilization mechanism, mentioned in the Introduction and identified in the low cutoff phase. It was found that the one-loop effective action may stabilize the vacuum and generate a condensate of the metric tensor even in the absence of matter. It remains to be seen if such a condensate involves the conformal modes only or the spin 2 components, too. Another open question is whether the new vacuum with condensate is weakly or strongly coupled. The dimensionless ratio $\\kappa_B/\\Lambda$ available in this nonrenormalizable theory may generate small or large dimensionless numbers and this question can be settled by detailed computation only. Though the other stabilization mechanisms were not pursued in this work it is clear that the excitations above the condensed vacuum produced by either mechanism listed in the Introduction are different than those above the naive, 'empty' vacuum and the success of general relativity in the classical domain requires more careful justification.  One arrives at a surprising conjecture by following the analogy offered by the scalar model. It is well known that the speed of sound is vanishing in the mixed phase. This is due to the strongly coupled soft modes, the translations of the domain walls. In more physical terms the domain walls absorb the sound. In a similar manner one expects the speed of the propagation of gravitons be vanishing in the low cutoff phase. This situation where the naive plane waves ceases to propagate and drop out from the asymptotical sector is called confinement. If this mechanism is realized then gravitons will be confined in the low cutoff phase due to the quantum liquid nature of the vacuum. The confinement radius which can be approximated by the length scale of the IR Landau pole is in the order of magnitude of the size of the stable de Sitter space, indicating simply that the gravitons should be confined by the horizon.  Finally we mention other natural extensions of this work, the inclusion of matter fields and the cosmological constant in the bare action. Condensate modifies the vacuum structure and it may stabilise the gravitational sector by breaking conformal invariance. The interplay of condensates in the coupled Einstein-Yang-Mills system is a particularly interesting problem. The bare cosmological constant, included from the very beginning in bare action should preserve the topology of the phase structure but make the result more realistic. We plan to report on such results soon."
    },
    "0404/hep-ph0404278_arXiv.txt": {
        "abstract": "We study dynamics of flat directions in the minimal supersymmetric standard model taking account of its constituent fields. It is found that there exist new instabilities due to the D-term potential and the nature of these instabilities depends on the eccentricity of the orbit. For a roughly circular orbit, it is similar to  narrow-band resonance described by the Mathieu equation. For an  elliptic orbit,  the instabilities exhibit tachyonic nature. In the narrow-band regime, we show that the existence of the instabilities is crucial to the formation of $Q$-balls. The $Q$-ball formation proceeds through  quasi-stable state called  ``$I$-ball''. The transition from $I$-balls to $Q$-balls is made efficient by the $D$-term instability and during this process some fraction of the charge of the $Q$-ball is emitted. This discovery may revive the scenario which relates the baryon number and dark matter density of the universe. Furthermore, the tachyonic $D$-term instability causes the drastic decay of the flat direction moving in an orbit with relatively large eccentricity. Thus the evolution of flat directions is completely altered by the appearance of this instability. ",
        "introduction": "\\label{sec:intro} The scalar potential in the minimal supersymmetric standard model (MSSM) has flat directions along which there is no classical potential in supersymmetric (SUSY) limit. In the real world, the potential along the flat directions is not completely flat because SUSY breaking induces soft masses of the order $\\sim O(100)$~GeV. Such flat directions  play important roles in cosmology. In the inflationary universe some fields corresponding to flat directions have large field values during inflation, and start oscillation after their mass becomes comparable to the Hubble parameter. The flat directions generally consist of squarks, sleptons and Higgs, and hence have baryon and/or lepton numbers. Thus, the oscillation of the flat directions and subsequent decay can produce large baryon and/or lepton numbers in the universe if some baryon(lepton) number violating terms exist in the potential. This mechanism was discovered by Affleck and Dine~\\cite{Affleck:1984fy} and it has been one of promising scenarios for explaining the baryon asymmetry of the present universe. Hereafter we call scalar fields corresponding to flat directions Affleck-Dine (AD) fields.  Recently, it was found that  condensates of some AD fields evolve into non-topological solitons called $Q$-balls~\\cite{Coleman:1985ki} through strong instabilities during oscillation~\\cite{Kusenko:1997si,Enqvist:1998en}. In fact, numerical simulations showed that initial tiny fluctuations of the AD field rapidly grow and deform into large $Q$-balls~\\cite{Kasuya:1999wu,Kasuya:2000wx}. The existence of $Q$-balls complicates dynamics of the AD field and significantly changes the scenario of the AD baryogenesis~\\cite{Kusenko:1997si,Enqvist:1998en, Kasuya:2000sc,Kasuya:2001hg}. $Q$-balls also provide several interesting possibilities. For example, $Q$-balls can account for the dark matter and baryon asymmetry of the universe simultaneously~\\cite{Kasuya:2000sc, Kasuya:2001hg}, or they may produce very large lepton asymmetry~\\cite{Kawasaki:2002hq,Ichikawa:2004pb}.  So far, dynamics of  flat directions has been studied with use of ``single-field parametrization'', i.e. a flat direction is described by one complex scalar field. In fact, however, the flat direction is actually composed of several scalar fields. When one only considers the dynamics of a homogeneous AD field, the single-field parametrization may be good enough. However, we now know that the fluctuations of the AD field is crucial to $Q$-ball formation. Since the field fluctuations increase the number of degrees of freedom in the dynamical system, the single-field parametrization is not always adequate to describe the evolution of the fluctuations and multi-field analysis is necessary. In the previous paper~\\cite{Kawasaki:2004th}, we demonstrated the importance of the multi-field analysis when one of component fields that form $Q$-balls has a large decay rate.  In this paper, we study the dynamics of the AD field taking account of its constituent fields. We find that there exist instabilities due to the $D$-term potential and the nature of these instabilities is determined by the eccentricity of the orbit. For a roughly circular orbit, it is described by narrow-band resonance. However, as the eccentricity increases the instability exhibits another character: tachyonic instability.  In the former case, since the instability is the narrow-band resonance, its effect on the cosmological evolution of the homogeneous mode of the AD field is not so important. However, the existence of the instabilities is crucial to the formation of $Q$-balls. In general, in gravity mediated SUSY breaking models, the $Q$-ball formation is complete through quasi-stable state  ``$I$-balls''~\\cite{Kasuya:2002zs} inside which the AD field rotates in an elliptical orbit. It is found that the D-term instability quickly makes the orbit of the AD field circular and a $Q$-ball configuration is reached. During this process some fraction of the charge confined in the $Q$-ball is emitted, which is different from the previous result obtained with use of the single-field parametrization. As a result  the baryon density of the universe is accounted for by both $Q$-ball and residual baryon emitted during the $Q$-ball formation. On the other hand the dark matter is explained by the lightest SUSY particle (LSP) produced in the $Q$-ball decay. This may revive the scenario which relates the baryon number and dark matter density of the universe.  Meanwhile,  if the instability becomes tachyonic,  its effect on the cosmological evolution of the AD field is no longer negligible. Irrespective of whether $Q$-balls are formed or not, the homogenous motion always suffers from  rapid grow of the instabilities. Therefore  the AD field moving in an orbit with large eccentricity decays soon after it begins to oscillate. We expect that the similar instabilities will exist in many other cases. For instance, the instabilities might change the reheating processes of some supersymmetric inflation models.  In the next section, we investigate the D-term instability, and the effect on $Q$-ball formation is studied in Sec.~\\ref{sec:qb_form}. Sec.~\\ref{sec:tachyonic} is devoted to the discussion on the fate of the flat directions under the effect of the tachyonic instability and finally we give discussions and conclusions in Sec.~\\ref{sec:discuss}. ",
        "conclusions": "\\label{sec:discuss} In this paper we have investigated the dynamics of the flat directions taking account of the $D$-term potential. To this end, we adopted the multi-field parametrization of the flat directions, which enables us to deal with the relative fluctuations between the constituent fields. It is found that there exist instabilities due to the D-term potential and the nature of these instabilities depends on the axial ratio $\\varepsilon$ of the oscillating flat direction. If $\\varepsilon$ is close to $1$, the $D$-term instability is similar to narrow resonance described by the Mathieu equation with $q \\ll 1$. On the other hand, the instabilities show the tachyonic nature for $\\varepsilon <1$. Due to this tachyonic instability, the cosmological evolution of the flat direction are always under the influence of $D$-term instability. In particular we have shown that the existence of the instabilities is crucial to the formation of $Q$-balls. Since the $D$-term instability violates the conservation of the adiabatic charge, $I$-balls are no longer quasi-stable, and therefore the transition from $I$-balls to $Q$-balls occurs very efficiently. It is important to note that  during this process some fraction of the charge of the $Q$-ball is emitted, which may revive the relation between the baryon number and dark matter density of the universe. Furthermore, the $D$-term instability drastically changes the evolution of the flat directions. Unless $Q$-balls are formed, the flat direction completely decays relatively soon after it begins to oscillate. We believe that the existence of such instabilities is generic to the system of scalar fields with the $D$-term potential. For instance, in the reheating stage of the $D$-term inflation models~\\cite{Binetruy:1996xj}, the scalar fields are oscillating in the similar potential. Thus the reheating processes of inflation models of this type might be affected by the $D$-term instability. However, further investigation is clearly necessary and is beyond the scope of this paper.      \\subsection*"
    },
    "0404/astro-ph0404267_arXiv.txt": {
        "abstract": "Modeling of radio and X-ray observations of supernovae interacting with their circumstellar media are discussed, with special application to SN 1993J and SN 2002ap. We emphasize the importance of including all relevant physical mechanisms, especially for the modeling of the radio light curves. The different conclusions for the absorption mechanism (free-free or synchrotron self-absorption), as well as departures from an $\\rho \\propto r^{-2}$ CSM, as inferred by some authors, are discussed in detail. We conclude that the evidence for a variation in the mass loss rate with time is very weak. The results regarding the efficiencies of magnetic field generation and relativistic particle acceleration are summarized. ",
        "introduction": "\\label{CF-intro} The interaction of supernovae (SNe) with their circumstellar medium (CSM) offer important clues to both the nature of the SN progenitors, the hydrodynamics of the explosion, the environment of the SN, and the physics of high velocity shock waves. The now convincing connection of Type Ic SNe and GRB's has also made the study of the SN environment especially interesting. Basically, the standard picture of the SN interaction with the surroundings is only a non-relativistic version of the standard afterglow scenario (e.g., Chevalier, this Proceedings).  We will discuss a few issues related to the SN interaction.  For a more detailed exposition see the recent review in \\cite{CF-CF03}. ",
        "conclusions": ""
    },
    "0404/astro-ph0404051_arXiv.txt": {
        "abstract": "The results of the lowest frequency spectral survey carried out toward a molecular cloud and sensitive observations at selected frequencies are presented. The entire Arecibo C-band (4--6~GHz) was observed towards the cyanopolyyne peak of TMC-1 with an rms sensitivity of about 17--18 mK (about 2--2.5~mJy). In addition, a number of selected frequency ranges within the C-band and X-band (8--10~GHz) were observed with longer integration times and rms sensitivities  7--8 mK ($\\approx$ 2 mJy) or higher.  In the spectral scan itself, already--known H$_2$CO and HC$_5$N lines were detected. However, in more sensitive observations at selected frequencies, lines of C$_2$S, C$_3$S, C$_4$H, C$_4$H$_2$, HC$_3$N and its $^{13}$C substituted isotopic species, HC$_5$N, HC$_7$N, and HC$_9$N were found, about half of them detected for the first time. The rotational temperatures of the detected molecules fall in the range 4--9 K. Cyanopolyyne column densities vary from $5.6\\times 10^{13}$~cm$^{-2}$ for HC$_5$N to $2.7\\times 10^{12}$~cm$^{-2}$ for HC$_9$N. Our results show that for molecular observations at low frequencies (4--10 GHz) to be useful for studying dark clouds, the sensitivity must be of the order of 5--10 mK or better. To date, observations at around 10~GHz have been more productive than those at lower frequencies. ",
        "introduction": "Most known interstellar molecules have been detected in the millimeter wavelength range. This is clearly related to the fact that, in general, the simplest and hence the lightest molecules are the most abundant cosmic molecular species. Their strongest rotational lines at the temperatures characteristic of the molecular interstellar medium arise at mm and submm wavelengths. Heavier molecules, often with large permanent dipole moments, have detectable rotational lines at lower frequencies in the microwave range. Transitions between low energy levels are more favorable in terms of line intensities, especially in regions of low or modest density and low temperature. The latter include, in particular, the inner parts of dark clouds, which are among the coldest regions of the interstellar medium, because they have no internal energy sources and are shielded from external sources of radiation by their own gas and dust.  Another reason why observations at low frequencies are needed is the fact that the line intensities of molecules having fractional abundance below approximately $10^{-11}$ relative to H$_2$ fall below the confusion level in the millimeter wavelength range.  In such cases, sensitive observations at low frequencies, which are less ``contaminated'' by the emission of simple molecules, may help. In addition, hyperfine splitting is often negligible at high frequencies, but detectable at low frequencies (e.g. for species HC$_5$N, HC$_7$N, etc.). Hyperfine splitting may be used to determine line opacities and thus to obtain more accurate molecular column densities and abundances.  Spectral scans, covering wide ranges of frequencies, are powerful tools for molecular searches. Typically many lines of the same molecule fall within the observed spectral window, making it possible to more reliably identify observed spectral features. A number of spectral scans of molecular clouds have so far been carried out in atmospheric windows between 17~GHz and 700~GHz.  A selection of the frequency ranges covered and references are 17.6--22 GHz \\citep{bel93}; 72--144 GHz  \\citep{cum86,joh84,tur91}; 330--360 GHz \\citep{jew89}; 455--507 GHz, \\citep{whi03}; 607--725 GHz, \\citep{sch01}; 780--900 GHz, \\citep{com02}. Recently, \\cite{kai04} carried out a spectral scan of the dark cloud TMC-1 in the frequency range between 8.8 and 50 GHz. However, no spectral scan has been performed at frequencies below 8.8~GHz. Therefore we undertook a spectral survey of TMC-1 at 4--6 GHz. This is the lowest frequency spectral survey carried out toward a molecular cloud\\footnote{A similar spectral scan of IRC+10216 has been performed using the Arecibo telescope by Araya et al. (2003).}.  TMC-1 has a shape of a ridge elongated in the northwest-southeast direction. The structure and chemical composition of TMC-1 have been studied by \\cite{hir92}, \\cite{pra97}, \\cite{tur00}, and \\cite{dic01}. The temperature of this object is about 10~K, which is typical for dark clouds; the H$_2$ density estimates towards the ``cyanopolyyne peak'' (see below) vary between $4\\times 10^3$~cm$^{-3}$ \\citep{tur00} and $2\\times 10^4$~cm$^{-3}$ \\citep{pra97}.  The study by Peng et al.~(1998) revealed 45 clumps which are grouped into 3 cylindrical features oriented along the ridge. The LSR velocities of these features are approximately 5.7, 5.9, and 6.1~km~s$^{-1}$. TMC-1 has proven to be an excellent object for the study of chemistry in dark clouds under quiescent conditions. It shows a carbon--rich chemistry with chemical gradients across the densest part of the ridge extending over 0.2 pc x 0.6 pc. The most prominent chemical feature of TMC-1 is the high abundance of various carbon--chain molecules, such as cyanopolyynes HC$_{2n+1}$N and radicals C$_n$H. The peak position of the line intensities of cyanopolyynes (the cyanopolyyne peak) was found to be located $\\approx 7'$ southeast from another prominent position in TMC-1, the peak of the ammonia line emission. The heaviest known interstellar molecule, HC$_{11}$N, as well as a number of other heavy molecules have been found exactly towards the cyanopolyyne peak, making it an especially interesting object for observations at low frequencies. ",
        "conclusions": "The results of the lowest frequency spectral survey carried out toward a dark cloud, studying TMC-1 in the 4--6 GHz range using the Arecibo radio telescope, and those from sensitive observations at selected frequencies in  C-- and X--bands can be summarized as follows:  \\begin{itemize} \\item A number of molecular lines were detected. The majority of the detected lines belong to the cyanopolyynes HC$_3$N, HC$_5$N, HC$_7$N, and HC$_9$N. \\item The rotational temperatures of the detected molecules fall in the range 4--9 K. Cyanopolyyne column densities vary from $5.6\\times 10^{13}$~cm$^{-2}$ for HC$_5$N to $2.7\\times 10^{12}$~cm$^{-2}$ for HC$_9$N. \\item Molecular observations at low frequencies (4--10 GHz) can be useful for studying dark clouds, but to be really effective, the sensitivity must be increased by a factor of 2 to 3. To date, observations at around 10~GHz have been more fruitful than those at lower frequencies.  \\end{itemize}"
    },
    "0404/astro-ph0404098_arXiv.txt": {
        "abstract": "We present new and accurate Near-Infrared $J$ and $K_s$-band data of the Large Magellanic Cloud cluster Reticulum. Data were collected with SOFI available at NTT  and covering an area of approximately $5\\times5$ arcmin around the center of the cluster. Current data allowed us to derive accurate mean $K$-band magnitudes for 21 fundamental (\\RRab) and 9 first overtone (\\RRc) RR Lyrae stars. On the basis of the semi-empirical $K$-band Period-Luminosity-Metallicity ($PLZ_K$) relation we have recently derived, we find that the absolute distance to this cluster is $18.52\\pm0.005 (random) \\pm0.117 (systematic)$. Note that the current error budget is dominated by systematic uncertainty affecting the absolute zero-point calibration and the metallicity scale. ",
        "introduction": "Absolute distances of RR Lyrae stars play a fundamental role in stellar astrophysics, and the reasons are manifold. They are widely adopted to establish the absolute age of globular clusters (GCs) and to properly compare predicted and empirical Color-Magnitude diagrams. Moreover, they are robust stellar tracers of low-mass, old populations and are ubiquitous across the Galactic spheroid as well as in Local Group galaxies. During the last few years a paramount theoretical and observational effort has been devoted to the RR Lyrae distance scale to improve not only the intrinsic accuracy but also to constrain the occurrence of deceptive systematic errors (Caputo et al. 2000; Walker 2000; Bono 2003; Cacciari \\& Clementini 2003).  In a series of previous papers (Bono et al. 2001,2002,2003, hereinafter B01, B02, and B03) we have drawn the attention to the $K$-band Period-Luminosity-Metallicity ($PLZ_K$) relation of RR Lyrae stars. This relation was originally discovered by Longmore et al. (1990, hereinafter L90) and presents several undisputable advantages when compared with other methods available in the literature. In particular, the $PLZ_K$ relation is only marginally affected by evolutionary effects as well as by a spread in stellar mass inside the instability strip. Moreover, it is only weakly affected by uncertainties in the reddening correction and shows a linear trend when moving from metal-poor to metal-rich objects. Current theoretical predictions appear in very good agreement with empirical observations, and indeed the pulsation parallax to RR Lyr itself estimated by B02, is in very good agreement with the trigonometric parallax measured by Benedict et al. (2002) with the Fine Guide Sensor on board the Hubble Space Telescope (HST).  Although the theoretical framework concerning the $PLZ_K$ relation has been thoroughly discussed, only a few investigations have been devoted to near-infrared (NIR) observations of cluster and field RR Lyrae stars. In fact, with the exception of the $K$-band data collected by L90 for cluster RR Lyrae and Baade-Wesselink RR Lyrae (see B03 for a detailed list) only one investigation has been devoted to RR Lyrae stars in the Galactic Bulge (Carney et al. 1995). The observational scenario has been recently complemented with new $K$-band data collected by Butler (2003), using the adaptive optics camera available on the 3.5m Calar Alto telescope, for a small sample of RR Lyrae stars located in the very center of the GC M3. The above discussion shows quite clearly that we lack new and accurate NIR data to supply a sound empirical estimate of both the slope and the zero-point of the $PLZ_K$ relation but also to improve current empirical constraints on the metallicity dependence. To fill this gap we undertook an observational project aimed at providing a comprehensive observational investigation of cluster and field RR Lyrae stars in the NIR $J$ and $K_s$-bands.  In this investigation we present $K$-band data for RR Lyrae stars in the LMC cluster Reticulum. We selected this cluster for the following reasons: it presents a sizable sample of well-studied RR Lyrae stars (32, Walker 1992), its central density is very low ($\\log \\rho_0\\approx 0\\; M_\\odot pc^{-3}$, Peterson \\& Kunkel 1977), and therefore the NIR photometry is not seriously affected by crowding, and the metal content of this cluster has been estimated using medium resolution spectra ($[Fe/H]\\approx=-1.71\\pm0.1$, Suntzeff et al. 1992). In section 2, we present the observations and the fit to the $K_s$-band light curves, while in section 3 we discuss the new distance determination together with intrinsic and systematic uncertainties affecting the current estimate. ",
        "conclusions": "Based on the mean $K$-magnitudes from the previous section we now proceed to determine the distance to the Reticulum cluster. To further improve the theoretical calibration of the $PLZ_K$ relation, B03 devised a new pulsation approach that relies on mean $K$-band magnitudes and $V-K$ colors. In particular, they derived new $PLZ_K$ relations (see their relations 3 and 4) and Period-Luminosity-Color-Metallicity, $PLCZ_{(V,K)}$, relations (see their relations 7 and 8) that include the luminosity term. Note that these relations when compared with theoretical predictions by B01 present three main advantages: {\\em i)} they rely on a larger set of pulsation models;  {\\em ii)} independent theoretical relations have been derived for fundamental and first overtone pulsators; {\\em iii)} the new $PLZ_K$ and $PLCZ_{(V,K)}$ relations were derived by adopting predicted mass values for each assumed chemical composition, while B01 adopted an ensemble average. To root the calibration of the zero-point as well as of the coefficient of the metallicity term of the new $PLZ_K$ relation on empirical data, B03 adopted a sample of field RR Lyrae stars for which good $V$ and $K$-band light curves, accurate reddening corrections, and metal abundances were available in the literature. For fundamental mode pulsators they found:  \\begin{equation} M_K = -0.770 (\\pm 0.044) - 2.101 \\log P + 0.231 (\\pm 0.012) \\FeH \\end{equation}  Suntzeff et al. (1992) have determined the metallicity of the Reticulum cluster to be $\\FeH=-1.71\\pm0.1$ based on spectroscopy of the Calcium II triplet at around 860~nm for 9 individual stars in the cluster. Entering this value in equation 2 we can determine the distance modulus, $(m-M)_K=\\langle K \\rangle - M_K$, to each star. Forming the weighted average of these estimates we find for the \\RRab stars a value of $(m-M)_K=18.531\\pm0.006$, while for the \\RRc stars, using the fundamentalized periods, we find $(m-M)_K=18.539\\pm0.009$ and for the complete sample we find $(m-M)_K=18.534\\pm0.005$.  The reddening towards Reticulum has been determined by Walker (1992) who finds a very low value of $E(B-V)=0.03\\pm0.02$. We combine this value with the reddening law from Cardelli, Mathis, \\& Clayton (1989) of $A_K = 0.114 \\times 3.1 \\times E(B-V)$ which gives a minuscule absorption in the $K$-band of $A_K = 0.011$~mag. Correcting the modulus for the reddening we find a best estimate of $(m-M)_0=18.523\\pm0.005$, where the error estimate only consider intrinsic random errors. The referee suggested to perform a detailed check of the systematic uncertainties affecting distance estimates based on the semi-empirical relation derived by B03. To accomplish this goal we selected the Galactic Globular Clusters (GGCs) for which are available both optical and NIR mean magnitudes. To avoid subtle uncertainties in the reddening corrections we only selected GGCs with reddening corrections smaller than 0.2. We ended up with four GGCs whose metallicity ranges from $[Fe/H]=-1.27$ (M~5) to $[Fe/H]=-2.26$ (M~15). To estimate the difference in the distance moduli we adopted distances based on the Baade-Wesselink (BW) calibration provided by Fernley (1994) as well as on the First Overtone Blue Edge (FOBE) provided by Caputo et al. (2000). Data listed in the last two columns of Table 2 show that the difference ranges from $\\Delta \\mu=0.32$ for the less metal-poor cluster (M~15) to $\\Delta \\mu=0.10$ for the more metal-rich cluster (M~5).  This finding further strengthens the result obtained by B03, i.e. the difference between distances based on $PLZ_K$ and on the BW method depend on the mean metallicity and it decreases in the metal-rich regime (see figures 6 and 8 in B03). Unfortunately, we still lack accurate mean K-band magnitudes for RR Lyrae stars in a larger sample of metal-rich GGCs. Moreover, current uncertainties affecting distance moduli based on mean V-band magnitudes are systematically larger than 0.1 mag. This hampers any precise quantitative conclusion on the systematic uncertainty affecting the $PLZ_K$ distance scale.  It is noteworthy, that cluster distances presented here appear to be systematically larger than distances based on the FOBE method and on the $PLZ_K$ calibration based on theoretical models provided by B01. The semi-empirical calibration by B03 presents a steeper dependence on the metallicity (0.231 vs 0.167) and the reasons for the difference between the empirical and theoretical scales are not clear yet (see also Carney et al. 1992). Finally, we note that the current distance estimate to $\\omega$ Cen is also 0.1 mag larger than the distance modulus derived by Thompson et al. (2001) on the basis of the eclipsing binary OGLE17. Although, the analysis of current uncertainties affecting the RR Lyrae distance is far from being definitive we decided to adopt a systematic uncertainty of 0.1 in the calibration of the zero-point.  The uncertainty on the reddening correction is quite negligible, but we have to account for the uncertainty in the metallicity and in the metallicity scale. Equation 2 was derived using the metallicity estimates provided by Fernley et al. (1998) and by Zinn (1985) for field and cluster RR Lyrae stars. The metallicities of field RR Lyrae rely on a calibration of the $\\Delta S$ method provided by Fernley \\& Barnes (1997), i.e. $[Fe/H]=-0.13 - 0.195 \\Delta S$, that is quite similar to the calibration provided by Clementini et al. (1995), $[Fe/H]=-0.08 - 0.195 \\Delta S$. These calibrations supply metal abundances quite similar to the Zinn \\& West (1984) metallicity scale. On the other hand, Gratton (1999) provided a new calibration, $[Fe/H]=-0.03 - 0.176 \\Delta S$, that relies on the Carretta \\& Gratton (1997) metallicity scale. By accounting for these empirical uncertainties on the zero-point, on the slope as well as on the systematic uncertainties affecting the metallicity scale (Rutledge, Hesser, \\& Stetson 1997; Kraft, \\& Ivans 2003) we estimate the uncertainty on the metallicity to be of the order of 0.25 dex. This means that the uncertainty on the mean metallicity in equation 2 as well as in the cluster metallicity introduces an uncertainty of the order of 0.06 mag in the current distance determination. Therefore, we end up with a distance to Reticulum of $18.52\\pm0.005 (random) \\pm0.117 (systematic)$.  Reticulum is located 11 degrees away from the center of LMC, and therefore we can only supply weak constraints on the absolute distance to LMC itself. According to our optical photometry and by assuming a mean reddening of $E(B-V)=0.03\\pm0.02$ (Walker 1992) the mean visual magnitude of RR Lyrae stars in Reticulum is: $V_0=18.96\\pm0.05$ mag. Note that this estimate is in very good agreement with the estimate provided by Walker (1992), i.e. $V_0=18.98\\pm0.04$. Recent estimates of the mean visual magnitude of RR Lyrae stars in the LMC bar range from $V_0=19.05\\pm0.06$ (108 stars) by Clementini et al. (2003) to $V_0=18.90\\pm0.02$ (7110 stars) derived by the OGLE team (Soszynski et al. 2003) and to $V_0=18.99\\pm0.02 (random)\\pm 0.16 (systematic)$ (80 first overtones) derived by the MACHO team (Alcock et al. 2004). The key differences in these estimates are: statistics, absolute zero-point calibration, and reddening correction. It goes without saying that according to these estimates the position of Reticulum ranges from $\\approx 3$ kpc in front of LMC to $\\approx 2$ kpc beyond LMC. New and homogeneous NIR data for a sizable sample of RR Lyrae in the LMC bar are mandatory to settle the uncertainty on the relative position of this cluster.   The empirical evidence brought out by current data suggests that the $PLZ_K$ relation is characterized by a global accuracy better than 0.12 mag. Note that the current discrepancy between HB evolutionary models constructed by adopting different assumptions concerning the input physics is of the order of 0.1-0.15 mag (Cassisi et al. 1999; VandenBerg et al. 2000). However, the most recent systematic survey of $K$-band data for cluster RR Lyrae stars dates back to L90. Therefore new and homogeneous NIR data for GGCs with sizable samples of RR Lyrae stars are strongly required to shed new lights on this {\\em vexata questio}. At the same time, a new empirical calibration of the $PLZ_K$ relation based on RR Lyrae stars for which heavy-element abundances has been measured on high-resolution spectra is also necessary.  The intrinsic accuracy of the $PLZ_K$ relation seems also very promising in providing a new distance scale for GCs that can allow us to improve their absolute age estimate. It is noteworthy that new cluster data will allow us to derive an independent $PLZ_K$ relation for first overtone RR Lyrae stars, and in turn to further improve the intrinsic accuracy of this method. Finally, due to the marginal dependency on reddening corrections it could also be soundly adopted to derive the 3D-structure not only of the Magellanic Clouds, but also of the Galactic bulge."
    },
    "0404/astro-ph0404321_arXiv.txt": {
        "abstract": "Properties of the grain-scattered X-ray halo of the eclipsing X-ray binary pulsar 4U 1538-52 are derived from a 25 ksec observation by the {\\it Chandra} X-ray observatory extending from just before its eclipse immersion to near mid-eclipse. Profiles of the observed halo, compiled in two energy ranges, 2 to 4 keV and 4 to 6 keV, and three time intervals before and after the eclipse, exhibit a three-peak shape indicative of a concentration of the interstellar dust grains in three discrete clouds along the line of sight.  The observed profiles are fitted by the profiles of a simulated halo generated by a Monte Carlo ray-tracing code operating on a model of three discrete clouds and a spectrum of the photons emitted by the source over a period of time extending from 270 ksec before the observation began till it ended. In the model, the spectrum before the observation began is expressed as a function of the orbital phase of the pulsar and is derived as an average over 6.3 years of data accumulated by the All Sky Monitor of the {\\it Rossi X-Ray Timing Explorer}.  The distances of the two nearer dust clouds are fixed at the distances of the peaks of atomic hydrogen derived from the 21-cm spectrum in the direction of the X-ray source, namely at 1.30 and 2.56 kpc. With these constraints, a good fit is achieved with the source at a distance 4.5 kpc, the distance of the third cloud at 4.05 kpc, the total scattering optical depth of the three clouds equal to 0.159 at 3 keV, and the column density of hydrogen set to $4.6\\times 10^{22}$ cm$^{-2}$. With $A_{\\rm V}=6.5\\pm0.3$~mag for the binary companion star, QV Nor, the ratio of the scattering optical depth at 3 keV to the visual extinction is $0.0234\\pm0.0010$ mag$^{-1}$.  The column density of hydrogen in the model is much greater than the column density of atomic hydrogen derived from the 21-cm spectrum, which indicates that most of the hydrogen is in molecules, probably concentrated in the three dust clouds. ",
        "introduction": "The grain-scattered halos of X-ray sources contain clues to the nature and distribution of interstellar dust grains.  The {\\it Chandra} X-ray observatory provides the means to achieve substantial improvements over previous observations of X-ray halos.  Here we present results from a {\\it Chandra} observation and a Monte Carlo simulation of the halo of the eclipsing X-ray binary pulsar 4U 1538-52 (hereinafter x1538), an especially interesting halo because of its comparatively high intensity, evidence in the shapes of its radial profiles for a concentration of dust in discrete clouds along the line of sight, and high variability caused by eclipses and flares.  The theory of grain-scattered halos of steady sources, first discussed by \\citet{ove65}, has been developed by various authors, e.g. \\citet{sly69,hay70,mar70,mau86,smi98,dra03}.  The effects of intensity variations of a source on the variation of its halo and the possibilities this can present for measuring the distance of the source and the distribution of the scattering grains along the sight line have been discussed by \\citet{tru73,alc78,xu86,klo94}. \\citet{klo91} calculated the expected profiles of halos that would be produced by grains with various assumed size distributions. \\citet{mat91} carried out an analytical study of the possibilities of using halos as diagnostics of grains.  \\citet{pre96} discussed the possibilities and the practical limitations of such diagnostics. \\citet{smi98} showed that the Rayleigh-Gans approximation to the Mie scattering theory, commonly employed in theoretical treatments X-ray halos, yields substantial overestimates of halo intensities for photon energies less than 1 keV.  Early evidence of X-ray halos was reported by \\citet{rol83} and by \\citet{cat83} from images recorded by the {\\it Einstein} observatory. Mauche and collaborators \\citep{mau86, mol86, mau89}, employing the same observatory, found a correlation between the fractional intensity of the halo of an X-ray source and the visual extinction of its optical counterpart that was consistent with the model of solid silicate and graphite grains proposed by \\citet{mat77}. \\citet{pre95} carried out an extensive survey of X-ray halos with the {\\it ROSAT} X-ray observatory. \\citet{wit01} found evidence for a distribution in grain size extending to 2 $\\mu$m and possibly more in a {\\it ROSAT} observation of the halo of Nova Cygni 1992. \\citet{pre00} derived an estimate of the distance to Cyg X-3 by an analysis of its X-ray halo. \\citet{cla94} employed the non-imaging {\\it Ginga} observatory to observe x1538 and its halo.  From a comparison of the counting rates before and after an eclipse of x1538 they derived an estimate of the fractional halo intensity that was substantially less than that expected on the assumption that the visual extinction of its optical companion is caused by solid silicate grains with a specific density near 3. They interpreted the result as consistent with a grain structure in the form of loose aggregates of smaller particles.   Following its discovery in the {\\it Uhuru} galactic plane survey \\citep{gia74}, x1538 was observed to pulse with a period of 529 s and to have an orbital period of 3.73 days with an eclipse lasting about 12 hours \\citep{dav77,bec77}. An accurate position determination with {\\it SAS 3} \\citep{app78} led to identification of the highly reddened BO supergiant, QV Nor, as its companion \\citep{cra78}. The X-ray spectrum has a prominent iron K line \\citep{mak87} and a pulse phase-dependent cyclotron resonance feature at 20 keV \\citep{cla90}. Properties of the binary system derived from various observations are listed in Table 1.  For the present study X-ray images of x1538 and its halo were recorded before and after an eclipse immersion. Radial profiles of simulated halos, generated by a Monte Carlo ray-tracing code, were fitted to the three-peaked time-varying radial profiles of the observed halo by adjustment of the parameters of a model of the source spectrum, the dust grains, and the interstellar photoelectric absorption.  The model parameters that yielded a good fit are compared with the optical extinction of QV Nor and with the column density of atomic hydrogen derived from the spectrum of 21-cm radiation in the direction of the source to derive information about the scattering efficiency of the grains and the total column density of hydrogen.  Section 2 describes the observations and data reduction.  Properties of the halo are presented in section 3. The Monte Carlo simulation code is described briefly in section 4 together with the model and the fitting procedure.  Section 5 is a summary and discussion of results. Section 6 lists conclusions drawn from this study. ",
        "conclusions": "From this study we conclude:  1) Monte Carlo simulation of a grain-scattered X-ray halo can provide an effective means for extracting information about the distance of the source and the properties and distribution of interstellar dust grains in circumstances involving discrete dust clouds, source variability, and multiple scattering.  2) To reduce uncertainty in such a simulation it is essential to have an accurate record of the source spectrum for an extended period before the halo observation begins.  3) The halo of x1538 is an especially interesting object because of its relatively high surface brightness, complex structure, and high variability due to eclipses and flares.  4) The time-varying three-peaked halo profiles of x1538 before and after an eclipse can be fitted by simulated halo profiles generated by a Monte Carlo ray-tracing code operating on a model with the source at 4.5 kpc, three discrete dust clouds of which two are at the distances of the two greatest concentrations of atomic hydrogen derived from the 21-cm spectrum, and a third cloud closer to the source.  5) The total scattering optical depth of the three dust clouds in the model is 0.159 for $E=3$ keV.  The ratio of this quantity to the visual extinction of the optical companion, QV Nor, is $0.0234\\pm0.0010$~mag$^{-1}$, somewhat less than the theoretical value of 0.0272 derived for solid spherical carbonaceous and silicate grains by \\citet{dra03}."
    },
    "0404/astro-ph0404392_arXiv.txt": {
        "abstract": "We describe the morphological and fluctuation properties of the pulsars in the double neutron star system, PSR J0737--3039. Pulsar \\B is seen in almost all orbital phases, except in the range of $\\sim 6\\deg$ to $65\\deg$. This may be interpreted as an {\\it eclipse} of \\pBn's signal by its own magnetopause region produced by interaction with \\pAn's relativistic wind. No modulation of the emission of \\pB is found at the period of \\pAn. This places a constraint on the models that propose that \\pAn's beamed radiation is directly responsible for \\pBn's emission. Modulation index values indicate that the pulse to pulse variations in the two objects are mostly intrinsic. Pulsar A shows significant differential modulation index within its pulse profile. ",
        "introduction": "\\label{sec-intro} Recent discovery of the double pulsar system, reported by Burgay et al. (2003) and Lyne et al. (2004) is one of the most important discoveries in the history of pulsar astronomy. This high inclination system, which consists of a short period recycled pulsar and an ordinary slow pulsar revolving with an orbital period 2.45 hrs, gives us almost an ideal situation to probe many fundamental properties of neutron star astrophysics.  The distance between the two pulsars is so small (2.8 light sec), that the light cylinder volume (magnetosphere of a solitary pulsar) of PSR J0737--3039B (hereafter \\Bn) subtends about 20 degrees from PSR J0737--3039A (hereafter \\An).  It is unclear whether the measured value of $\\dot{P}$ of \\pB is intrinsic to the pulsar, given the mechanical luminosity ($\\dot{E}$) of \\pAn. It is conceivable that the relativistic wind from \\pA influences the characteristics of \\pBn.  In spite of this, \\pB continues to emit radio radiation, clearly demonstrating the robustness of the fundamental radio emission mechanism.  According to Lyne et al. (2004), this system may well have indicated that the seed acceleration responsible for this emission comes from deep inside the magnetosphere, clearly ruling out outer magnetospheric emission models.  The measured flux of \\pB significantly changes as a function of orbital phase\\footnote{Orbital phase ($\\phorb$) is given as $(\\omega+\\eta)$, where $\\omega$ and $\\eta$ are the position angle of periastron and the ``true'' anomaly.} (Lyne et al. 2004).  Models have been proposed to suggest that the wind from \\pA creates a magnetopause region around \\pB due to the interaction with \\pBn's magnetosphere (Arons et al. 2004). In this model, the synchrotron absorption optical depth of this magnetopause region is the primary reason for the observed eclipse of \\pAn, and possibly for the light curve of \\pBn.  In this work, we present some of the basic morphological and fluctuation properties of these two pulsars. Section \\ref{sec-flux} describes the properties of average pulses of \\pA and \\pBn. The fluctuation properties of the two pulsars are described in Section \\ref{sec-fluct}. Implications of our results are discussed in Section \\ref{sec-discuss}. ",
        "conclusions": "\\label{sec-discuss} The measured flux from \\pB shows systematic variation as a function of orbital phase. The pulsar is visible is all phase ranges except in the range of $\\phi_{\\rm orb}\\sim 6\\deg$ to $65\\deg$, where we have obtained only an upper limit to the flux. This void region may be interpreted as an {\\it eclipse} of \\pB by its own magnetopause region which is produced by interaction with the relativistic wind from \\pAn. It is this magnetopause region that may also be responsible for the eclipse of \\pA around its superior conjunction (Arons et al. 2004). Although it is difficult for us to produce reliable light curve around this \\pBn's apparent {\\it eclipse} owing to signal-to-noise considerations, it is clear that the drop in flux is very rapid (within several pulses!)  With our observations at 800 and 1400 MHz, we see some indication of frequency dependence in \\pBn's light curve. During windows III and IV, the flux at 1400 MHz seems to be systematically lower than the value at 800 MHz. This needs to be validated with further high signal-to-noise observations.  Modulation index measurements indicate variation of single pulse flux in both these objects. These variations are significantly more than what is expected from the measured decorrelation bandwidth (due to interstellar scattering) of this system. Therefore, these variations are intrinsic to the two pulsars. The inner regions in \\pA seems to have significantly higher modulation index than the brighter outer regions.  Pulsar \\A exhibits a complex evolution of its average profile as a function of radio frequency. However, these properties do not seem to be different from the characteristics that short period pulsar population seems to show.  Pulsar \\Bn's profile shows dramatic variations as a function of orbital phase, making it impossible to solve for its geometric orientation without sensitive polarization measurements. The profile evolution as a function of frequency seems to reveal mainly a double component profile, where the components become more pronounced at higher frequencies. This is quite the opposite of what one would expect from simple dipole geometry, combined with radius to frequency mapping.  Given the possible geometrical orientation of \\pA derived on the basis of polarimetric observations, it appears that there will be an orbital phase dependent wind pressure on \\B (Demorest et al. 2004). It is also plausible that the diameter of \\pBn's emission cone changes due to the effect of orbital phase dependent pressure, there by changing the fractional value of impact angle ($\\beta$) with respect to the radius of the emission cone ($\\rho$). This can introduce complex profile changes as a function of orbital phase, as we see for \\pBn. It is very important to understand the geometrical orientation of \\pBn, for which sensitive polarimetric observations are necessary.  Within our sensitivity limits, there is no significant intermodulation between signals from the two objects. We do not detect any significant feature in the L-R \\& H-R spectra, indicating that the spectral characteristics of the intrinsic pulse fluctuations is ``noiselike''. This constrains the models proposing direct influence on the emission mechanism of \\pB by radiation from \\pAn."
    },
    "0404/astro-ph0404447_arXiv.txt": {
        "abstract": "We have tabulated lists of upper red giant, horizontal, and asymptotic giant branch (RGB, HB, AGB) stars in the globular cluster M5 that are complete to over $10\\arcmin$ from the core for the RGB and AGB samples, and $8\\arcmin$ for the HB sample. The large samples give us the most precise value of $R_2 = N_{AGB} / N_{HB}$ to date for a single globular cluster ($0.176 \\pm 0.018$). This is incompatible with theoretical calculations using the most recent physical inputs. The discrepancy can probably be attributed to the dependence of observed \\rat values on horizontal branch morphology.  We identify the cluster M55 as being another possible example of this effect. Samples of HB and AGB stars in populous clusters may provide a means of calibrating the masses of horizontal branch stars in globular clusters. The cumulative luminosity function of the upper red giant branch shows an apparent deficit of observed stars near the tip of the branch. This feature has less than a 2\\% chance of being due to statistical fluctuations.  The slope of the cumulative luminosity function for AGB stars is consistent with the theoretically predicted value from stellar models when measurement bias is taken into account. We also introduce a new diagnostic $R_{clump}$ that reflects the fraction of the AGB lifetime that a star spends in the AGB clump. For M5, we find $R_{clump} = 0.42 \\pm 0.05$, in marginal disagreement with theoretical predictions. Finally, we note that the blue half of M5's instability strip (where first overtone RR Lyraes reside) is underpopulated, based on the large numbers of fundamental mode RR Lyraes and on the nonvariable stars at the blue end of the instability strip. This fact may imply that the evolutionary tracks (and particularly the colors) of stars in the instability strip are affected by pulsations. ",
        "introduction": "Evolved-star populations in clusters have long provided important tests of stellar evolution theory because they present us with coeval ensembles of stars having nearly identical chemical compositions. Because the ratio of the numbers of stars in two different evolved phases relates to the the ratio of the intervals of time that a star spends in the two phases, we can get a glimpse at the stellar evolution clock. However, it is only in the most massive clusters that we can expect to sample, and therefore test model timescales for, stars in the shortest evolutionary phases. Even in the most massive clusters, the huge range of stellar densities requires that wide-field data be taken to cover the extended halo of the cluster, while high spatial-resolution (often space-based) photometry is gathered for the crowded cores.  This study is an analysis of evolved stars in the massive Milky Way globular cluster M5. We have compiled large samples of relatively rare populations on the upper red giant branch (RGB) and the asymptotic giant branch (AGB). Our intent is to test whether the observed evolution conforms to our expectations from stellar models with current physics inputs. In \\S 2, we describe the observational material we used in identifying and the stars, and the procedures used to classify them based on their photometric properties. In \\S 3, we describe various diagnostics of the bright populations of globular clusters, including the population ratio \\rat, the distribution of HB stars, cumulative luminosity functions, and a new population ratio $R_{clump}$. ",
        "conclusions": "We have compiled a nearly complete list of bright RGB, HB, and AGB stars for the globular cluster M5 reaching from the core of the cluster to $8 - 10 \\arcmin$ from the center. We have used these samples to conduct a thorough comparison with theory in order to test stellar interior physics under the conditions prevalent in these bright stars. We have introduced a new diagnostic $R_{clump}$ for evaluating the evolutionary timescale in the early part of the AGB phase, finding that there is a marginal disagreement between the observations and the models of \\citet{dro}. On the whole we find good agreement between observations and theory for the cumulative LFs of the RGB above the RGB bump and for the AGB, indicating that the evolutionary timescales in both phases are predicted accurately. An apparent exception appears near the tip of the RGB, where there appear to be too few giants compared to theoretical predictions. The chance that this is a statistical fluctuation is less than 2\\%. This may be an indication that the neutrino emission rates in the cores of these bright stars are underestimated.  The most significant result is the large value for the population ratio \\rat compared to recent theoretical values. The high value is probably the result of the particular HB morphology of the cluster and not continuing uncertainties in physical inputs like the cross section for the \\cag reaction or the core mixing algorithm. We encourage new calculations of HB and AGB phases using updated physics because AGB stars provide a means of testing predictions for the morphology of evolutionary tracks for stars fusing He into C and O and of constraining the \\cag reaction rate. In particular, they stand the greatest chance of identifying the color at which HB stars change from going into or not going into a typical AGB phase, and the range of colors for which the AGB phase has its maximum duration.  We also call attention to a peculiarity of the distribution of stars on M5's HB. We find that the mass distribution for HB stars peaks at a position corresponding to the blue edge of the instability strip. The instability strip is heavily populated due to the large dispersion in HB masses ($0.02 - 0.03 \\msun$). However, the distribution of stars within the instability strip is heavily biased toward the red half of the instability strip. This is consistent with M5 being an Oosterhoff group I cluster, but it means that the first overtone instability strip is underpopulated compared to the fundamental strip to the red and compared to the nonvariable stars to the blue. Because M5 has one of the bluest HB morphologies of the Oo I clusters (including the more metal-poor M3), it is a severe test of potential explanations of the Oosterhoff dichotomy. Some mechanism for reducing the evolutionary timescale of stars in the first overtone instability strip seems to be needed. The blue half of M5's instability strip {\\it should be} more heavily populated --- the big question is why it isn't.  \\newpage"
    },
    "0404/nlin0404043.txt": {
        "abstract": "In this review article we will describe recent developments in statistical theory of magnetohydrodynamic (MHD) turbulence. Kraichnan and Iroshnikov first proposed a phenomenology of MHD turbulence where Alfv\\'{e}n time-scale dominates the dynamics, and the energy spectrum $E(k)$ is proportional to $k^{-3/2}$. In the last decade, many numerical simulations show that spectral index is closer to 5/3, which is Kolmogorov's index for fluid turbulence. We review recent theoretical results based on anisotropy and Renormalization Groups which support Kolmogorov's scaling for MHD turbulence.  Energy transfer among Fourier modes, energy flux, and shell-to-shell energy transfers are important quantities in MHD turbulence. We report recent numerical and field-theoretic results in this area. Role of these quantities in magnetic field amplification (dynamo) are also discussed. There are new insights into the role of magnetic helicity in turbulence evolution. Recent interesting results in intermittency, large-eddy simulations, and shell models of magnetohydrodynamics are also covered. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404501_arXiv.txt": {
        "abstract": "Astrophysical gamma-ray spectroscopy is a most valuable tool for studying nuclear astrophysics, as well as recent star formation  in the Milky Way. After a short, historical, introduction to the field, I present a brief review of the most important current issues. Emphasis is given to radioactivities produced by massive stars and associated supernova explosions, and in particular, those related to observations presently carried out by INTEGRAL: short-lived \\Ti \\ from CasA and SN1987A and long-lived \\Al \\ and \\Fh \\ from massive stars; various candidate sources of positrons for the 511 keV emission of the Galactic bulge are also critically discussed. ",
        "introduction": "The starting point of $\\gamma$-ray line astronomy with cosmic radioactivities is usually considered to be the landmark paper of Clayton, Colgate and Fishman (1969). That work clarified the implications of the production of \\Ni \\ (a doubly magic, and yet unstable nucleus) during explosive Si-burning in supernovae (SN). In particular, it opened exciting perspectives for $\\gamma$-ray line astronomy, by suggesting that any supernova within the local group of galaxies woud be detectable in the characteristic $\\gamma$-ray lines resulting from the radioactive decay of \\Ni \\ and its daughter nucleus \\Co.  In the 70's D. Clayton identified most of the radionuclides of astrophysical interest (i.e. giving a detectable $\\gamma$-ray line signal); for that purpose, he evaluated their average SN yields by assuming that the corresponding daughter stable nuclei are produced in their solar system abundances. Amazingly enough (or naturally enough, depending on one's point of view) his predictions of average SN radionuclide yields (Table 2 in Clayton 1982) are in excellent agreement with modern yield calculations, based on full stellar models and detailed nuclear physics (see Fig. 1). Only the importance of \\Al \\ escaped Clayton's (1982) attention, perhaps because its daughter nucleus $^{26}$Mg is produced in its stable form, making the evaluation of the parent's yield quite uncertain. That uncertainty did not prevent Arnett (1977) and Ramaty and Lingenfelter (1977) from arguing (on the basis of Arnett's (1969) explosive nucleosynthesis calculations) that, even if only 10$^{-3}$ of solar $^{26}$Mg is produced as \\Al, the resulting Galactic flux from tens of thousands of supernovae (during the $\\sim$1 Myr lifetime of \\Al) would be of the order of 10$^{-4}$ \\cs.   \\begin{figure*} \\centering \\includegraphics[height=0.67\\textwidth,angle=-90]{Prantzos_f0.eps} \\caption{Yields of radioactive nuclei from massive exploding stars. {\\it On the left},  Clayton's (1982) predictions for average SN yields are given (based on the assumption that the corresponding stable daughter isotopes are produced in their solar abundances as radioactive progenitors); the uncertainty in \\Fh \\ yield stems from the unknown percentage of its contribution to the production of stable $^{60}$Ni (which may also be produced as unstable $^{60}$Co; the limits of the error bar correspond to 0 percent and to 100 percent contribution, respectively). {\\it In the middle}, recent theoretical results are plotted as a function of stellar mass ({\\it filled symbols}, from Rauscher et al. 2002), while {\\it open parallelograms} indicate observational constraints from SN1987A, a 18-20 \\ms \\ star. {\\it On the right} appear other observational constraints, either  from CasA (for \\Ti) or from extragalactic SN (for \\Ni, see also Fig. 3). } \\end{figure*}    In the case of \\Al \\  nature appeared quite generous, providing a \\ga-ray flux even larger than the optimistic estimates of Ramaty and Lingenfelter (1977): the HEAO-3 satellite detected the corresponding 1.8 MeV line from the Galactic  center direction at a level of 4 10$^{-4}$ \\cs (Mahoney et al. 1984). That detection, the first ever of a cosmic radioactivity, showed that nucleosynthesis is still active in the Milky Way; however, the implied large amount of galactic \\Al \\ ($\\sim$2 \\ms \\ per Myr, assuming steady state) was difficult to accomodate in conventional models of galactic chemical evolution if SN were the main \\Al \\ source (Clayton 1984), since $^{27}$Al would be overproduced in that case; however, if the  ``closed box model'' assumption is dropped and {\\it infall} is assumed in the chemical evolution model, that difficulty is removed, as subsequently shown by Clayton and Leising (1987).  Another welcome mini-surprise came a few years later, when the \\Co \\ \\ga-ray lines were detected in the supernova SN1987A, a $\\sim$20 \\ms \\ star that exploded in the Large Magellanic Cloud. On theoretical  grounds, it was expected that a SNIa (exploding white dwarf of $\\sim$1.4 \\ms \\ that produces $\\sim$0.7 \\ms \\ of \\Ni) would be the first to be detected in \\ga-ray lines; indeed, the large envelope mass of massive exploding stars  ($\\sim$10 \\ms) allows only small  amounts of \\ga-rays to leak out of SNII, making  the detectability of such objects problematic. Despite the intrinsically weak \\ga-ray line emissivity of SN1987A, the proximity of LMC allowed the first detection of the tell-tale \\ga-ray line signature from the  radioactive chain  \\Ni\\ra\\Co\\ra\\Fe; this confirmed a 20-year old conjecture, namely that the abundant \\Fe \\ is produced in the form of radioactive \\Ni.  Those discoveries laid the observational foundations of the field of \\ga-ray line astronomy with radioactivities. The next steps were made in the 90ies, thanks to the contributions of the Compton Gamma-Ray Observatory (CGRO). First, the {\\it OSSE} instrument aboard CGRO detected the \\Ci \\ \\ga-ray lines from SN1987A (Kurfess et al. 1992); the determination of the abundance ratio of the isotopes with mass numbers 56 and 57 offered a unique probe of the physical conditions in the innermost layers of the supernova, where those isotopes are synthesized (Clayton et al. 1992). On the other hand, the {\\it {\\it COMPTEL}} \\ instrument mapped the Miky Way in the light of the 1.8 MeV line and found irregular emission along the plane of the Milky Way and prominent ``hot-spots'' in directions approximately tangent to the spiral arms (Diehl et al. 1995); that map implies that massive stars (SNII and/or WR) are at the origin of galactic \\Al \\ (as suggested by Prantzos 1991, 1993) and not an old stellar population like novae or AGB stars.  Furthermore, {\\it {\\it COMPTEL}} \\  detected the 1.16 MeV line of radioactive \\Ti \\ in the Cas-A supernova remnant (Iyudin et al 1994). That discovery offered another valuable estimate of the yield of a radioactive isotope produced in a massive star explosion (although, in that case the progenitor star mass is not known, contrary to the case of SN1987A). On the other hand, it also created some new problems, since current 1D models of core collapse supernova do not seem able to account for the yield inferred from the observations; however, recent axisymmetric models of rotating star explosions offer interesting perspectives in that respect (see next section).   \\begin{table*}[t] \\caption{Important stellar radioactivities for gamma-ray line astronomy \\label{fig:Table1}} \\vskip -2.8 cm \\includegraphics[width=\\textwidth,height=19.cm]{Prantzos_f1.eps} \\vskip -3.4 cm \\end{table*}   Finally, after CGRO and before INTEGRAL, another important discovery was made in the field: the RHESSI experiment confirmed the detection of \\Al \\ in the Galaxy, but it also detected the characteristic decay lines of \\Fh, another long-lived isotope with $\\sim$2 Myr period (Smith 2004, and this volume). The detected \\Fh \\ emissivity is, however,  considerably  lower than expected from current models of nucleosynthesis in core collapse SN. The implications of that discovery are studied in Prantzos (2004) and in Sec. 4.  Table 1 summarizes our current picture of stellar radioactivities and associated $\\gamma$-ray lines. In the following we shall focus on the radioactivities produced by massive stars and associated supernova explosions, and in particular, those related to observations presently carried out by INTEGRAL. Similar reviews have been presented by Diehl and Timmes (1998) and Kn\\\"odlseder and Vedrenne (2001). The current status of the physics of supernova explosions is reviewed by Hillebrandt (this volume). ",
        "conclusions": ""
    },
    "0404/astro-ph0404209_arXiv.txt": {
        "abstract": "This is the second in a series of papers on the effects of dust on multifluid, magnetohydrodynamic shock waves in weakly-ionized molecular gas. We investigate the influence of dust on the critical shock speed, \\vcrit, above which C shocks cease to exist. Chernoff showed that \\vcrit\\ cannot exceed the grain magnetosound speed, \\vgms, if dust grains are dynamically well coupled to the magnetic field. Since $\\vgms \\simeq 5$\\,\\kms\\ in a dense cloud or core, the potential implications for models of shock emission are profound. We present numerical simulations of steady shocks where the grains may be well- or poorly coupled to the field. We use a time-dependent, multifluid MHD code that models the plasma as a system of interacting fluids: neutral particles, ions, electrons, and various ``dust fluids'' comprised of grains with different sizes and charges. Our simulations include grain inertia and grain charge fluctuations but to highlight the essential physics we assume adiabatic flow, single-size grains, and neglect the effects of chemistry. We show that the existence of a phase speed $\\vphi$ does not necessarily mean that C shocks will form for all shock speeds $\\vs$ less than $\\vphi$. When the grains are weakly coupled to the field, steady, adiabatic shocks resemble shocks with no dust: the transition to J~type flow occurs at $\\vcrit\\approx 2.76\\vnA$, where \\vnA\\ is the neutral \\Alf\\ speed, and steady shocks with $\\vs>2.76\\vnA$ are J~shocks with magnetic precursors in the ion-electron fluid. When the grains are strongly coupled to the field, $\\vcrit=\\min\\left(2.76\\vnA,\\vgms\\right)$. Shocks with $\\vcrit<\\vs<\\vgms$ have magnetic precursors in the ion-electron-dust fluid. Shocks with $\\vs>\\vgms$ have no magnetic precursor in any fluid. We present time-dependent calculations to study the formation of steady multifluid shocks. The dynamics differ qualitatively depending on whether or not the grains and field are well coupled. ",
        "introduction": "Shock waves in weakly-ionized plasmas, such as the molecular outflows around young stellar objects, may have a multifluid structure, in which the charged and neutral components of the plasma behave as distinct fluids. If the shock speed is less than the speed with which the charged fluid communicates compressive disturbances, the charged particles are accelerated ahead of the neutral gas in a ``magnetic precursor'' (Mullan 1971; Draine 1980). The streaming of charged particles through the neutrals also accelerates, compresses and heats the neutral fluid. If it remains sufficiently cool, the neutral fluid is everywhere supersonic in a frame at rest in the post-shock gas. In this case all hydrodynamical variables are continuous and the entire structure is termed a ``C~shock'' (Draine 1980). Alternatively the neutral fluid must undergo a transonic transition. This may occur smoothly via continous flow through a sonic point (a ``C$^*$ shock'', see Chernoff 1987 and Roberge \\& Draine 1990, but note that Chi\\`{e}ze, Pineau des For\\^{e}ts, \\& Flower 1998 question the reality of C$^*$ shocks) or abruptly at a viscous jump front (a ``J shock'').  Simulations of the different types of multifluid MHD shocks reveal distinct observational signatures, including the strengths and velocity profiles of atomic and molecular emission lines (Chernoff, Hollenbach, \\& McKee 1982; Draine \\& Roberge 1982; Draine, Roberge, \\& Dalgarno 1983, henceforth DRD83; Smith \\& Brand 1990a,b; Smith 1991a,b; Smith, Brand, \\& Moorhouse 1991a,b; Kaufman \\& Neufeld 1996a,b; Timmermann 1996; Neufeld \\& Stone 1997; Pineau des For\\^{e}ts et al.\\ 2001), the ortho/para ratio of H$_2$ (Timmermann 1998, Wilgenbus et al.\\ 2000), and the abundances of various chemical species in the post-shock gas (e.g., DRD83; Flower, Pineau des For\\^{e}ts, \\& Hartquist 1985; Pineau des For\\^{e}ts, Flower, \\& Dalgarno 1988; Flower \\& Pineau des For\\^{e}ts 1994; Flower et al.\\ 1996; Bergin, Neufeld, \\& Melnick 1998).  When the simulations are used to model the rich array of spectral observations now available for molecular outflows, a combination of C- and J shocks is broadly satisfactory, but an acceptable explanation for important aspects of the observations remains elusive (e.g., Smith, Eisl\\\"{o}ffel, \\& Davis 1998; Rosenthal, Bertoldi, \\& Drapatz 2000). Consequently, the theory of multifluid shocks remains of considerable interest.  In this paper we reexamine a fundamental prediction of the theory, namely, the ``critical speed,'' \\vcrit, above which C type solutions fail to exist. According to most prior work, the transition from C- to J-type solutions occurs when the neutral flow becomes too hot to remain subsonic. For shocks in dense molecular gas, the transition is caused either by the onset of \\Htwo\\ dissociation (which decreases radiative cooling of the neutral gas) or by self-ionization (which increases heating by ion-neutral scattering). The precise value of \\vcrit\\ is sensitive to the rates of various atomic and molecular processes. DRD83 estimated that $49 \\ge \\vcrit \\ge 40$\\,\\kms\\ for preshock densities in the range $10^4 \\le \\nH \\le 10^6$\\,\\cmMMM. Analogous calculations by Smith (1994) give similar values, though recent simulations by Le~Bourlot et al.\\ (2002) predict somewhat larger critical speeds (e.g., $\\vcrit \\approx 70$\\,\\kms\\ for $\\nH=10^4$\\,\\cmMMM). In any case, shocks in dense molecular gas would be C~type for speeds $\\gtrsim$40--50\\,\\kms\\ if the critical speed was determined solely by the temperature of the neutrals.  However \\vcrit\\ is also constrained by the ``signal speed'' for compressive waves in the charged fluid, since shocks faster than the signal speed must necessarily be J type. In a cold, weakly-ionized molecular gas, $\\vsig \\simeq B/\\sqrt{4\\pi\\rhoc}$, where $B$ is the magnetic field strength and \\rhoc\\ is the mass density of charged particles. If \\rhoc\\ is taken to be the density of atomic and molecular ions, then $\\vsig=\\viA$, where $\\viA\\simgt 1000$\\,\\kms\\ is the ion \\Alf\\ speed. In this view the constraint imposed by the signal speed would be of no practical significance. However dust grains are also (mostly) charged  (e.g., see Fig. 1 of Paper I) and, to the extent that the grains are ``loaded'' onto the field lines, their contribution to \\rhoc\\ also should be included (McKee, Chernoff, \\& Hollenbach 1984). For example, if \\rhoc\\ is taken to be the {\\it total}\\/ dust density, then $\\vsig=\\vgms$, where $\\vgms\\approx 5$\\,\\kms\\ (see eq. [\\ref{vgmsexpresseq}]) is the grain magnetosound speed. A critical speed this small would have profound implications for models of shocked molecular outflows. Flower \\& Pineau des For\\^{e}ts (2003) have computed the abundance of charged grains in clouds including polycyclic aromatic hydrocarbons (PAHs) and cosmic-ray-induced UV photons. They suggest that only a small fraction of the total grain mass in a plasma containing PAHs is charged and coupled to the field, generally yielding magnetosound speeds $\\sim 50~\\kms$. However, in the environments we wish to study, namely, dense, dark clouds and cores, it is likely that most PAHs have condensed or frozen onto the icy mantles of larger grains (Ehrenfreund \\& Charnley 2000; Bernstein et al.  2003; Gudipati \\& Allamondola 2003; Kr\\\"{u}gel 2003), and therefore should not affect the charge state of grains to such a large degree.  The critical speed $\\vcrit$ would equal the grain magnetosound speed $\\vgms$ only if all the grains are well coupled to the magnetic field lines (Chernoff 1987). However, the grains may be only weakly coupled (e.g., Pilipp et al. 1990; Draine \\& McKee 1993) and grain charge fluctuations limit the effective mass of the loaded grains to some fraction of the total dust mass (Ciolek \\& Mouschovias 1993). Consequently, detailed numerical simulations are generally required to determine \\vcrit. Although dust has been included in many detailed studies of shocks, most have neglected the inertia of the grains (i.e., the momentum equation for the dust has been omitted, on the assumption that the dust motion is force free). Because propagating waves in the grain fluid are suppressed when the grain inertia is neglected, studies of this type cannot be used to examine the connection between grain loading and \\vcrit.  In this paper we present numerical calculations which include the grain inertia, and extend Chernoff's work by also including partial grain-field coupling and grain charge fluctuations. Our plan is as follows: In \\S~2 we summarize our assumptions and the free parameters in our models. The signal speed in a dusty plasma is discussed in \\S~3, where we examine the effects of dust on the phase velocities and damping rates of compressive waves. In \\S~4 we present numerical solutions for steady and time-dependent shock waves. Our results are summarized in \\S~5.  \\vspace{-4ex} ",
        "conclusions": "The principal results of this paper are as follows:  \\begin{enumerate}  \\item We calculated the dispersion relations for compressive waves in a dusty plasma for two models, where the dust is weakly coupled and strongly coupled to the magnetic field, respectively. In the ``weakly coupled'' (WC) model, the only propagating modes in the charged fluids are ion magnetosound waves at very short wavelengths and neutral magnetosound waves at very long wavelengths. In the ``strongly coupled'' (SC) model, grain magnetosound waves appear at intermediate wavelengths.  \\item We simulated the evolution of small-amplitude, Gaussian wave packets to see how different plasma components respond to disturbances on various length scales. When the length scales are small enough for decoupling of the charged and neutral particles to occur, the dynamics are dramatically different in the WC and SC models. In the WC model there is no wave propagation as such. Propagating ion magnetosound waves are allowed but strongly damped, and so have virtually no affect on the dynamics. Rapid diffusion of ions and electrons through the neutral gas and dust occurs, and the disturbance decays without propagating. In the SC model, packets of grain magnetosound waves appear, and the disturbance propagates (with weak damping) at the grain magnetosound speed. Because the ions and electrons are tied to the magnetic field, the coupling of grains and magnetic field in the SC model suppresses diffusive ion motions.  \\item We simulated steady, adiabatic shock waves to compare the effects of dust in the WC and SC models. The WC model gives shocks that are virtually identical to shocks with no dust: steady shocks are C type for \\Alf\\ Mach numbers $\\Mna<2.76$ and J type with magnetic precursors in the ion-electron fluid for $\\Mna>2.76$. We say that steady, adiabatic C shocks are ``cooling limited'' in the WC model because the C-J transition occurs when the neutrals become too hot to remain everywhere supersonic.  \\item Steady, adiabatic C shocks are also cooling limited in the SC model if $\\vgms>2.76\\vnA$. Their structures resemble their WC counterparts, with differences attributable to additional momentum transfer caused by grain-neutral scattering.  \\item In the SC model steady, adiabatic C shocks are ``signal limited\" if $\\vgms<2.76\\vnA$. In this case the C-J transition occurs when the shock speed exceeds the grain magnetosound speed, as predicted by Chernoff (1985). Shocks faster than the grain magnetosound speed have no magnetic precursor in any fluid.  \\item We simulated time-dependent shocks to see how magnetic precursors form in the WC and SC models. Consistent with our study of small-amplitude disturbances, the dynamics are dramatically different in the two cases. In the WC case, a magnetic precursor forms via diffusion of the ion-electron plasma through the dust and neutrals. In the SC case, the precursor forms via the propagation of grain magnetosound waves. \\end{enumerate} Our calculations incorporated some very unrealistic assumptions in order to isolate the essential physics: Real grains have a spectrum of sizes and real shocks have chemistry and radiative cooling. Having understood the physics, our next investigation will determine realistic values of \\vcrit\\ by relaxing these assumptions."
    },
    "0404/astro-ph0404453_arXiv.txt": {
        "abstract": "The Burst and Transient Source Experiment (BATSE), aboard the {\\it Compton Gamma Ray Observatory} (CGRO), provided a record of the low-energy gamma-ray sky ($\\sim$20-1000 keV) between 1991 April and 2000 May (9.1y). BATSE monitored the high energy sky using the Earth occultation technique (EOT) for point sources whose emission extended for times on the order of the CGRO orbital period ($\\sim$92 min) or greater.  Using the EOT to extract flux information, a catalog of sources using data from the BATSE large area detectors has been prepared. The first part of the catalog consists of results from the all-sky monitoring of 58 sources, mostly Galactic, with intrinsic variability on timescales of hours to years.  For these sources, we have included tables of flux and spectral data, and outburst times for transients. Light curves (or flux histories) have been placed on the world wide web.  We then performed a deep-sampling of these 58 objects, plus a selection of 121 more objects, combining data from the entire 9.1y BATSE dataset.  Source types considered were primarily accreting binaries, but a small number of representative active galaxies, X-ray-emitting stars, and supernova remnants were also included. The sample represents a compilation of sources monitored and/or discovered with BATSE and other high energy instruments between 1991 and 2000, known sources taken from the HEAO 1 A-4 \\citep{Levine84} and \\citet{MacombGehrels99} catalogs. The deep sample results include definite detections of 83 objects and possible detections of 36 additional objects.  The definite detections spanned three classes of sources: accreting black hole and neutron star binaries, active galaxies and supernova remnants. The average fluxes measured for the fourth class, the X-ray emitting stars, were below the confidence limit for definite detection.  Flux data for the deep sample are presented in four energy bands: 20-40, 40-70, 70-160, and 160-430 keV.  The limiting average flux level (9.1 y) for the sample varies from 3.5 to 20 mCrab (5$\\sigma$) between 20 and 430 keV, depending on systematic error, which in turn is primarily dependent on the sky location. To strengthen the credibility of detection of weaker sources ($\\sim$5-25 mCrab), we generated Earth occultation images, searched for periodic behavior using FFT and epoch folding methods, and critically evaluated the energy-dependent emission in the four flux bands. The deep sample results are intended for guidance in performing future all-sky surveys or pointed observations in the hard X-ray and low-energy gamma-ray band, as well as more detailed studies with the BATSE EOT. ",
        "introduction": "The {\\it Compton Gamma Ray Observatory} (CGRO)\\footnote{Acronyms and abbreviations used throughout the text and tables are listed in Appendix A.}, one of the four missions in NASA's ``Great Observatories\" series, was launched in 1991 April and operated in low Earth orbit until its controlled re-entry in 2000 June.  CGRO was responsible for many discoveries in the study of gamma-ray bursts, accreting binaries, active galaxies, and pulsars \\citep[see general reviews by][]{GehrelsShrader97, KniffenGehrels97, LeonardWanjek00}.  The quest for the origin of gamma-ray bursts \\citep{Meegan92, FishmanMeegan95} led to the development and flight of the Burst and Transient Source Experiment (BATSE) on the CGRO. In addition to BATSE's primary science goals, its nine years of nearly continuous operation and all-sky capability allowed monitoring of the low-energy gamma-ray sky using the Earth occultation technique (EOT). Several new objects were discovered in the BATSE dataset with the EOT, and numerous outbursts of transient and periodic sources were detected during the mission.  Here we present a compilation of BATSE EOT observations that can be used as the basis for more detailed studies on individual sources, for which it is important to have a knowledge of the long-lived emission in the BATSE energy range. It is our hope that these results will stimulate and guide future observations of the low-energy gamma-ray sky. In this work, we performed a careful assessment of systematic error in deriving results for the entire mission. The intent was to push the instrumental and technical limits of the present application of the BATSE Earth occultation technique (EOT). These data are made available on-line to other investigators through the High Energy Astrophysics Science Archive Research Center (HEASARC) at NASA Goddard Space Flight Center, the web site of the BATSE Earth Occultation Team at the NASA Marshall Space Flight Center, and the Compton Observatory Science Support Center (COSSC).  To make the presentation of this work tractable, we divided the catalog into two parts.  First, we emphasized the all-sky monitoring capability of BATSE. For the bright sources that vary on timescales greater than a few hours, we give an overview of activity and representative flux and spectral behavior.  We also briefly make note of the science investigations performed with the EOT during the mission. Second, we performed a ``sky survey\" of a comprehensive set of low-energy gamma-ray sources. We note that this is not a true survey, in that it is biased by preselection of sources. The emphasis is on obtaining good quality results along the Galactic plane where accreting sources dominate the low-energy gamma-ray sky.   We note that other missions with all-sky or wide-field coverage were operational between 1991 and 2000. A non-exhaustive list includes the Japanese {\\it Ginga} satellite, with its All-Sky Monitor, functioning between 1987 February and 1991 October \\citep{Tsunemi89}, and covering the bandpass of 1$-$20 keV, the scanning and imaging instruments aboard the French-Russian spacecraft {\\it Granat} \\citep{Brandt90, Roques90, Sunyaev90, Paul91}, which operated between 1989 December and 1997 October, and the U.S. {\\it Rossi X-Ray Timing Explorer} (RXTE) mission, with its All-Sky Monitor (2-12 keV) \\citep{Bradt93, Levine96}, which began operation in 1995 December and remains operational. {\\it Gra\\-nat}/Sig\\-ma \\citep{Roques90}, a coded mask, sodium iodide imaging detector system, provided considerable coverage and detailed images of the central region of the Galaxy in the 35$-$1300 keV bandpass. There were also scanning instruments of common design on board {\\it Granat} and Danish {\\it EURECA} spacecraft known as WATCH \\citep{Lund86, Brandt94, Castro94} that monitored a portion of the X-ray sky in the 5-120 keV band.  The powerful Wide Field Cameras \\citep{Jager97} on the Italian-Dutch {\\it BeppoSax} satellite (1.8-28 keV) provided coded aperture imaging for both surveying and monitoring of large areas of the sky. BATSE complemented these instruments by full-sky monitoring of the energy range from 20 keV up to about 1 MeV.  Much potential remains for investigations using combinations of these datasets.   BATSE served as an all-sky monitor using the Earth occultation technique \\citep{Harmon02} between 1991 April 23 and 2000 May 28 ($\\sim$9.1y).  This capability was used to trigger numerous pointed observations with the CGRO and other observatories. We extracted Earth occultation measurements from the BATSE dataset at a maximum rate of twice per 92m orbit of CGRO in sixteen energy channels between 20 keV and 1 MeV.  Intensity data for any part of the sky, subject to the geometric constraints of occultation, were available for the duration of the mission with the exception of a few brief intervals (see Appendix B).  Although the BATSE EOT was limited in sensitivity by high backgrounds and modest spatial resolution ($\\sim$1$^{\\circ}$), we could investigate intensity and spectral variations over a dynamical range from days to years. Although $\\sim$daily timescales have been accessible with pointed instruments, BATSE was the first all-sky monitor since those of {\\it Vela 5B} \\citep{Conner69} and {\\it Ariel 5} \\citep{Holt76}, to cover the upper end of this time range.   On the other hand, no low-energy gamma-ray full-sky surveys have been performed since the pioneering HEAO 1 A-4 experiment in 1977-79 \\citep{Levine84}.  That experiment covered the energy range between 13 and 180 keV, and resulted in the detection of about 70 sources (7 of them extragalactic). Such a survey with BATSE has been attempted by our group in the past \\citep{Grindlay96, Harmon97} using the Earth occultation imaging technique \\citep{Zhang93, Harmon02}. This did not prove to be an efficient way to survey the sky due to the limitations on computer time combined with residual features and the small size (a few degrees) of the images. A new effort, incorporating more sophisticated background modeling techniques developed for the {\\it International Gamma-Ray Astrophysics Laboratory}(INTEGRAL)\\citep{Lei99} and detector point spread functions for limb geometry, holds much promise for achieving this goal \\citep{Shaw01, Shaw03}. In Fig. 1 we compare the sensitivity of the BATSE deep sample (this work) with that of the HEAO 1 survey and a proposed future low-energy gamma-ray sky survey mission the {\\it Energetic X-ray Imaging Survey Telescope} (EXIST) \\citep{Grindlay01, Grindlay03}. The binning of the BATSE sensitivity curve reflects the maximum energy resolution obtainable with the EOT.  In addition to the list of monitored sources accumulated during the operation of BATSE, we reviewed the HEAO 1 A-4 hard X-ray \\citep{Levine84} catalog and the recent gamma-ray source compilation of \\citet{MacombGehrels99} for completeness of the known high energy Galactic sources, and performed a literature search of other high-energy objects above the BATSE sensitivity threshold.  We did not include the EGRET ``unidentified\" sources from \\citet{MacombGehrels99} due to the probable extragalactic nature of many of them, and the relatively large uncertainty in location ($\\sim$1$^{\\circ}$) that would strongly affect the measured flux using the EOT.  The sources addressed in this catalog are shown on a sky map in Galactic coordinates in Fig. 2. We have included 179 sources in our deep sample of the sky. The majority of sources are accreting binaries; however, we include the brighter sources from the other object classes for comparison. Of the 179 sources, most (76\\%) are accretors, i.e., black hole or non-pulsed neutron star binaries (97), or binaries containing pulsars (41). Again we stress that our deep sampling of the known gamma-ray sources should not be considered as a true background-limited survey to a given brightness threshold. However, by virtue of our comparisons to HEAO 1 A-4 survey data, we believe that we have a reasonably complete sampling of high energy accreting systems in the Milky Way Galaxy down to an average flux level of $\\sim$30 mCrab (20-100 keV).  We also note that to date, samples of objects have been performed with BATSE for active galaxies \\citep{Connaughton99, Malizia99, Bassani00}, supernova remnants \\citep{McCollough97}, globular clusters \\citep{Ford96}, X-ray stars \\citep{White94}, and prompt emission from nearby novae \\citep{Hernanz00}.  The most promising classes of sources, other than accreting binaries, for which BATSE has shown positive detections, are active galaxies, such as Seyferts and blazars, and supernova remnants. As part of our 179-source sample, we included selected supernova remnants (17) (some of which contain pulsars or soft gamma-ray repeaters (SGRs)), active galaxies (12), and X-ray emitting stars (12), that overlap these other studies.  In Sec. 2, we give a brief overview of BATSE and the use of the Earth occultation technique.  We then discuss production details for the catalog. In Sec. 3 we present highlights of the all-sky monitoring effort, and give an overview of science obtained for the different classes of detected sources. Sec. 4 presents the deep sample results and how they are organized. We describe a sky-dependent systematic error model and apply it to the measured source fluxes.  We then examine the consistency of source detections by focusing on some of the weak sources in the few mCrab intensity range.  This is done by testing with methods that are independent of the occultation step fitting routines.  Finally, in Sec. 5, we briefly discuss applications of the catalog data. ",
        "conclusions": ""
    },
    "0404/astro-ph0404386_arXiv.txt": {
        "abstract": "We analyze in detail the peculiar velocity field traced by 56 clusters within 120 \\hmpc\\ in the ``Streaming Motions of Abell Clusters'' (SMAC) sample.  The bulk flow of the SMAC sample is $687\\pm203$ \\kms, toward \\lberr{260}{13}{0}{11}. We discuss possible systematic errors and show that no systematic effect is larger than half of the random error.  The flow does not drop off significantly with depth, which suggests that it is generated by structures on large scales.  In particular, a Great Attractor as originally proposed by Lynden-Bell et al. cannot be responsible for the SMAC bulk flow. The SMAC data suggest infall into an attractor at the location of the Shapley Concentration, but the detection is marginal (at the 90\\% confidence level).  We find that distant attractors in addition to the Shapley Concentration are required to explain the SMAC bulk flow.  A comparison with peculiar velocities predicted from the \\emph{IRAS} PSCz redshift survey shows good agreement with a best fit value of $\\beta\\sbr{I} = \\omsix/b\\sbr{I} = 0.39\\pm0.17$. However, the PSCz density field is not sufficient to acount for all of the SMAC bulk motion. We also detect, at the 98\\% confidence level, a residual bulk flow of $372\\pm127\\,\\kms$ toward \\lb{273}{6} which must be generated by sources not included in the PSCz catalogue, that is, either beyond 200 \\hmpc, in the Zone of Avoidance or in superclusters undersampled by \\emph{IRAS}. Finally, we compare the SMAC bulk flow with other recent measurements. We argue that, at depths ranging from 60 to 120 \\hmpc, flows of order 600 \\kms\\ are excluded by multiple data sets.  However, convergence to the CMB frame by a depth of 60 \\hmpc\\ is also excluded by multiple data sets.  We suggest that a bulk flow of $225\\,\\kms$ toward \\lb{300}{10} at depths greater than 60 \\hmpc\\ is consistent with all peculiar velocity surveys, when allowance is made for sparse sampling effects. ",
        "introduction": "\\label{sec:intro}  The source of the Local Group's (LG) peculiar velocity of $627\\pm22 \\kms$ toward \\lb{276}{30} \\citep{KogLinSmo93} with respect to the Cosmic Microwave Background (CMB) has been a puzzle since the detection of the CMB dipole \\citep{SmoGorMul77}.  In order to resolve this fundamental question, it is necessary to map the peculiar velocities of nearby galaxies.  If a volume which is at rest with respect to the CMB can be identified, then the masses responsible for the motion of the LG must be contained within that volume.  An important milestone in peculiar velocity surveys was the study of \\citet{LynFabBur88}, who used the \\dnsig\\ distance indicator and claimed detection of a ``Great Attractor'' (GA) at a distance of approximately 45 \\hmpc\\ believed to be responsible for most the LG's motion.  Subsequent peculiar velocity surveys of the nearby Universe undertaken in the early 90s did not reveal the expected infall signature on the far side of the GA \\citep{MatForBuc92} and hinted at a large coherence length for the flow \\citep{Wil90,Cou92}. Analyses of redshift surveys \\citep{Dre88,StrDavYah92,Hud93,Hud94b} did identify a significant excess of galaxies near the proposed GA, but with an abundance insufficient to generate the LG's motion.  For example, \\citet{Hud94b} concluded that $400\\pm45\\,\\kms$ of the LG's motion arose from sources beyond 80 \\hmpc.  The location of the peak of the GA in the Zone of Avoidance (ZOA), as given by \\citet{KolDekLah95}, suggests the possibility that most of the GA's overdensity may be obscured. Searches in the ZOA have revealed the presence of massive clusters such as A3627 \\citep{KraWouCay96} but the integrated overdensity within the ZOA appears still insufficient to generate substantial infall at the LG \\citep{StaJurHen00}.  \\citet[hereafter LP]{LauPos94} used the photometric properties of brightest cluster galaxies as standard candles and claimed that a much larger region of space extending to 150 \\hmpc\\ was moving at a velocity of $689\\pm178\\,\\kms$ toward \\lb{343}{52}.  The LP result was puzzling, first because the amplitude of the flow on such a large scale was higher than expected in popular cosmological models \\citep{StrCenOst95}, and second because the direction of the LP flow was significantly different from previous bulk flow measurements \\citep{StrWil95}.  Recent surveys of field galaxies to depths $R \\la 60\\,\\hmpc$ generally suggest that the bulk flow within these nearby volumes is in the range from $100$ -- $300\\,\\kms$ in the direction \\lbapr{300}{10} \\citep{WilCouFab97, GioHayFre98, CouWilStr00, daCBerAlo00b,TonBlaAjh00}. Clusters surveys on large ($R \\ga 100 \\hmpc$) scales have not supported the LP result, but, at face value, also appear to suggest a wide range of values for the large-scale motion \\citep{MouAkeBot93,HudSmiLuc99, Wil99b, DalGioHay99b, ColSagBur01}.  This apparent disagreement is due to neglect of sparse sampling on the quoted flow errors \\citep{WatFel95}. When the effects of sparse sampling are properly taken into account, the large-scale surveys are not in conflict \\citep{Hud03}.  For recent reviews, the reader is directed to \\citet{CouDek01} and \\citet{Zar02b}.  This paper is the fifth in a series based on the ``Streaming Motions of Abell Clusters'' (SMAC) project. The aim of this project is to obtain distance estimates, via the Fundamental Plane (FP) method, for elliptical galaxies in clusters and to map the peculiar velocity field within 120 \\hmpc\\ of the Local Group.  First results from this work were presented in \\citet{HudSmiLuc99}, who quoted a bulk flow of $630\\pm200\\,\\kms$ toward \\lb{260}{-1} from a preliminary analysis of the same sample studied here.  Previous papers in this series have reported new data for this project. \\citet[hereafter SMAC-I and SMAC-II respectively]{SmiLucHud00, SmiLucSch01} presented spectroscopic and photometric data, respectively; these data were compared and combined with other data available in the literature to obtain the final data set used for this project \\citep[hereafter SMAC-III]{HudLucSmi01}. Smith et al. (in preparation, hereafter SMAC-IV) reports FP distances for 56 clusters.  In this paper, we analyze the peculiar velocity field in the Universe as traced by the SMAC clusters.  In \\secref{data}, we summarize the SMAC peculiar velocity data. In \\secref{bulk}, we model the peculiar velocity field as a simple bulk flow (the dipole moment of the velocity field).  An important aspect of this paper is to evaluate the robustness of the results to systematic effects (\\secref{sys}).  In \\secref{othermod} we consider more complicated flow models, including toy models based on simple attractors.  \\secref{pscz} compares the peculiar velocities of SMAC clusters to the predictions from the \\emph{IRAS} PSCz density field.  \\secref{comparison} compares the flow of the SMAC sample to other results in the literature. ",
        "conclusions": "We have examined in detail the $687\\pm203$ \\kms\\ bulk flow of the SMAC sample and find that it is robust to systematic errors.  We have shown that most of the bulk flow is not generated by nearby sources such as the Great Attractor but rather arises from structures at depths greater than 100 \\hmpc.  The Shapley Concentration is identified as one likely source of the large scale flow but is unlikely to be responsible for all of the SMAC flow.  When we compare the SMAC motion to the predictions of the PSCz survey we find that $\\beta = 0.39\\pm0.17$, consistent with the ``concordance'' value $\\beta \\sim 0.5$.  However, the \\emph{IRAS} PSCz survey can only explain $50\\pm20$\\% of the amplitude of the SMAC flow. This suggests that there are gravitational sources not well-mapped by the PSCz survey.  Evidence from other redshift surveys suggests that these sources may be located in the ZOA or in superclusters such as Shapley which are undersampled by \\emph{IRAS}.  Finally, we have shown that the SMAC survey is not inconsistent with other sparsely-sampled surveys of the large scale velocity field, such as \\citet[EFAR]{DalGioHay99b,Wil99b} and the SNe sample of \\citet{TonSchBar03}. Taken together, all surveys suggest a large scale flow approximately 225 \\kms\\ towards \\lb{300}{10}.  Further analyses of the effects of sparse sampling, and detailed comparisons with the predictions of redshift surveys such as PSCz are needed to more accuractely quantify this result.  The NOAO Fundamental Plane Survey \\citep{SmiHudNel04} will measure the peculiar velocities of 100 X-ray selected clusters.  The total number of FP distances will be $\\sim 4000$, about six times the number in SMAC.  When results from the NFPS are compared to, for example, predictions of the 2MASS redshift survey \\citep{HucJarSkr03}, we expect to identify the sources the LG's motion."
    },
    "0404/astro-ph0404103_arXiv.txt": {
        "abstract": "We present {\\sc budda} (BUlge/Disk Decomposition Analysis), a new code devoted to perform a two-dimensional bulge/disk decomposition directly from the images of galaxies. The bulge component is fitted with a generalized S\\'ersic profile whereas disks have an exponential profile. No other components are included. Bars and other sub-structures, like lenses, rings, inner bars and inner disks, are studied with the residual images obtained through the subtraction of bulges and disks from the original images. This means that a detailed structural analysis of galaxies may be performed with a small number of parameters, and sub-structures may be directly studied with no a priori assumptions. As has been already shown by several studies, two-dimensional fitting is much more reliable than one-dimensional profile fitting. Moreover, our code has been thoroughly tested with artificial data and we demonstrate it to be an accurate tool for determining structural parameters of galaxies. We also show that our code is useful in various kinds of studies, including galaxies of, e.g., different morphological types, and inclinations, which also may be observed at different spatial resolutions. Thus, the code has a broader range of potential applications than most of the previous codes, which are developed to tackle specific problems. To illustrate its usefulness, we present the results obtained with a sample of 51 mostly early--type galaxies (but covering the whole Hubble sequence). These results show some of the applications in which the code may be used: the determination of parameters for Fundamental Plane and structural studies, quantitative morphological classification of galaxies and identification and study of hidden sub-structures. We have determined the structural parameters of the galaxies in our sample and found many examples of hidden inner disks in ellipticals, secondary bars, nuclear rings and dust lanes in lenticulars and spirals, and also wrong morphological classification cases. We now make {\\sc budda} generally available to the astronomical community. ",
        "introduction": "It is widely assumed that the light distribution in galaxies may be described by one or more components, which might be dynamically and/or chemically distinct. While most of ellipticals (E) may be photometrically described by only one major spheroidal component, lenticular (S0) and spiral (S) galaxies in general need at least two components: a spheroidal bulge and a flattened disk. However, in almost every object we can trace the presence of other relevant components, such as primary bars, secondary bars (in spirals and S0s), inner disks, lenses and rings, none of which have been as carefully studied as the bulges and disks. Although there is an unclear and poorly understood interplay between these components, which is not generally taken into account in most structural analyses, bulge/disk decomposition has been a major motif behind a number of important contributions to our understanding on galaxy formation and evolution. Examples include the Tully-Fisher relation \\citep{tul77}, the Fundamental Plane \\citep[hereafter FP]{fab87,dre87,djo87,ben92}, the correlation between the scale lengths of bulges and disks \\citep{cou96,dej96b,anj04} and the evolution of galaxy morphologies \\citep{mar98}. See also \\citet{pen02} and \\citet{gra02a,gra02b} for further and more recent examples.  Two important issues are related to the subject of bulge/disk decomposition: the fitting method and the mathematical functions used to describe each component. Generally speaking, there are two fitting methods, the one-dimensional (1D) and the two-dimensional (2D). In the 1D method, an azimuthally averaged surface brightness profile of the galaxy under study, or the major or minor axis profiles, is fitted by one or more components. This method has the advantage of being simple and fast, and works in low S/N conditions. However, it has two major problems. Firstly, the procedure to obtain the galaxy profile may vary, and each procedure has its own drawbacks. In isophote fitting \\citep{jed87}, which is commonly used, changes in ellipticity and position angle at consecutive isophotes turn the profile poorly defined, since it is extracted not along a straight line but an arc. This is especially true when a bar or a strong spiral arm is present. Alternatively, one may extract the profile along the galaxy major axis, which may be itself poorly defined, but in this case much of the information in the image is thrown away \\citep{bag98}. Moreover, minor axis profiles are also essential in this case since the bulge dominates the profile along this axis. Secondly, 1D fitting averages out in a complicate way non-axisymmetric components, like bars. It is possible, though, to obtain reasonable results with 1D fitting, as shown by \\citet{ken89}, which study the SB0 galaxy NGC 936 by means of 1D decomposition but also incorporating further intricate steps \\citep[see also][]{ken86}. In 2D fitting, information in the whole image is used to build a model for each component and thus differences in the projected ellipticities and position angles become an advantage, since they help to constrain the parameters of the different components. Moreover, non-axisymmetric components are now taken into account. There are several examples in the literature of studies showing that the 2D method is much more reliable than the 1D \\citep{dej96a}, retrieving more accurate structural parameters.  While there are several examples in the literature of studies which use 1D fitting \\citep{bag98}, only in the last decade 2D fitting algorithms become more popular, mainly because of the higher computational speed of today's machines. As examples, we may point out the algorithms developed by \\citet{dej96a}, by \\citet{sim98}, and more recently by \\citet{pen02} and \\citet{tru01}. Most of the available algorithms were developed to tackle specific problems, like, e.g., face--on late--type spirals \\citep{dej96a}, distant galaxies \\citep{sim98}, or nearby galaxies observed with very high spatial resolution \\citep{pen02}. On the other hand, it is a common practice to include several components other than the spheroid and the disk in an attempt to get better fits. In this way, one can model bars and, in the case of \\citet{pen02}, also nuclear point sources, inner disks and inner bars, and lenses. In the Peng et al. algorithm 40 free parameters are used, and in several cases a single component is fitted by various different functions.  A caveat behind this philosophy of performing structural analysis is related to the fact that we do not know how well a single general function could fit these sub-structures with a given accuracy. Therefore these fits might be subject to an unknown systematic effect. For instance, as some bars show an exponential profile, especially in late--type galaxies, and others have a flat light distribution, especially in early--type galaxies \\citep{elm96}, with a steep outer drop, it is difficult to foresee how a single general function could be used to fit bars and give meaningful structural parameters. This is also the case for lenses and other sub-structures. In our code, we explore the alternative route of fitting the two major components, representing the disk and the bulge contributions, and to perform a detailed study of sub-structures in the residual images obtained after the subtraction of these major components. In principle, this procedure could allow us to study the minor structures and assess the possibility to describe them as families of functions. It should be noted, however, that residual images can also be obtained with 1D fits if one can derive reasonable estimates of the apparent axis ratios of the bulge and disk components. For instance, \\citet{but92} show a residual image of NGC 4622 based on 1D decomposition expanded to 2D using assumptions about its components projected shapes. Also, 2D bulge/disk decomposition is not the only way to bring out hidden features. Using Fourier decomposition \\citep[see, e.g.,][]{con88}, \\citet{but92} were able to bring out remarkably well the amazing counter-winding spiral arms in NGC 4622. To a lesser extent, a similar analysis can be performed via unsharp masking and color maps \\citep[see, e.g.,][for a recent application]{erw02}.  Therefore, following these guidelines, we have developed a new and general purpose algorithm, {\\sc budda}\\footnote{See \\anchor{ http://www.astro.iag.usp.br/~dimitri/budda.html}{\\url{ http://www.astro.iag.usp.br/$\\sim$dimitri/budda.html}} where {\\sc budda} is available to the astronomical community.} -- which stands for BUlge/Disk Decomposition Analysis -- devoted to perform a detailed 2D structural analysis in galaxy images. The code allows us to perform a direct study on sub-structures in galaxies by means of accurate residual images, besides determining structural parameters of the main components, namely the bulge and the disk. In Sect. 2 we present the details of how the algorithm works. In Sect. 3 we present a thorough analysis of the reliability of the algorithm using a set of artificial data images of galaxies having known structural parameters. In Sect. 4 we apply our algorithm to a sample of 51 galaxies and show the results. Finally, in Sect. 5 we finish with a summary and a brief discussion of our results and the potential applications of our code. ",
        "conclusions": "We have thoroughly presented {\\sc budda}, a code devoted for bulge/disk decomposition in galaxy images. As was shown in details above, its development and performance tests show that it is a very suitable tool for the structural analysis of galaxies. The parameters for the bulge and disk models can be retrieved with great confidence as well as the residual images obtained from them. Its simplicity to use and its broad range of applications add to its qualities.  Applying {\\sc budda} on our sample of 51 galaxies we could check how it works concerning its many potential applications. One obvious use of the code is to obtain bulge and disk structural parameters to study e.g. the Fundamental Plane and other scaling relations. Such a study regarding the galaxies in our sample will appear in a forthcoming paper. Such relations may also be strengthened if wrong morphological classification cases are identified, which is reliably done with our code, as we showed in Sect. 4.3. Galaxies classified as Es, but actually being S0s, which was the case for 6 of our galaxies, may add to the spread in scaling relations. Obviously, the opposite is also true. In the FP, for instance, S0s have generally lower mass-to-light ratios than Es \\citep{ben92}.  Another use of {\\sc budda}, in which we've found it to be very powerful, is to reveal hidden sub-structures, especially in the inner regions of galaxies. In Sect(s). 4.1 and 4.3 we revealed inner disks in many elliptical galaxies, which may be the result of a recent merger. Or are they related to galaxy formation processes? {\\sc budda} has found also several examples of inner spirals, inner and nuclear rings, as well as inner and secondary bars in S0s and Ss. In some cases a hint of them shows up in the direct images or in the profiles, but in others there is no a priori sign of them at all. These sub-structures, which appear conspicuously in the residual images, are relevant for e.g. dynamical studies, since they are related to non-circular orbits and Lindblad resonances. Moreover, these sub-structures may possibly be relevant to the long standing problem of fueling active galactic nuclei \\citep{com01}, and may be directly studied in the residual images."
    },
    "0404/astro-ph0404429_arXiv.txt": {
        "abstract": "The most accurate way to measure the energy levels for the O~{\\sc ii} $2p^3$ ground configuration has been from the forbidden lines in planetary nebulae.  We present an analysis of modern planetary nebula data that nicely constrain the splitting within the $^2$D term and the separation of this term from the ground $^4$S$_{3/2}$ level. We extend this method to H~{\\sc ii} regions using high-resolution spectroscopy of the Orion nebula, covering all six visible transitions within the ground configuration.  These data confirm the splitting of the $^2$D term while additionally constraining the splitting of the $^2$P term.  The energies of the $^2$P and $^2$D terms relative to the ground ($^4$S) term are constrained by requiring that all six lines give the same radial velocity, consistent with independent limits placed on the motion of the O$^+$ gas and the planetary nebula data. ",
        "introduction": "} In this paper we determine the four energies which describe the separation of the five ground configuration energy levels of O~{\\sc ii} shown in Figure~\\ref{figene}, through analysis of observations of the visible forbidden transitions.  Being forbidden transitions, these are weak and not accessible in the laboratory.  The ground configuration levels can be determined from laboratory data on permitted ultraviolet transitions from higher states, but in practice the uncertainties in the derived energy levels are larger than those obtained from the nebular visible transitions.  Thus the published compilations of data on these energy levels (Eriksson 1988; Martin, Kaufman \\& Musgrove 1993) rely heavily on the pioneering work of \\citet{bow60} and de Robertis, Osterbrock, \\& McKee (1985) on planetary nebulae (see Table~\\ref{tbl1}).  Our analysis of new data improves upon these energies and provides a discussion of the uncertainties.  \\citet{bow60} observed the two blue [O~{\\sc ii}] lines in seven bright planetary nebulae (a total of sixteen photographic plates) and the four red [O~{\\sc ii}] lines (blended into two lines) in one nebula (three separate plates).  The O$^+$ gas of the planetary nebulae was assumed to have an average bulk velocity identical to the measured velocity of the H$^+$ gas (the H$^+$ energy levels being well known). The observed [O~{\\sc ii}] wavelengths adjusted to this velocity give absolute wavenumbers for the energy level differences for the transitions.  These early data resolved the $^2$D term into its two $J$ components, but were not able to determine the splitting of the $^2$P term. The latter splitting was determined in \\citet{der85} using high-resolution digital spectra of the four red lines in one bright planetary nebula (NGC 7027).  As can be appreciated from Figure~\\ref{figene}, the red transitions ending on a common lower level constrain the splitting of the $^2$P term while the transitions from a common upper level constrain the splitting of the $^2$D term, independently of the blue data.  \\cite{der85} found a value for the latter splitting somewhat different than that found by Bowen. Since absolute velocities or wavenumbers were not measured, \\cite{der85} calculated the $^2$D--$^2$P separation by assuming that the $\\lambda$7319/$\\lambda$7320 line blend measured by \\citet{bow60} was actually a measure of the stronger $\\lambda$7320 line, and that \\citet{bow60}'s roughly symmetric $\\lambda$7330/$\\lambda$7331 blend was a measure of this line pair's average.  We show that H~{\\sc ii} regions can also be used to determine the energy levels of O~{\\sc ii}.  With high-resolution echelle spectroscopy of the Orion nebula (resolving all red and blue lines), we are able to measure not only the splitting of both the $^2$P and $^2$D terms but also the separation between these two split terms. The observations are described in \\S~\\ref{obs} and the full analysis, including uncertainties, in \\S~\\ref{con}.  But first we update Bowen's work by analysis of more recent digital data on 23 planetary nebulae (\\S~\\ref{bow}).  This provides important insight into the method of analysis of the full Orion data and also provides improved values of the two energies constrained by the blue lines \\`a la Bowen.  We conclude by comparing our results with recent work by \\citet{sha04}. ",
        "conclusions": "\\label{conc}  From a detailed analysis of Orion nebula and planetary nebulae data, we find new energies for the ground configuration. Both model-fitting uncertainties and systematic uncertainties are presented. This work confirms the utility of astrophysical measurements in determining accurate energies for the O~{\\sc ii} $2p^3$ ground configuration. As a by-product, we determine a revised set of air wavelengths of the [O~{\\sc ii}] visible lines (see Table~\\ref{tblvel}) and the UV lines (see \\S~\\ref{con}). As \\citet{sha04} point out, there are revisions required in the standard NIST values.  Using these revised wavelengths, it is now possible to constrain the velocities of all three oxygen zones (O$^0$,O$^+$,O$^{++}$), and in turn, the spatial and physical origin of oxygen permitted lines can be inferred \\citep{bla03b}."
    },
    "0404/astro-ph0404335_arXiv.txt": {
        "abstract": "In an effort to elevate to higher grounds our understanding on the impact of the formation and evolution of bars in the formation and evolution of galaxies, we have developed a diagnostic tool to distinguish between recently formed and evolved bars. Our method was applied in the study of a sample of 14 galaxies and revealed that, apparently, AGN activity tends to appear in galaxies which have young bars rather than evolved bars. This suggests that the time scale for the fueling of AGN by bars is short, and may help to explain, for instance, why there is not a clear correlation between the presence of bars and AGN in galaxies. ",
        "introduction": "It has long been known that once a bar develops in a galaxy it induces several evolutionary processes which severely modify many aspects of its host. Changes include the stellar and gas dynamics, by the onset of new relevant orbits, resonances and torques \\cite[(see, e.g., Sellwood \\& Wilkinson 1993; Combes 2001)]{Sellwood93,Combes01}; the evolution of the stellar populations, by the attenuation of chemical composition gradients \\cite[(Martin \\& Roy 1994; Zaritsky, Kennicutt \\& Huchra 1994)]{Martin94,Zaritsky94} and the triggering of nuclear starbursts \\cite[(e.g., Gadotti \\& dos Anjos 2001)]{Gadotti01}; and the global galactic structure, with the formation of rings and lenses, and the building of bulges \\cite[(see Athanassoula 1983; Buta \\& Combes 1996; Carollo, this volume)]{Athanassoula83,Buta96}. Bars also are very likely related to the problem of fueling AGN (see, e.g., Combes, this volume, and references therein) transferring outwards the angular momentum of gas and stars in the disk. However, these processes may have different time scales and, what is particularly relevant for the subject of this meeting, this may explain, at least partially, the absence of a clear correlation between the presence of bars and AGN in galaxies. For instance, a barred galaxy may not have an AGN simply because its bar did not yet have the time to funnel the disk material to the center. Inversely, if the time scale for AGN fueling by bars is short, then we may be observing barred galaxies which had already their AGN phases. The availability of AGN fuel, the presence of inner Lindblad resonances and nuclear spirals, the possibility of destruction and rejuvenation of bars, and the recurrence of AGN phases, add complexity to the theme. Indeed, it is really naive to expect a clear correlation between bars and AGN given all the details involved.  While the time scale issue is recognized now as a very important point to better understand the observations (which means that it is also a key ingredient in theoretical modeling), very little has been done on this regard, considering bars. In spite of its relevance to many aspects of the formation and evolution of galaxies, we still lack the means to evaluate the age of a bar, i.e., for how long its presence is altering the evolutionary course of its host. Thus, we have developed a diagnostic tool to distinguish between recently formed and evolved bars, and observationally estimate the age of a bar in a galaxy. This allows us, for instance, to directly tackle the issue of the time scale in the AGN fueling by bars. ",
        "conclusions": "From the 14 galaxies of our sample we found 8 evolved bars and 5 young ones (1 case is not conclusive). Only 2/8 (25\\%) galaxies with evolved bars have AGN, while 3/5 (60\\%) galaxies with young bars have AGN. Then we have to conclude that AGN appear more often in galaxies with young rather than evolved bars, which in turn suggests that the time scale for AGN feeding by bars is short. This means that a clearer correlation between the presence of bars and AGN in galaxies may be observed when using a sample of barred galaxies whose bars are recently formed. This result is obviously based on a small number statistics and should be confirmed by further investigation. Note that to determine if a galaxy has a young or evolved bar is a task which is expensive in terms of telescope time.  We now may ask how young are recently formed bars or what is the age difference between young and evolved bars. A more difficult task, which needs deep theoretical analysis, is to obtain the age of the bar from the kinematical measurements. To investigate these questions, we have performed several $N$-body simulations in which bars form and evolve for 2 Gyr. These experiments consider only gravity physics and stars, and show that boxy/peanut bulges appear in time scales of less than about 1 Gyr, and even after 2 Gyr the values for $\\sigma_z$ in the bar are still lower than about 50 Km/s. These results lead us to conclude that, to explain the values we measured for $\\sigma_z$, another mechanism, rather than the ones which originate boxy/peanut bulges, namely vertical resonances and/or the hose instability, is playing a relevant role to give bars an expressive vertical structure. We suggest that the \\cite[Spitzer \\& Schwarzschild (1951,1953)]{Spitzer51,Spitzer53} mechanism may be the answer if we assume that giant molecular clouds are twice as much concentrated along the bar as in the disk. We have analytically estimated that in this case bars reach the observed $\\sigma_z$ in time scales of $5 - 10$ Gyr. These questions were also tackled by multiband (BVRIKs) imaging, which allowed us to measure the colors of the bars. Interestingly, the mean difference in the colors of young and evolved bars suggests a difference in age of about the same time scale. These results will appear in a forthcoming paper, but see \\cite{Gadotti03a}.  \\begin{figure} \\centering \\includegraphics[width=6.55cm]{dgadotti-fig2.eps}\\includegraphics[width=6.45cm]{dgadotti-fig3.eps} \\caption{Left: frequency of strongly barred galaxies (B's), weakly barred galaxies (X's) and all barred galaxies (B's + X's) along the Hubble sequence of all galaxies brighter than 14 B mag and with a bar classification in the RC3 (\\cite{deVaucouleurs91}). Right: frequency of all Seyferts (1 and 2) and LINERs along the Hubble sequence in the catalog of \\cite{Lipovetsky88}.} \\end{figure}  In this regard, it is interesting that we have found that weak bars have bluer colors than strong bars. This suggests that the former are in the process of formation rather than in dissolution. To further analyse this issue, we have studied the frequency of bars along the Hubble sequence, and found an excess of weakly barred galaxies in morphological classes T=5$\\pm$1, which corresponds to Sbc, Sc and Scd. Moreover, we have also found that, while Seyferts are more common in early-type spirals (T=1), LINERs are preferentially found in late-type ones (T=5, see Figure 2). If we distinguish, among galaxies with AGN, those which are interacting (i.e., either in pairs, groups or clusters) and those which have bars, like in Table 1, we can see that early-type galaxies are mostly interacting, and that the fraction of barred galaxies rises when we consider late-type galaxies. Taken altogether, these results seem to indicate that young bars appear more often and are more important to feed AGN (especially LINERs) in late-type spirals, rather than in early-type ones, where interactions seem to play a more relevant role.  \\begin{table} \\begin{center} \\begin{tabular}{ccc} T & Interacting & Barred \\\\ \\\\ 0 - 1 & 68 (56\\%) & 35 (29\\%) \\\\ 2 - 3 & 51 (53\\%) & 34 (35\\%) \\\\ 4 - 5 & 25 (44\\%) & 26 (46\\%) \\\\ \\end{tabular} \\caption{Galaxies with AGN which are either interacting and/or have bars in the catalog of \\cite{Lipovetsky88}. The numbers in parentheses are the fraction of galaxies satisfying these conditions in each morphological bin.} \\end{center} \\end{table}"
    },
    "0404/astro-ph0404045_arXiv.txt": {
        "abstract": "The spherically symmetric accretion of dark and phantom energy onto Morris-Thorne wormholes is considered. It is obtained that the accretion of phantom energy leads to a gradual increase of the wormhole throat radius which eventually overtakes the super-accelerated expansion of the universe and becomes infinite at a time in the future before the occurrence of the big rip singularity. After that time, as it continues accreting phantom energy, the wormhole becomes an Einstein-Rosen bridge whose corresponding mass decreases rapidly and vanishes at the big rip. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404273_arXiv.txt": {
        "abstract": "Most of the soft and a growing fraction of the harder X-ray background has been resolved into emission from point sources, yet the resolved fraction above $7\\kev$ has only been poorly constrained. We use $\\sim700\\ks$ of \\xmm\\ observations of the Lockman Hole and a photometric approach to estimate the total flux attributable to resolved sources in a number of different energy bands. We find the resolved fraction of the X-ray background to be $\\sim90$ per cent below $2\\kev$ but it decreases rapidly at higher energies with the resolved fraction above $\\sim7\\kev$ being only $\\sim50$ per cent. The integrated X-ray spectrum from detected sources has a slope of $\\Gamma\\sim1.75$, much softer than the $\\Gamma=1.4$ of the total background spectrum. The unresolved background component has the spectral signature of highly obscured AGN. ",
        "introduction": "The cosmic X-ray background (XRB) extends from $\\sim0.1$ to several hundred $\\kev$. At the lowest energies ($\\lesssim0.25\\kev$) the dominant contribution is the thermal emission from hot gas in the local bubble and the Galaxy. Above $1\\kev$ it is mainly extragalactic in origin and is well-described by a power-law with a spectral slope of $\\Gamma=1.4$. The spectrum peaks in $\\nu I_{\\nu}$ at $\\sim30\\kev$ before decreasing exponentially at higher energies.  It is now clear that the extragalactic XRB is due to the integrated emission of point sources, principally Active Galactic Nuclei (AGN), along with some thermal emission from galaxy clusters and starbursts (that contribute at soft X-ray energies). Since the early years of the \\einstein\\ observatory \\citep{maccacaro91}, a larger and larger fraction of the XRB has been resolved into discrete objects. Most ($\\sim70-80$ per cent) of the $0.5-2\\kev$ background was resolved by \\rosat\\ \\citep{hasinger98}. Using \\asca, \\citet{ueda98} could account for $\\sim30$ per cent of the $2-10\\kev$ background whilst an early \\chandra\\ survey of the SSA13 field resolved some $\\sim60-80$ per cent \\citep{mushotzky00}.  The \\chandra\\ Deep Fields North and South have now been able to resolve some $\\sim70-90$ per cent of the hard XRB \\citep[CDF-N;][]{hornschemeier00,hornschemeier01,brandt01_2,brandt01_1,miyaji02,alexander03}, \\citep[CDF-S;][]{campana01,tozzi01,giacconi01,giacconi02,rosati02,moretti02}. Significant uncertainty in this fraction results from field-to-field variations.  With the first deep X-ray survey conducted by the \\xmm\\ observatory \\citet{hasinger01} resolved $\\sim60$ percent of the $5-10\\kev$ hard XRB in the Lockman Hole. This field is one of the most well studied regions of the sky and has been the target of many different multi-wavelength observations. It has an unusually low level of neutral Galactic absorption, only $\\sim5\\times10^{19}\\pcmsq$ compared to the typical $2-3\\times10^{20}\\pcmsq$ at high latitudes \\citep{lockman86,elvis94}. This makes it an excellent window through which to observe distant sources at X-ray energies.  The $\\sim100\\ks$ of \\xmm\\ data used in the \\citet{hasinger01} work has now been extended to $\\sim700\\ks$. The large accumulated observation time; the low level of Galactic absorption; and, most importantly, the large effective area of \\xmm\\ at energies above $7\\kev$ make the data a unique resource. Whilst other studies will focus on source identification and spectra, the aim here is to produce a robust result for the resolved fraction of the XRB, as a function of energy, by summing broad-band fluxes for the detected sources.  Recent synthesis models \\citep[e.g.][]{gilli01,gandhi03,ueda03} suggest there may be an unresolved population of highly obscured, faint AGN. Such objects would have negligible flux in the softer energy bands with the population only presenting itself as a missing fraction in the resolved flux at higher energies. \\citet{moretti03} compile $\\rm{log}\\thinspace N-\\rm{log}\\thinspace S$ distributions and integrate these to determine the resolved fluxes to be $\\sim94.3^{+7.0}_{-6.7}$ and $\\sim88.8^{+7.8}_{-6.6}$ per cent in the $0.5-2\\kev$ and hard $2-10\\kev$ bands respectively. Using the recent XRB intensity measurement by \\citet{deluca03}, with a higher value of the background normalisation, reduces the hard band resolved fraction to only $\\sim78$ per cent. $100$ per cent resolution is not necessarily expected, with maybe $\\sim1$ per cent in the hard band resulting from truly diffuse emission \\citep{soltan03}. This is not enough, however, to make-up the significant shortfall in the hard band resolved fraction.  Figures based on the broad $2-10\\kev$ band can be deceptive since the sensitivity of most instruments in the $7-10\\kev$ range is poor, e.g. the effective area of \\chandra\\ is negligible above $7\\kev$; however, \\xmm\\ remains sensitive (at least in the PN camera) to $\\sim12\\kev$ and by calculating the resolved fraction separately in a number of finer bands we can analyse much more carefully the behaviour of the fraction with energy. ",
        "conclusions": "We have shown that only $\\sim50$ per cent of the hard X-ray background above $7\\kev$ is resolved and we are missing a considerable population of faint, hard AGN. This is supported by the hardening of the XRB intensity between the spectrum measured from bright sources compared to that of the fainter sources (Fig.~\\ref{intensity_split}). The missing component (Fig.~\\ref{residual}) has the shape of a heavily absorbed AGN spectrum. Given the negligible X-ray flux, the scattered soft component must be weak, consistent with the emission from almost complete ($4\\pi\\sr$) coverage by obscuring matter.  Fig.~\\ref{obscuration} shows the column density required to harden an intrinsic $\\Gamma=2$ emission spectrum into the shape of the faint, hard, unresolved population which has a power-law slope of $\\Gamma\\sim0.6$ (although note the spectrum is rather curved with extremes of $\\sim-0.2$ and $\\sim0.9$ at the ends). The column density is greater than $10^{23}\\pcmsq$ and Compton-thick objects are required at redshifts $\\gtrsim2$. Such AGN are common at low z -- two of the three nearest AGN; NGC\\thinspace4945 and the Circinus galaxy, are Compton thick \\citep[e.g.][]{matt00} -- few have been found at high redshift \\citep{stern02,norman02,crawford03,gandhi04}. Distant ones are difficult to detect individually \\citep{fabian02}. The unresolved hard XRB component we have detected has the signature of a substantial obscured AGN population.  \\begin{figure} \\rotatebox{270}{ \\resizebox{!}{\\columnwidth} {\\includegraphics{obscuration.ps}}} \\caption{The column density required to harden an unobscured $\\Gamma=2$ spectrum into that of the residual XRB component, as a function of redshift. The residual component (Fig.~\\ref{residual}) was approximated to a power-law, the upper and lower bounds plotted are the results for the extremes of the power-law slope for the component ($\\Gamma=-0.2$ and $\\Gamma=0.9$). The numbers within the plot are the approximate maximum value of the unosbcured, rest-frame, $2-10\\kev$ luminosity of the undetected sources ($\\rm{erg~s^{-1}}$) in order to remain below the sensitivity limit (for that redshift and column density).} \\label{obscuration} \\end{figure}"
    },
    "0404/astro-ph0404534_arXiv.txt": {
        "abstract": "We review recently developed models of galactic discrete sources of high energy neutrinos. Some of them are based on a simple rescaling of the TeV $\\gamma$-ray fluxes from recently detected galactic sources, such as, shell-type supernova remnants or pulsar wind nebulae. Others present detailed and originally performed modeling of processes occurring close to compact objects, i.e. neutron stars and low mass black holes, which are supposed to accelerate hadrons close to dense matter and radiation fields. Most of the models considered in this review optimistically assume that the energy content in relativistic hadrons is equal to a significant part of the maximum observable power output in specific sources, i.e. typically $\\sim 10\\%$. This may give a large overestimation of the neutrino fluxes. This is the case of models which postulate neutrino production in hadron-photon collisions already at the acceleration place, due to the likely $e^\\pm$ pair plasma domination. Models postulating neutrino production in hadron-hadron collisions avoid such problems and therefore seem to be more promising. The neutrino telescopes currently taking data have not detected any excess from discrete sources yet, although some models could already be constrained by the limits they are providing. ",
        "introduction": "\\label{sec:intro}  Galactic sources can potentially produce interesting event rates in neutrino telescopes. Since they are at shorter distances to the Earth ($\\sim 1 \\div 10$ kpc) compared to extra-galactic sources, the source luminosity required for a galactic source to produce the same event rate as an extra-galactic one, is orders of magnitude smaller (see Sec.~\\ref{sec1}). Given the large photon luminosities observed from some of the sources, it is possible to single out interesting candidate neutrino emitters in the Galaxy. A rough estimate of the source luminosity required to produce a certain event rate in a neutrino telescope (Halzen~2003) is shown in Sec.~\\ref{sec1}.  In order to produce high energy neutrino fluxes, galactic sources must accelerate particles at sufficiently high energies. Hillas derived the maximum energy E at which a particle of charge Z can be accelerated, from the simple argument that the Larmor radius of the particle should be smaller than the size of the acceleration region (Hillas 1984). If energy losses inside sources are neglected, this maximum energy E (in units of $10^{18}$ eV) is related to the strength of the magnetic field B (in units of $\\mu$Gauss) and the size of the accelerating region R (in units of kpc) by the following relationship: \\begin{equation} \\rm E_{18} \\sim \\beta~Z~B_{\\mu G}~R_{kpc} \\label{hill} \\end{equation} \\noindent where $\\beta$ is the velocity of the shock wave or the acceleration mechanism efficiency. Hence, the maximum energy up to which particles can be accelerated depends on the $B R$ product. Particle acceleration may occur in many candidate sites, with sizes ranging from kilometers to megaparsecs.  Most of the models discussed below assume that neutrinos are produced by charged pion decays which in turn are produced in hadron-hadron ($pp$) and hadron-photon (p$\\gamma$) collisions. As a matter of fact, in order to produce neutrinos with energies of the order of 10 TeV, the primary hadron energy has to be about two orders of magnitude larger in the case of $pp$ collisions, and about one order of magnitude larger in the case of p$\\gamma$ collisions. Neutrinos produced as decay products of neutrons (which originate in hadronic collisions) are usually much less important since they are produced with much lower multiplicities (with respect to charged pions from $pp$ collisions) and typically take only $\\sim 10^{-3}$ of the initial neutron energy. Therefore, the fluxes of neutrinos from neutron decay are significantly lower than those produced in hadronic collisions (through meson decays) due to the steepness of the spectrum of accelerated hadrons (described usually by a $E^{-2}$ power law). Only in some special conditions interesting fluxes of neutrinos from decay of neutrons might be expected, e.g. in the case of neutrons produced in heavy nuclei photo-desintegration which do not lose efficiently energy in pion production in the radiation field.  In this review paper, we concentrate on the models of galactic neutrino sources which have been developed over the past few years. For reviews on earlier works, the reader is referred to, e.g., Gaisser et al. (1995) and Learned \\& Mannheim (2000). In Section~\\ref{sec:sn} we consider high energy neutrino fluxes produced during the early phase of supernova explosions, whereas Section~\\ref{sec:pwn} will be devoted to pulsar wind nebulae. Shell-type supernova remnants are also good candidate sources for high energy neutrino emission, and this will be illustrated in Section \\ref{sec:shell}. In Section~\\ref{sec:pulsars} we will discuss neutrino emission from pulsars close to high density regions, such as those observed in the Galactic Center or massive stellar associations, and from neutron stars in binary systems in Section~\\ref{sec:binary}. Finally, in Sections \\ref{sec:mq} and \\ref{sec:magnetars} we will discuss models of microquasars and magnetars. Section \\ref{sec:conc} will be devoted to our conclusions.  The predicted event rates by the models reviewed in this paper should vary in the presence of neutrino oscillations. The expected flavor ratios at sources of high energy cosmic neutrino fluxes from $pp$ or $p\\gamma$ collisions are $\\phi_{\\nu_e}:\\phi_{\\nu_{\\mu}}:\\phi_{\\nu_{\\tau}}=1:2:<10^{-5}$. Given the solar neutrino results (see for instance Fogli et al.~2004), and particularly the recent results by SNO (Ahmed et al.~2004) with best fit oscillation parameters $\\Delta m^2 = 7.1 \\cdot 10^{-5}$ eV$^2$ and $\\theta = 32.5$, the strong evidence for $\\nu_{\\mu}$ and $\\nu_{\\tau}$ maximal mixing from atmospheric neutrino experiments which prefer $\\Delta m^2 \\sim 2 \\cdot 10^{-3}$ eV$^2$ (Super-Kamiokande (Fukuda et al.~1998), MACRO (Ambrosio et al.~1998) and Soudan 2 (Sanchez et al.~2003)), and the limit on the mixing matrix element $|U_{e3}|\\ll 1$ from CHOOZ (Apollonio et al.~2003), these flavor ratios are expected to transform into $1:1:1$ (Learned \\& Pakvasa~1995, Athar et al.~2000) along $\\gtrsim kpc$ distances in a direct hierarchy scenario. For models considering electron antineutrinos from neutron decay (such as Anchordoqui et al.~2003a) described in Sec.~\\ref{sec:massive}), about 60\\% of the electron antineutrinos should reach the Earth and $\\sim 40\\%$ should oscillate into $\\bar{\\nu}_{\\mu}$ and $\\bar{\\nu}_{\\tau}$ in about the same percentage.  In this paper we give an estimate of the muon neutrino event rates in the absence of neutrino oscillations. If neutrino oscillations are to be considered, a part of the muon neutrinos from mesons decay convert into other flavors and could be detected. Nevertheless, the estimate of event rates for $\\tau$ events in neutrino telescopes needs a calculation including the regeneration effect in the Earth due to the fast tau decay (see Bugaev et al. 2004 and references therein). Moreover, it should be considered that the topologies induced by $\\nu_{\\tau}$ and $\\nu_e$ are different than $\\nu_{\\mu}$ induced events and hence have different detection efficiencies. As a matter of fact, neutrino telescopes measure neutrinos indirectly thanks to their interactions with the matter inside or surrounding instrumented regions (for a review on recent results see Montaruli~2003). The new generation of neutrino telescopes aims at the detection of the Cherenkov light produced by charged secondaries induced in neutrino interactions. 3-D arrays of optical modules, pressure resistant glass spheres containing phototubes (PMTs), are deployed in ice or sea/lake water depths. For muon neutrinos (and also for extremely high energy $\\nu_{\\tau}$), the effective target mass is much larger than the instrumented region, since the technique profits of the increase of the muon (tau) range with energy. On the other hand, neutrinos of other flavors and neutral current interactions produce showers in the instrumented regions. ",
        "conclusions": "\\label{sec:conc} In general we can distinguish transient, persistent, and periodic neutrino sources. Some of these sources should produce strongly collimated neutrino beams, e.g. supernova models (Protheroe et al.~1998), binary system models (Anchordoqui et al.~2003b, Bartosik et al.~2003), or magnetar model (Zhang et al.~2003). Therefore, a specific geometry is required in order to have a chance to detect them.  The neutrino fluxes expected from transient sources such as the supernovae can produce from a hundred up to a few thousands of neutrino events in a 1 km$^2$ detector during the first year after supernova explosion within our Galaxy, provided that the accelerated hadrons take a significant part of the available supernova energy (Waxman \\& Loeb~2001, Beall \\& Bednarek~2002, Nagataki~2004).  However, a more realistic estimation of the hadron acceleration efficiency, which is based on the condition that the observed cosmic ray flux in the Galaxy should not be exceeded (Blasi et al.~2000, Beall \\& Bednarek~2002), produce significantly reduced neutrino event rates in a 1 km$^2$ detector. A few events can be expected in the case of pulsar formation with typical values of the surface magnetic fields of radio pulsars, but with initial periods of the order of a few milliseconds. If a pulsar with extremely high surface magnetic field is born in a supernova explosion, a so called magnetar, then the estimated event rates might be higher. However, in this case the pulsar can additionally accelerate the supernova envelope, dropping faster the column density of matter for relativistic hadrons and consequently lowering the expected neutrino event rates. This effect has not been taken into account in the calculations presented by Beall \\& Bednarek~(2002) and Nagataki~(2004). In general, detection of neutrinos from the early stage of supernovae in a reasonable time in not very promising due to the low rate of supernova explosion in the Galaxy $\\sim 10^{-2}$ yr$^{-1}$. Estimated event rates from Galactic supernovae do not give also much hope of detecting such supernovae in the nearby galaxies. For example, the expected event rates from the supernova SN1987A in LMC, at distance of $\\sim 5$ times larger than a typical Galactic supernova, are below or at the level of sensitivity of a 1 km$^2$ detector in most of the considered realistic models.  The microquasar and magnetar models (Levinson \\& Waxman~2001, Zhang et al.~2003) predict efficient production of neutrinos during the outbursts observed in the radio and/or X-ray bands. The predicted strong correlation between the expected neutrino fluxes and this low energy emission should eventually help in extracting the neutrino signal from the atmospheric neutrino background. This will require long multi-wavelength campaigns since the expected detection of the neutrino signal from , e.g. GRS 1915+105, will need the observation of about 30 outbursts.  More promising detection of neutrinos is expected from persistent sources such as the PWNe and the shell-type SNRs, provided that a significant part of their TeV $\\gamma$-ray emission have hadronic origin. The expected neutrino event rates from these sources are relatively low in a 1 km$^2$ detector, of the order of a few events yr$^{-1}$ (Alvarez-Muniz \\& Halzen~2002, Bednarek~2003a, Amato et al. 2003), but we know exactly where to look for them. Therefore, even a relatively weak signal from the direction of the Crab Nebula or TeV $\\gamma$-ray shell-type SNRs will give important constraints on hadronic processes. Another likely persistent neutrino source is expected in the Galactic Centre region. If excess of CRs from the Galactic Centre at $\\sim 10^{18}$ eV, found in the AGASA and SUGAR data, is confirmed (Hayashida et al.~1999, Bellido et al.~2001), then it should be accompanied by a neutrino source able to produce up to several neutrino events in a 1 km$^2$ detector per year (Bednarek~2002, Anchordoqui et al.~2003a, Biermann et al.~2004). In this case, neutrino fluxes would be correlated to the measured CR flux.  Also the models involving hadron acceleration inside massive binary systems are expected to be persistent and periodic (with the orbital period) neutrino sources. Bartosik et al. (2003) consider in more detail the old scenario in which energetic pulsar is responsible for acceleration of hadrons which interact with the massive companion matter. The model predicts that  detectable neutrino fluxes can be observed, also from non-eclipsing binary systems, due to the propagation of relativistic hadrons in the magnetic field of the massive star, which allows them to fall onto the star at large angles to the plane of the binary system. Anchordoqui et al.~(2003b) argue that the relativistic hadrons impinging on the surface of the accretion disk around the neutron star, which is companion of the massive star, can produce a neutrino signal above the atmospheric neutrino background in a 1 km$^2$ detector near the periastron for very dense accretion disks, e.g. in the case of A0535+26.  Models for astrophysical neutrino sources in our Galaxy usually, very optimistically, assume efficient transfer of the energy generated inside the source into relativistic hadrons. It looks that the acceleration efficiency is going to be constrained for some sources by the upper limits derived on the basis of AMANDA observations (Ahrens et al.~2003, Hauschildt et al.~2003). For example, the microquasar model (Distefano et al.~2002) prediction is about a factor of 2 below the AMANDA-II 90\\% c.l. upper limit for the persistent source SS433 (Ahrens et al.~2004a).  In general, models which postulate the production of neutrinos in hadron collisions with the radiation field already inside the acceleration region (e.g. the early stage of the model by Beall \\& Bednarek~(2002), the microquasar model by Levinson \\& Waxman (2001), or the magnetar model by Zhang et al.~(2003)) may have problems with efficient conversion of energy from the acceleration mechanism into relativistic hadrons. This is due to the fact that the threshold for $e^\\pm$ pair production is about two orders of magnitude below the threshold for pion production and the cross section for $e^\\pm$ production is significantly larger than that for pion production in the case of hadron-photon collisions. Therefore, pair production is more efficient than pion production by protons, with the possibility to have $e^\\pm$ pair domination of accelerated plasma. Such $e^\\pm$ pair domination of the plasma during the acceleration process in the electric field (induced in reconnection regions) has been discussed by Bednarek \\& Kirk~(1995) and in the case of shock accelerated particles by Mastichiadis \\& Kirk (1995). As a result, most of the energy from the acceleration mechanism is transferred to leptons and it is radiated in the form of electromagnetic radiation but not of neutrinos. In the case of models which postulate production of neutrinos in hadron collisions with matter the situation is more promising. The cross section for $e^\\pm$ pair production in collisions of two hadrons is significantly lower (even for nuclei as heavy as iron) than the cross section for pion production. Therefore, the domination in the acceleration region by the $e^\\pm$ plasma is not possible, provided that other processes, e.g. leptonic, do not contribute significantly. Moreover, in these models the region of hadron acceleration is usually different than the region of their interaction with matter."
    },
    "0404/astro-ph0404472_arXiv.txt": {
        "abstract": "Deep $K'$ images with 0.1 arcsec angular resolution, obtained with ALTAIR$+$NIRI on Gemini North, are used to investigate photometric variablity among the brightest asymptotic giant branch (AGB) stars in the central regions of M32. Based on a comparison with brightnesses obtained from the $K-$band data discussed by Davidge et al. (2000, ApJ, 545, L89), it is concluded that (1) at least 60\\% of bright AGB stars near the center of M32 are photometrically variable, and (2) the amplitudes of the light variations are similar to those of long period variables in the Galactic bulge. We do not find evidence for a population of large amplitude variables, like those detected by IRAS in the Galactic bulge. The technique discussed here may prove useful for conducting an initial reconnaisance of photometric variability among AGB stars in spheroids in the Virgo cluster and beyond, where the required long exposure times may restrict observations to only a few epochs. ",
        "introduction": "Studies of nearby galaxies lay the groundwork for understanding the evolution of more distant systems. While lacking a massive classical elliptical galaxy, the Local Group contains the compact elliptical galaxy M32, which served as the basis for early pioneering efforts to understand the stellar contents of distant galaxies (e.g. O'Connell 1980; Rose 1985). However, the nature of M32, and its relation to classical elliptical galaxies, remains a matter of on-going debate. While it has been suggested that interactions with M31 played a critical role in defining the present-day properties of the galaxy (e.g. Faber 1973, Burkert 1994; Bekki et al. 2001), the central structural characteristics of M32 suggest that its gross properties have not changed drastically since it formed (Graham 2002; Choi, Guhathakurta, \\& Johnston 2002).  Studies of the stellar content of M32 will provide insights into its evolution, and the brightest stars are convenient probes for this purpose. The brightest red stars in M32 are evolving on the asymptotic giant branch (AGB), and a number of studies have investigated the resolved stellar content in the outer (Freedman 1989, 1992; Elston \\& Silva 1992; Davidge \\& Jones 1992, Grillmair et al. 1996, Davidge 2000) and inner (Davidge et al. 2000) regions of this galaxy. The number density, normalized with respect to red surface brightness, and peak brightness of the brightest AGB stars do not change with radius (Davidge 2000; Davidge et al. 2000), indicating that these stars are well mixed throughout the main body of M32, and likely belong to a population that formed when its present-day morphology was imprinted. The bright AGB content of M32 also appears to be representative of the inner bulge of M31 and more distant systems (Davidge 2001, 2002).  Many of the brightest AGB stars are expected to be long period variables (LPVs), and studies of these objects will provide clues into the stellar content of M32. For example, the detection of very large amplitude LPVs, such as those that were detected by IRAS in the Galactic bulge (e.g. Glass et al. 1995, 2001), would be one signature of an intermediate age population. In the present paper, high angular resolution $K-$band observations of the brightest AGB stars in the central regions of M32 are used to investigate photometric variablity among these objects. A comparison is made with the LPV content of the Galactic bulge. ",
        "conclusions": "We have measured the change in $K-$band magnitude of bright AGB stars near the center of M32 over a 3 year baseline. The distribution of brightness differences, which is a measure of the amplitude of light variations, has been compared with that expected for LPVs detected at visible-red wavelengths in the Galactic bulge. It is concluded that at least 60\\% of the bright AGB stars near the center of M32 are LPVs with amplitudes similar to LPVs in the Galactic bulge.  Rejkuba et al. (2003) discuss the properties of LPVs in the nearby elliptical galaxy NGC 5128 (Cen A). They find that (1) the majority of LPVs are like those in the Galactic bulge, and (2) in one of the fields they studied at least 70\\% of the AGB stars are LPVs. These results are very similar to what we find in M32. However, Rejkuba et al. (2003) also found that 10\\% of the LPVs had periods that were comparable to the IRAS variables detected in the Galactic bulge, which presumably come from a population that is younger than the shorter period LPVs. Very long period, large amplitude LPVs were intentionally excluded from the comparison in \\S 3. The distribution of $\\Delta K$ values computed from the IRAS variables studied by Glass et al. (1995) is very different from that computed from the variables discovered at visible-red wavelengths, and the standard deviation in the $\\Delta K$ distribution computed solely from IRAS variables is $\\sigma = \\pm 1.0$ mag, which is almost three times that measured for LPVs detected at visible-red wavelengths. There are some bright objects in the $(K, \\Delta K)$ diagram that have $\\Delta K = \\pm 1$, but these are probably the extreme members of a lower amplitude LPV distribution. There is a selection effect working against the detection of large amplitude variables in the current study, as they may fall below the detection limit when at the faint points in their light curves; however, these variables should be easy to detect when near the peak of their light curves. While we can not discount the presence of a modest number of large amplitude variables in M32 with these data, it is worth noting that these objects tend to have very red colors, and the narrow, well-defined AGB sequence see near the center of M32 (Davidge et al. 2000) and at moderately large radii (Davidge 2000) argues against the presence of such a population.  Surveys of variable stars in Baade's Window (BW) reveal objects with a range of characteristics that match those seen in the Magellanic Clouds (Glass \\& Schultheis 2003). Some of the variables in BW have modest amplitudes, and such a population is likely to be present in M32; indeed, some fraction of the `non-variable' component included in the models may actually be small amplitude variables. However, the number of small amplitude variables in our sample is likely modest. Assuming that the distance modulus of M32 is $\\mu_0 = 24.4$ (van den Bergh 2000), and that the distance modulus of the Sgr I field is $\\mu_0 = 14.7$ (Glass et al. 1995), then the range of brightnesses used to construct the M32 $\\Delta K$ distribution in Figure 2 corresponds to K$_0$ between 6.6 and 7.1 in BW. It is evident from Figure 7 of Glass \\& Schultheis (2003) that the variable stars in BW in this brightness interval tend to have log(P) $> 2$, and hence are larger amplitude variables. Lower amplitude variables occur in large numbers in BW only when $K_0 > 7.1$, and it can be anticipated that these objects will become more significant in M32 at fainter brightnesses than considered here.  We close by noting that with the up-coming generation of AO-equipped $30+$ metre telescopes it will be possible to resolve bright stars in spheroids in the Virgo cluster and beyond (e.g. Anthony et al. 2003). A traditional characterization of LPVs in galaxies at this distance, involving observations over a large number of nights to obtain light curves, may only be practical for a modest number of systems because of the long integration times involved and time assignment pressures imposed by the need for nights with superb imaging conditions. The technique employed here, which requires observations covering only two epochs, provides an economical means of surveying the amplitude characteristics of LPVs in a large number of systems.  \\parindent=0.0cm  \\clearpage"
    },
    "0404/astro-ph0404158_arXiv.txt": {
        "abstract": "{ An analysis of the radial distribution of  ab-type RR Lyrae star periods in the outer region of the globular  cluster M3 at r $\\ge$ 0.83$^\\prime$ has been performed. That analysis points towards a real shortage of stars with long periods  in the radial distance range  3.5$^\\prime$ to  6$^\\prime$ (or about 7 to 12 core radii). A brief discussion is presented. The origin of the phenomenon remains an open question. ",
        "introduction": "Accurate distances to resolved stars in the present-day universe are of  basic importance for the  observational study of the stellar content and formation of individual galaxies over time (e.g. Dolphin et al. 2002; Mackey \\& Gilmore 2003; Gratton et al. 2003). By providing a means to evaluate the zero point of the distance scale, RR Lyrae stars, in addition to classical Cepheid stars,  provide a powerful constraint on the physics of the formation of the Galaxy and Local Group of galaxies. For example, for reliable information on stellar masses and their epochs of formation, accurate distances are crucial (e.g. Dolphin et al. 2002). There is however a great need to remove systematic errors both in empirical and theoretical studies of RR Lyrae stars (e.g. see Carretta et al. 2000; Caputo et al. 2000 and references therein). As their periods  are  more accurately known than other parameters (e.g. magnitudes, colours  and amplitudes), the distributions of RR Lyrae star periods in different systems should provide a cleaner test of the theoretical models of RR Lyrae stars. The purpose of the present paper is to use the database of RR Lyrae star periods for M3 from Corwin \\& Carney (2001; referred to hereafter as CC01) and to explore the distribution of periods in the outer region of that star cluster.  There have been  past studies of the periods of RR Lyrae stars on different spatial scales. For example, on Galactic scales the mean periods of such stars vary systematically with Galactic latitude (e.g. Fig.9a in Cseresnjes 2001). Differences in period distributions (e.g., mean period and width) are also known to occur in Galactic sub-systems and satallite galaxies. For example, while period distributions are quite similar in some systems (e.g., Sgr and the Large Magellanic Cloud) they are significantly different in others: e.g., compare  Sgr with Carina, Leo II or M15 for instance (Cseresnjes 2001). Do  differences occur within any galactic globular clusters ? This will be examined in the present paper using the  radial distribution of RR Lyrae stars in the outer region of M3, a star cluster  with one of the largest RR Lyrae star populations in the   Galaxy.  The layout of this short paper is as follows. In Sec.~\\ref{sample1}, the  period data and its radial distribution is presented and analyzed. That is followed by a summary of the  result and a basic discussion  in Sec.~\\ref{discuss}. ",
        "conclusions": "\\label{discuss} In the previous section it was shown  clearly that there is evidence for a possible shortage of long period stars of the type ab in the outer region of M3. It  relies on the high precision and accuracy of periods even in the innermost regions, unlike  mean colours and amplitudes that may be affected by systematic photometric uncertainties due to crowding. It is stated here as only a possible shortage because although random errors are assessed well  by the boot-strap technique applied earlier, it is not possible to rule out the possibility of a  systematic error  arising from the finite size of the RR Lyrae star population. With this a caveat, it  is  concluded that the general trend for the dispersion of period values is a drop from  the  inner region sampled (0.83$^{\\prime}$ to 3.5$^{\\prime}$)  to smaller values in an intermediate radial region (3.5$^{\\prime}$ to  6 $^{\\prime}$), and an increase again toward the outer region sampled (beyond about 6$^{\\prime}$). Beyond a radial distance of about 9$^{\\prime}$, there is little information due to a decreasing surface density of stars.   Due to the potential importance of the result in the context of single star evolution  in star clusters, an important parameter worth speculating about briefly is chemical composition. Both observationally (e.g. Sandage 1982) and theoretically (e.g. Bono et al. 2001) it is known that there is a correlation between period, heavy element abundance and K-band magnitude, i.e. the PKZ relation\\footnote{The  apparent shortage of long periods at r = 3.5-6$^{\\prime}$ seems to be reflected in the  B- and V-band  amplitude data from CC01. It is  expected because of a known correlation between optical amplitudes and periods in M3 (e.g., Roberts \\& Sandage 1955, Fig. 6; CC01, Fig. 9). Herein  an amplitude analysis  is excluded  because of likely systematic effects caused by light curve scatter and/or stellar crowding.} Thus, {\\it if} taken at face value, the apparent shortage of long periods should be  reflected in K-band magnitudes and/or Z. For example, applying the tight  PKZ relation  (e.g. Eqn. (3) of Bono et al. (2001)),   one finds that for M$_{\\rm K}$ = -0.3\\,mag at a mid-way period of 0.5\\,d, a period increase of 0.05\\,d would correspond to $\\Delta$log\\,Z $\\sim$ 0.6\\,dex. The current census of [Fe/H] values is however too incomplete to warrant an examination of its radial distribution at the present time.  Another issue worth mentioning is the spatial distribution of RR Lyrae stars in M3. A Kolmogorov-Smirnov test of the cumulative radial distribution of the RR Lyrae stars and  red giant branch (RGB)  stars, using RGB data from Ferraro et al. (1997; and references therein), indicates a probability of 82\\% that they are from the same parent populations (proof of the null-hypothesis). That suggests that the dynamical histories of the RR Lyrae star and  RGB star populations in M3 are  similar in a statistical sense over the range of radial distances examined in the present paper. Consequently, if there is some form of selection mechanism that affects RR Lyrae stars, and if that mechanism is linked to the dynamics of cluster evolution, its effect on RGB stars, if any, may vary with radial distance in the cluster.     Lastly, although there is no supporting empirical evidence, it is of academic interest to sketch very briefly one intriguing potential selection mechanism; that is the issue of gas giant planets or brown dwarfs that may be interacting with stars evolving along the RGB (Soker 1998). It is believed that  the mass loss rate  of RGB stars would be increased by the release of energy and angular momentum on entering the envelope of the RGB star. In the reported scenario, increasing mass loss would cause a star to become bluer as it starts towards the HB  (Soker 1998).  Speculatively, one could imagine that this might preferentially select short RR Lyrae star periods, and that the efficiency of such a mechanism might well vary significantly with radial distance in a star cluster.  \\begin{figure}% \\centering {\\includegraphics[angle=0,height=8cm,width=8cm]{./ButlerFig1.ps}} \\caption{ (a) Radial distribution of  the periods of Blazhko stars (squares), and non-Blazhko ab-type stars (diamonds) at r $\\ge$ 0.83$^{\\prime}$.   The short-dash line marks the median period of the ensemble data. (b) Standard deviation  of periods in annular bins as a function of radial position for  Blazhko plus non-Blazhko stars (solid) and  non-Blazhko stars (dashed). See the text in Sec.~\\ref{sample1} for further information. }\\label{fig_rad_sig} \\end{figure}   \\begin{figure}% \\centering \\scalebox{0.5}{\\includegraphics[angle=0]{./ButlerFig2.ps}} \\caption{Frequency of occurrence of $\\sigma_{\\rm P}$  at r = 3.5$^{\\prime}$ to 6$^{\\prime}$ for non-Blazhko ab-type stars (dashed) as well as Blazhko plus non-Blazhko ab-type stars (solid). Arrows indicate the measured dispersion in the same radial range for both the former case  (long) and the latter case (short).}\\label{fig_prob_RRLy_A_VB} \\end{figure}   \\begin{table} \\begin{tabular}{llll} \\hline \\rm Annulus         & \\rm  N$_{\\rm RRab}$$^{\\rm a}$   & \\rm N$_{\\rm non-Blazhko}$ &  \\\\ &  & &  \\\\ (1)  &   (2)          &   (3)  \\\\ \\hline 0.83$^{\\prime}$-2.00$^{\\prime}$  & 40 &  25  \\\\ 2.00$^{\\prime}$-3.50$^{\\prime}$  & 29 &  21\\\\ 3.50$^{\\prime}$-6.00$^{\\prime}$  & 22 &  11    \\\\ 6.00$^{\\prime}$-9.00$^{\\prime}$  & 21 & 15  \\\\ \\hline \\end{tabular} \\noindent \\\\ $^{\\rm a}$   Non-Blazhko and Blazhko  ab-type RRLyrae stars \\caption[]{Numbers of RR Lyrae stars  in different annular regions in M3. Columns (2) to (3) have been produced using RR Lyrae star classifications from CC01/Bakos et al. (2000). See text in Sec.~\\ref{sample1} for further information.}\\label{tableRRLy} \\end{table}"
    },
    "0404/astro-ph0404314_arXiv.txt": {
        "abstract": "A multi-object spectroscopy follow-up survey of galaxy clusters selected from the Red-sequence Cluster Survey (RCS) is being completed.  About forty clusters were chosen with redshifts from 0.15 to 0.6, and in a wide range of richnesses.  One of the main science drivers of this survey is a study of internal dynamics of clusters. We present some preliminary results for a subset of the clusters, including the correlation of optical richness with mass, and the mass-to-light ratio as a function of cluster mass. ",
        "introduction": "The RCS is a large imaging survey in $R$ and $z^\\prime$, which uses galaxy overdensity along with colour information to efficiently detect clusters out to $z\\!\\!>\\!\\!1$ (\\cite[Gladders \\& Yee 2004]{rcs}).  Detection of a cluster with the RCS technique also provides $z_{phot}$, a photometric estimate of the cluster redshift.  In order to study cluster dynamics and galaxy populations, we are completing a spectroscopic follow-up survey of a subset of about $40$ RCS clusters.  These clusters have a wide range of optical richnesses (parametrized by the cluster-centre galaxy correlation amplitude $B_{gc}$; see \\cite[Yee \\& Ellingson 2003, hereafter Y03]{y03}), and are chosen with redshift estimates $z<0.6$ so as to allow efficient follow-up.  We have performed multi-object spectroscopy of galaxies in the fields of these clusters using the LDSS-2 spectrograph on the Magellan telescopes and the MOS/SIS on the CFHT.  We observe two masks per field, for improved spatial sampling as well as some redundant observations (which help in estimating redshift errors).  In order to obtain spatial coverage out to a few virial radii for each cluster, we observe between one and three fields depending on the redshift and richness of the cluster.  Our selection criteria for mask design depends only on geometry and magnitude: we make no colour cuts, so as to ensure a fair sample for use in studying galaxy populations in these clusters.  Overall, with two to six masks per cluster, roughly 30 spectra per mask with Magellan/LDSS-2 and 100 spectra per mask with CFH/MOS (thanks to a bandlimiting filter which we use to multi-tier our spectra), we have over 7000 spectra in total.  Of course we cannot obtain redshifts from all spectra, nor are all the galaxies cluster members; ultimately we expect to have about 1750 members for the clusters in this survey.  \\begin{figure} \\centerline{ \\includegraphics[width=0.4\\textwidth]{blindert.fig1a.eps} \\includegraphics[width=0.4\\textwidth]{blindert.fig1b.eps} \\\\} \\caption{Pie diagrams of two example clusters; diagram (a) shows a poor cluster, while (b) is one of our richest systems. Each square is a galaxy, with size related to apparent magnitude.} \\label{pieplots} \\end{figure} ",
        "conclusions": "We are completing a spectroscopic follow-up survey of about 40 RCS clusters in a range of redshifts and richnesses.  We have presented some first results from 17 of these systems: \\begin{itemize} \\item 16 of the systems are confirmed as real overdensities in redshift space.  Data for the remaining system are too poor to determine whether it is real or not. \\item Predicted redshifts from the RCS method are in good agreement with the spectroscopic redshifts.  Optical richnesses also correlate with velocity dispersions, though weakly. \\item Cluster mass-to-light ratio depends on cluster mass.  This anti-biasing effect is in qualitative agreement with other datasets. \\item Scaled cluster-centric radii and velocities, shown for a single stacked ensemble, show a clear caustic. \\end{itemize}  The analysis we have presented here is quite simple, and several obvious improvements can be made.  Our interloper removal scheme is likely too strict, and we currently use only an estimate of $r_{200}$ in place of the virial radius.  Also, our masses are computed from the virial equation, neglecting the surface term; thus they overestimate the true masses enclosed within $r_{200}$ (C96).  We have 20 more systems to analyse, which will allow us to constrain differences in cluster properties such as mass and mass-to-light profiles with richness."
    },
    "0404/astro-ph0404408_arXiv.txt": {
        "abstract": "Following Tancredi, S\\`anchez and Roig (2001)'s criteria of chaos, two ways of setting initial velocities are used in the numerical surveys to explore the possible chaotic and regular orbits for the disc-star-planet systems. We find that the chaotic boundary does not depend much on the disc mass for Type I initial condition, but can change a lot for different disc masses for Type II initial condition. A few sample orbits are further studied. Both Poincar\\'e surface of section and the Lyapounov Exponent Indicator are calculated and they are consistent with each other. We also find that the influence from the disc can change the locations of equilibrium points and the orbital behaviors for both types of initial conditions. Because the chaotic orbits are less likely to become the stable resonant orbits, we conclude that the proto-stellar disc shall play important roles for the capture and depletion histories of resonant orbits of both Asteroid Belt and Kuiper Belt during the formation of Solar System. ",
        "introduction": "In the recent years, more than 100 extra-solar planetary systems are discovered and most of them exhibit interesting features. Many results have been presented in the dynamical studies trying to explain the observed orbital properties (for example, Laughlin \\& Adams 1999, Rivera \\& Lissauer 2000, Jiang \\& Ip 2001, Ji et al. 2002). Among these previous works, the restricted three-body problem is often used as a model for these planetary systems.  On the other hand, some of the discovered extra-solar planetary systems are claimed to have discs of dust and these discs are regarded as the young analogues of the Kuiper Belt. For instance, Greaves et al. (1998) found a dust ring around a nearby star $e$ Eri and Jayawardhana et al. (2000) detected the dust in the 55 Cancri planetary system. Particularly, $\\beta$ Pictoris planetary system has a warped disc and the influence of a planet might explain this warp (Augereau et al. 2001). Jiang \\& Yeh (2003) studied the effect of discs on the planetary orbits and concluded that the planets might prefer to stay near the inner part rather than the outer part of the disc.  Therefore, the modified restricted three-body problem proposed in Yeh \\& Jiang (2003) is interesting and might have important astronomical applications. In that model, the standard restricted three-body problem is generalized to the situation that the test particles are influenced by the additional gravitational forces from the disc. The influence from the disc makes the structure of equal-potential surfaces much more complicated and thus, the new equilibrium points exist in addition to the usual Lagrangian points. The conditions for the existence of these new equilibrium points were studied both analytically and numerically in that paper.  Having found the new equilibrium points, Jiang \\& Yeh (2004a, 2004b) studied the orbital behavior for this modified restricted three-body problem. The sensitivities to the initial conditions for some sample orbits clearly depend on the disc mass.  In this paper, we focus on the case that the central binary is a Sun-Jupiter system for the modified restricted three-body problem. Thus, it is a disc-star-planet system with the test particle moving around. Although the current masses of Asteroid Belt and Kuiper Belt in the Solar System are small, it could be much more massive during the early stage. Thus, our study might  have important implications for the formation of the dynamical structure near asteroid belts of any planetary systems. Moreover, the chaotic behaviors of the test particles in this system might be essential. Indeed, the chaotic orbits do play roles for planetary dynamical properties. For example, Wisdom (1983) successfully showed that the 3:1 Kirkwood gap of the asteroid belt might be related to the chaotic region of the system. Moreover, Astakhov et al. (2003) used the model of three-body interaction to explain the irregular satellites of the Jupiter. The possible chaotic orbits were claimed to be the reason for some Jupiter' satellites having irregular orbits.  We present our models in Section 2. Two types of initial conditions will be described in Section 3. The method to calculate the Lyapounov Exponent Indicator is summarized in Section 4. The results will be in Section 5 and 6 and Section 7 concludes the paper. ",
        "conclusions": "We have studied the chaotic orbits for the disc-star-planet systems through the calculations of Lyapounov Exponent and also Tancredi et al. (2001)'s criteria of chaos. Two types of initial conditions are used in the numerical surveys to explore the possible chaotic and regular orbits and thus determine the boundaries. We find that the chaotic boundary does not depend on the disc mass much for Type I initial condition, but can depend on the disc mass for Type II initial condition. That is, discs with different masses might change the sizes and locations of chaotic region. Some sample orbits are further studied and we find that the plots of Poincar\\'e surface of section are consistent with the results of Lyapounov Exponent Indicator.  On the other hand, we notice that the influence of the disc can change the locations of equilibrium points and also the orbital behaviors for both types of initial conditions. This point is particularly interesting. Without the influence from the disc, Wisdom (1983) showed that one of the Kirkwood gaps might be explained by the chaotic boundary. This is probably due to that the chaotic orbits are less likely to become resonant orbits. Indeed, we do find chaotic orbits in our disc-star-planet system and the disc, which can change orbital behaviors etc., shall play essential roles during the formation of Kirkwood gaps. For example, the disc might make some test particles to transfer from chaotic orbits to regular orbits or from regular orbits to chaotic orbits during the evolution.  The discs are often observed with various masses in both the young stellar systems and extra-solar planetary systems. It is believed that the proto-stellar discs are formed during the early stage of star formation and the planets might form on the discs through core-accretion or disc instability mechanism. The planetary formation time scale is not clear but the gaseous disc will be somehow depleted and the dust particles will gradually form. The above picture is plausible but unfortunately the time scales are completely unknown. For a system with a star, planets, discs and asteroids, it is not clear what is the order of different component's formation and what the duration of each component's formation process would be.  Nevertheless, one thing we can confirm is that the discs are there and can have different mass at different stages. From our results, it is clear that the dynamical properties of orbits, i.e chaotic or regular, can be determined quickly for different disc masses or models. Because those particles moving on chaotic orbits might have smaller probability to become the resonant objects in the system, the Kirkwood gaps, which are associated with the location of mean motion resonant, in the Solar System might be formed by the dynamical effect we have studied here. The detail processes would depend on the evolutionary time scale of the star, planet and also the disc. To conclude, it is important that the effect of discs on the chaotic orbits might influence the formation of the structures of planetary systems such as the Kirkwood gaps in the Solar System."
    },
    "0404/astro-ph0404122_arXiv.txt": {
        "abstract": "s{ We investigate the role of a possible coupling of dark matter and dark energy. In particular, we explore the consequences of such an interaction for the coincidence problem, i.e., for the question, why the energy densities of dark matter and dark energy are of the same order just at the present epoch. We demonstrate, that, with the help of a suitable coupling, it is possible to reproduce any scaling solution $\\rho_X \\propto \\rho_M a^\\xi$, where $a$ is the scale factor of the Robertson-Walker metric and $\\xi$ is a constant parameter. $\\rho_X$ and $\\rho_M$ are the densities of dark energy and dark matter, respectively. Furthermore, we show that an interaction between dark matter and dark energy can drive the transition from an early matter dominated era to a phase of accelerated expansion with a stable, stationary ratio of the energy densities of both components.} ",
        "introduction": "Our Universe seems to be made of roughly two third of something which is now called dark energy (DE) and of about one third of mainly cold dark matter (CDM). Usually, these components are considered as independent ingredients of the cosmic substratum. Given, that we neither know the nature of DE nor that of CDM, one cannot exclude that there exists a coupling between them. Such a situation is even the more general one compared with the uncoupled case. Here we explore some of the consequences that different types of interactions may have for the dynamics of the Universe. ",
        "conclusions": ""
    },
    "0404/astro-ph0404252_arXiv.txt": {
        "abstract": "{\\citet{gizis} reported a companion candidate of spectral type L1 near the K2 dwarf GJ 1048 using the Two Micron All-Sky Survey (2MASS). At that time it was not possible to verify companionship astrometrically using the 2MASS data alone due to the small proper motion of GJ 1048. We now show that both objects share the same proper motion by using data from the UK Schmidt Telescope Near-infrared (IVN) Southern Survey as the first epoch and data from 2MASS as the second epoch. Our technique of subtracting the PSF of the primary from the SuperCOSMOS I scans of the Southern Survey enables the astrometry of the companion candidate to be measured directly. ",
        "introduction": "Brown dwarfs are objects with masses below the mass limit for core hydrogen burning. They encompass a range of spectral types, from late M, for young brown dwarfs, down to the new spectral types L and T, defined by \\citet{kirk99}, \\citet{martin99} and \\citet{burg00}. So far a considerable number of such objects has been found, mainly by infrared (IR) surveys like the Two Micron All-Sky Survey (hereafter 2MASS) and the DEep Near Infrared Survey (hereafter DENIS). Although many field brown dwarfs are known, few are known as companions to stars. The majority of these stars are objects in the solar neighborhood and show a high proper motion. Confirmation of companionship can therefore be realized by determining a common proper motion of the star and the companion candidate. The significance and ease of this method generally increases with increasing proper motion and epoch difference between the observations.  ~\\citet{gizis} identified the 2MASS object 2MASSI J0235599--233120 as a companion candidate 11\\farcs9 separated from the K2 dwarf GJ1048. The JHK magnitudes of the object were consistent with an early L dwarf at the same (HIPPARCOS measured) distance of $21.3\\pm0.47$~pc as GJ1048. Thus, it is likely that the object is a companion of GJ1048. Gizis et al. referred to the companion as GJ1048B, the primary as GJ1048A. If GJ1048B is indeed a companion, then its mass (at the same age and distance as the primary) is near the sub-stellar limit \\citep{gizis}.  Due to the comparably small proper motion of GJ1048 ($\\mu_{\\alpha}=84.53\\pm0.98~$mas/yr, $\\mu_{\\delta}=13.04\\pm0.76~$mas/yr) it has so far not been possible to verify the companionship astrometrically. Since the epoch difference between the two major IR surveys 2MASS and DENIS is only in the range of a few years, stars showing only moderate proper motion like GJ 1048 are not suitable for a astrometric comparison within both mentioned surveys. It is therefore helpful to look for older red or IR surveys, sensitive enough to allow the detection of the companion candidates. The strength of this method is demonstrated by the discovery of $\\epsilon$ Indi B \\citep{scholz} on SuperCOSMOS Sky Survey images.  We report the missing confirmation of common proper motion companionship for GJ1048B. ",
        "conclusions": "Both methods presented in this paper show a consistent result for the common proper motion of GJ1048 A and B.  In combination with the previous results from \\citet{gizis} there can be no reasonable doubt that GJ1048B is physically associated with GJ1048A. Therefore the probability that GJ1048B is not a companion of GJ1048A is very small. However the final proof of this result has to be postponed until more of the orbital motion is detectable."
    },
    "0404/astro-ph0404587_arXiv.txt": {
        "abstract": "In this  second paper dedicated to the study of satellite galaxies we present broad-band photometry in the $B$, $V$, $R$ and $I$ filters of 49 satellite galaxies orbiting giant isolated spiral galaxies. First analysis of the properties of these objects are presented by means of color--color and color--magnitude diagrams for early- and late-type satellites.  Although we find differences in the slope of the $V-I$ vs. $M_v$ color magnitude diagram, as a whole, the relations are in agreement with the trends known to date for galaxies of similar magnitudes in nearby clusters of galaxies. Comparison with the relations found for satellites in the Local Group allows us to sample better the bright end of the luminosity function of satellite galaxies and extends for brighter objects the validity of the color--magnitude relation found for dwarf galaxies  in the Local Group. Most of the E/S0 galaxies in our sample show a negative color gradient with values similar to those known for early type galaxies in other environments. ",
        "introduction": "The study of the small scale structure in the Universe is a much discussed topic in current astronomy. One of the reasons for this wide interest is the discrepancies  between observations and the predictions of the standard cold dark matter (CDM) model. For instance, from $n$-body simulation (Moore et al.\\ 1999; Klypin et al.\\ 1999) the expected  number of satellites orbiting galaxies like the Milky Way or Andromeda is an order of magnitude larger than that observed in the Local Group (e.g., Mateo 1998). This is a potentially strong objection against the hierarchical scenario proposed in the standard CDM model. To solve this discrepancy several mechanisms that suppress the formation of satellites after the re-ionization  in the early epoch of the Universe have been proposed (e.g., Bullock, Kravtsov, \\& Weinberg 2000). So far, it  is unclear whether this or any of the other mechanisms proposed is able to reconcile fully the predictions from the models and the observations.  Satellite galaxies are also interesting for tracing the gravitational potential and   for estimating the mass of the parent galaxy at distances unreachable with other methods (Erickson, Gottesman \\& Hunter 1999). The standard models predict a decline in satellite galaxy velocity with distance to the primary. This has been observationally explored by Zaritsky et al.\\ (1997) and Prada et al.\\ (2003). The predicted effect was not found in the analysis of the first  group. However, Prada et al., with better statistics and the removal of interlopers,  have claimed the existence of this decline in the way predicted by the standard models.  Our knowledge of satellite galaxies beyond the Local Group is still very limited, owing to the intrinsic faintness of these objects. For instance, in the Local Group the known dwarf satellite galaxies have brightnesses in the range $-18\\le M_B \\le -8$. So detailed studies of external systems are limited to galaxies in the nearby Universe. Even surveys such as the Sloan Digital Sky Survey (Prada et al.\\ 2003) are  capable to sampling only the very bright part of the luminosity function of these objects, typically detecting   1--2 satellites orbiting each giant galaxy. For instance, the external isolated galaxies in which a larger number of galaxies have been found are NGC~1961 (Gottesman, Hunter \\& Shostak 1993) and NGC~5084 (Carignan et al.\\ 1997) in which 5 and 9 satellites respectively have been catalogued.  Bearing all the above points in mind, we decided to start a study of satellite galaxies orbiting external isolated galaxies. We extracted our  sample from the compilation of  such objects by Zaritsky et al.\\ (1997). Only spiral parents were included in that catalogue. A similar catalogue for parents with elliptical morphologies is being compiled by Smith \\& Martinez (2003). With photometry in optical broad ($BVRI$) and narrow bands  (H$\\alpha$), and near infrared ($J$ and $K$) of both the satellites and their parent galaxies we address the following questions: i) what are the properties (statistics, luminosity function,  morphology, structure, colors, etc.) of the satellite galaxies, and how do they compare with those found in the Local Group;  ii) how  the above properties are related to the relative positions of the satellite galaxies with respect to their parents; iii) what  the possible interactions are between satellites, and between satellites and their parents; and iv) the star formation rate and how it is related with the overall properties and relative position of the satellites. In the first paper of this series, Gutierrez, Azzaro, \\& Prada (2002, hereafter Paper~I), we presented the structural parameters and morphological classification of $\\sim 60$ of such objects.  Here, we present the photometry and first analysis of the colors of most of them. A full analysis combining structure and photometry and testing against the predictions of theoretical models will be conducted and presented in the near future. ",
        "conclusions": "\\begin{enumerate} \\item We have presented integrated photometry of 49 satellite galaxies orbiting external isolated spiral galaxies.  \\item The $B-V$ vs.\\ $M_V$ color magnitude relation is similar to the one found in the Fornax Cluster for galaxies with similar magnitudes. However the $V-I$ vs.\\ $M_V$ relation is steeper in our sample as compared with galaxies in Fornax.  \\item The $B-V$ vs.\\ $M_B$ color diagram extends to brighter magnitudes the relations found for satellite galaxies in the Local Group.  \\item We have measured the internal color gradients of nine early-type satellites. Excluding a galaxy with a strong positive gradient, we found $\\Delta (B-R)/\\Delta (\\log r)=-0.085\\pm 0.027$. \\end{enumerate}"
    },
    "0404/astro-ph0404064_arXiv.txt": {
        "abstract": "s{We calculate the atmospheric structure of an accretion disk around a Kerr black hole and obtain its X-ray spectrum, which exhibits prominent atomic transitions under certain circumstances. The gravitational and Doppler (red)shifts of the C~V, C~VI, O~VII, O~VIII, and Fe~I--XXVI emission lines are observable in active galaxies.  We quantify the line emissivities as a function of radius, to identify the effects of atmospheric structure, and to determine the usefulness of these lines for probing the disk energetics. The line emissivities do not always scale linearly with the incident radiative energy, as in the case of Fe~XXV and Fe~XXVI. Our model incorporates photoionization and thermal balance for the plasma, the hydrostatic approximation perpendicular to the plane of the disk, and general relativistic tidal forces. We include radiative recombination rates, fluorescence yields, Compton scattering, and photoelectric opacities for the most abundant elements.} ",
        "introduction": "One of the key signatures of strong gravity in astrophysical black holes is the peculiar shape of the Fe K$\\alpha$ emission line observed at 6.4~keV in active galactic nuclei and black hole X-ray binaries.\\cite{rev} The Fe K$\\alpha$ emission has been extensively studied in model calculations and Monte Carlo radiation transfer models. In such models, the Fe K$\\alpha$ line emission is produced by near-neutral plasma in the accretion disk that is irradiated by an external source of nonthermal X-rays and $\\gamma$-rays, located somewhere above the accretion disk. The Fe K$\\alpha$ line is skewed and redshifted down to $\\sim 3$ keV, but it only exhibits a small blueshift. Its line profile has been described by a raytracing model of a luminous disk in the Kerr metric.\\cite{laor}  The Fe K$\\alpha$ line provides evidence for an accretion disk extending very close to the event horizon of the black hole (however, alternative models exist, i.e. see Titarchuk in these proceedings). Although the origin of the nonthermal radiation bathing the disk is not fully understood, it is likely produced by inverse Compton scattering of $\\sim 100$~keV electrons with thermal UV radiation produced by the disk.  The relativistically-broadened emission lines of C~VI, N~VII, and O~VIII Ly$\\alpha$ have been tentatively identified in a handful of active galaxies.\\cite{branduardi} This identification has been controversial since the same X-ray spectra can be interpreted with photoelectric absorption features produced by plasma that is known to exist several parsecs away from the black hole.\\cite{lee} Our primary goal is to produce a theoretical prediction for such a recombination spectrum, accounting for relativistic effects. Since the raytrace calculation by Laor\\cite{laor} omitted the physics of the line emission mechanism in the disk, it introduced a power-law emissivity $\\epsilon \\propto r^{-q}$ as a function of radius $r$, with $q$ as a free parameter. Another purpose of our work is to test whether or not this line emissivity index $q$ is directly related to the emissivity index of the total radiative energy of the disk, $q_{T}$, as a function of $r$. The Shakura \\& Sunyaev\\cite{ss73} disk model implies that $q_{T} =3$, but in situations where the disk is heated by a coupling with the Kerr hole, $q_{T} = 3.5$.\\cite{agolk}  However, fits to the lines observed with the {\\it XMM-Newton} X-ray observatory show emissivity indices for Fe K$\\alpha$ of $q =4.3$--5.0\\cite{wilms}, that are larger than the theoretical predictions of $q_T$. Could $q \\neq q_T$ be due to atmospheric effects in the accretion disk? We test this hypothesis for the fluorescence and recombination X-ray lines of the most abundant elements. ",
        "conclusions": ""
    },
    "0404/astro-ph0404070_arXiv.txt": {
        "abstract": "Two FORTRAN programs are presented which plot the Hertzsprung-Russell diagram and the temporal evolution of such stellar quantities as: central and photospheric isotopic abundances, central densities and temperatures, luminosities,effective temperatures and photospheric radii for a wide range of stellar masses. The programs, which are modifications and extensions of some modules of the {\\it TYCHO 6.0} stellar evolution package, are accompanied by various control input files as well as by a library of data. The data library is actually one of the output files generated by {\\it TYCHO 7.0} (a modified version of {\\it TYCHO 6.0}), while the plots themselves are generated by the PGPLOT Graphics Subroutine Library which is also publicly available and well cited. ",
        "introduction": "{\\it TYCHO 6.0} \\cite{tycho60} is a general, one dimensional (spherically symmetric) stellar evolution code, designed for hydrostatic and hydrodynamic stages, using state of the art procedures and microphysics. It is written in structured FORTRAN (f77) and has extensive on-line graphics using the PGPLOT graphics package\\cite{pgplot} in X-windows environments. It has being developed as an open-source,community code and can run very effectively on Linux laptop and desktop computers. {\\it TYCHO 6.0} is being used by the University of Arizona astrophysics group (as well as by the author) for ongoing research in stellar evolution (for relevant papers see \\cite{tycho60} ) while it has also been used to generate low-noise initial models for multi-dimensional numerical hydrodynamic simulations. The {\\it TYCHO 6.0} package includes various auxiliary files and subprograms which can monitor both the evolution of a star and nuclear reaction networks under various conditions.  Although {\\it TYCHO 6.0's} auxiliary programs focus on various aspects of stellar evolution, such as the Hertzsprung-Russell (HR) diagram or the composition and thermodynamics of each evolutionary sequence, there is no auxiliary program which can investigate the temporal evolution of the star (neither during nor after the actual simulation process). Since typical HR diagrams and static composition profiles are actually bereft of any temporal information reliable programs are needed to satisfy this obvious need for all {\\it TYCHO 6.0} users. Moreover, such programs should not just be able to be used by TYCHO users. Anyone wishing to study the evolution of a star could benefit from an independent program which could be applied to a large library of data generated by TYCHO. This is the purpose of the programs presented here.  In the sections that follow we first describe the programs and then its input and output files. Then we implement the programs by performing a test-run for a 3M star and a particular initial composition. The present paper initiates an effort to improve the {\\it TYCHO 6.0} package and extend its capabilities in order to investigate various nuclear astrophysics topics. A similar effort is going on by the university of Arizona astrophysics group which has already derived a new version, called {\\it % TYCHO 6.11}, which is also publicly available\\cite{tycho60}. {\\it TYCHO 6.0} has now been equipped by the author with screening effects in the derivation of thermonuclear reaction rates as well as with various new auxiliary programs such as those presented in this paper. Screening corrections in the equation of state for completely degenerate environments (e.g. White-Dwarf progenitors) will soon be available These aspects are not included in the {\\it TYCHO 6.11} version which, however, has very interesting improvements over {\\it TYCHO 6.0}. To facilitate versioning and differentiation from the {\\it TYCHO 6.11} sequence (and all relevant work from the Univ.of Arizona group), the version independently evolved and implemented here will be referred to as {\\it TYCHO 7.0}, while all relevant improvements and modifications over the {\\it TYCHO 6.0} version will be consistently recorded and published, starting with the present paper. Note that whenever the stellar evolution code is referred to as {\\it TYCHO} then the relevant comment is applicable to all versions of the code. ",
        "conclusions": ""
    },
    "0404/astro-ph0404593_arXiv.txt": {
        "abstract": "The enrichment of Fe, relative to alpha-elements such as O and Mg, represents a potential means to determine the age of quasars and probe the galaxy formation epoch.  To explore how \\ion{Fe}{2} emission in quasars is linked to physical conditions and abundance, we have constructed a 830-level \\ion{Fe}{2} model atom and investigated through photoionization calculations how \\ion{Fe}{2} emission  strengths depend on non-abundance factors. We have split \\ion{Fe}{2} emission into three major wavelength bands, \\ion{Fe}{2} (UV), \\ion{Fe}{2}(Opt1), and \\ion{Fe}{2}(Opt2), and explore how the  \\ion{Fe}{2}(UV)/\\ion{Mg}{2}, \\ion{Fe}{2}(UV)/\\ion{Fe}{2}~(Opt1) and  \\ion{Fe}{2}(UV)/\\ion{Fe}{2}~(Opt2) emission ratios depend upon hydrogen density and ionizing flux in broad-line regions (BLR's) of quasars.  Our calculations show that: 1) similar \\ion{Fe}{2}(UV)/\\ion{Mg}{2} ratios can exist over a wide range of  physical conditions; 2) the \\ion{Fe}{2}(UV)/\\ion{Fe}{2}~(Opt1) and \\ion{Fe}{2}(UV)/\\ion{Fe}{2}~(Opt2) ratios serve to constrain ionizing luminosity and hydrogen density; and 3)  flux measurements of \\ion{Fe}{2} bands and knowledge of ionizing flux provide tools to derive distances to BLR's in quasars. To derive all BLR physical parameters with uncertainties, comparisons of our model with observations of a large quasar sample at low redshift ($z<1$) is desirable.  The  STIS and NICMOS spectrographs aboard the Hubble Space Telescope (HST) offer the best means to provide such observations. ",
        "introduction": "The observed ratio of restframe \\ion{Fe}{2}~UV emission from 2200$-$3000~{\\AA} (thereafter, \\ion{Fe}{2}(UV)) to that of the \\ion{Mg}{2} 2800~{\\AA} resonance doublet has been widely used to estimate Fe/Mg abundance ratios in BLRs of quasars (Wills et al. 1985, thereafter WNW85; Iwamuro et al.  2002; Dietrich et al. 2002, 2003; Freudling et al. 2003; Barth et al. 2003; Maiolino et al. 2003). Iron enrichment is not expected in evolution scenarios at redshifts $z \\leq 1$ (Hamann \\& Ferland 1993; Yoshii et al. 1996; Heger \\& Woosley 2002). Meanwhile, there is some observational evidence that narrow line AGN have higher metalicities compared to other low redshift AGN (Shemmer \\& Netzer 2002).  Measurements of \\ion{Fe}{2}(UV)/\\ion{Mg}{2}  emission ratios in quasars show large scatter from 1 to 20, and no redshift dependency up to $z \\sim 6.4$ (Iwamuro et al. 2002).  In a recent paper, we have used a sophisticated 830-level model for \\ion{Fe}{2}   to investigate abundance and microturbulence effects on these ratios (Verner et al. 2003a). Our modeling indicates strong \\ion{Fe}{2} emission from 1000 to 6800~\\AA, and reveals that \\ion{Fe}{2} (UV)/\\ion{Mg}{2} ratios (at {$n_H$ }= 10$^{9.5}$~cm$^{-3}$ and total hydrogen column density $N_{H}$ = 10$^{24 }~$cm$^{-2}$) increase from 5 to 30 for microturbulence $v_{turb}$ varied from 1 to 100 km~s$^{-1}$, while an increase in abundance by factor 10 only increases the same ratios by factor 2. We have further found that a reasonable range of microturbulence is between 5$-$10 km ~s$^{-1}$, and \\ion{Fe}{2} optical emission flux from 4000$-$6000~{\\AA}  is more sensitive to abundance than is the \\ion{Fe}{2}(UV) band. Conversely, the  \\ion{Fe}{2}(UV) band is more sensitive to  microturbulence than the \\ion{Fe}{2} optical band.  In this paper we continue our study of \\ion{Fe}{2} emission line formation in BLRs of quasars in an attempt to explain the nature of the large observed scatter. Since abundances, density, and excitation conditions are ill-determined in BLRs, we have performed a set of numerical calculations to study the effects of varying hydrogen and ionizing photon densities. We have split \\ion{Fe}{2} emission into three  wavelength bands that are most commonly measured in observations: \\ion{Fe}{2}(UV) (2000$-$3000~{\\AA}), \\ion{Fe}{2}(Opt1) (3000$-$3500~{\\AA}), and \\ion{Fe}{2}(Opt2) (4000$-$6000~{\\AA}). The major goal is to ascertain if \\ion{Fe}{2}(UV)/\\ion{Mg}{2}, \\ion{Fe}{2}(UV)/\\ion{Fe}{2}(Opt1), and \\ion{Fe}{2}(UV)/\\ion{Fe}{2}(Opt2) emission ratios present different trends due to non-abundance factors.  Although many factors contribute to increase \\ion{Fe}{2} emission fluxes, the effects can be separated by using several ratios simultaneously. Spectra of quasars at redshifts up to $z = 1$ are especially valuable, since they should exhibit no elemental overabundances, and the effects of physical parameters such as hydrogen density, ionizing luminosity, and even microturburbulence can be easily explored. We just showed general trends, while a detailed comparison between model and observations for each quasar spectrum should be done individually. The figures in our paper represent calculations for a range of reasonable BLR conditions, but do not cover the full range of possible variables and should be used with extreme care by observers. Such comparison will provide the first fundamental test of the \\ion{Fe}{2} emission model. This consistency check is a necessary step before one can confidently apply this methodology to quasars over a wide range of redshift as a means to probe galactic chemical evolution models. The required data for quasars up to $z = 1$ can be obtained with STIS and NICMOS aboard the HST. ",
        "conclusions": "  1. Our \\ion{Fe}{2} large model atom used in photoionization calculations predicts that large \\ion{Fe}{2}(UV)/\\ion{Mg}{2} ratios are not necessarily due to high iron abundance. The \\ion{Fe}{2}(UV)/\\ion{Mg}{2} diagram (Figure 2) demonstrates that the \\ion{Fe}{2}(UV)/\\ion{Mg}{2} ratios can have the same value over a wide range of physical conditions. This implies that the averaging of large numbers of quasars spectra may lead to properties of quasars that do not exist.  2. Since evolutionary models predict no iron overabundances for z $\\leq$ 1, observations to detect the optical iron band at these redshifts will be beneficial to calibrate the model assuming a solar abundance for iron. Such research must be accomplished with STIS and NICMOS aboard the HST to properly span the wavelength range where \\ion{Fe}{2} emission is seen and  to obtain constraints on the most uncertain parameter, microturbulence.   3. Our \\ion{Fe}{2} emission band analysis demonstrates that \\ion{Fe}{2} has a great potential for becoming an important diagnostic tool of the physical conditions in BLR's of quasars. We suggest three \\ion{Fe}{2} emission ratios, which can be used to obtain hydrogen density, flux of hydrogen-ionizing photons, microturbulent velocities, and distances to BLR's.  In this paper we have shown that even in the limited space of parameters (ionizing flux and hydrogen density) exact fits to observed spectra exist. The full investigation of our model in the whole parameter space including dust properties, Balmer continuum, and many others will be a subject of coming papers. Such investigation will provide all BLR physical parameters with the determined uncertainties."
    },
    "0404/astro-ph0404136_arXiv.txt": {
        "abstract": " ",
        "introduction": "The earliest detections of luminous X-ray sources ($\\Lx \\gtap 10^{36}$~\\ergsec) in globular clusters were made with the Uhuru and OSO-7 Observatories (Giacconi et al.\\ 1972 \\& 1974; Clark, Markert \\& Li, 1975; Canizares \\& Neighbours, 1975). About 10\\%\\ of the luminous X-ray sources in our Galaxy are found in globular clusters. This implies that the probability (per unit mass) of finding a luminous X-ray source in a globular cluster is about two to three orders of magnitude higher than of finding one in the rest of our galaxy (Gursky 1973, Katz 1975). Clearly, the conditions in globular clusters are very special in that they must be very efficient breeding grounds for X-ray binaries.  For reviews which reflect the ideas in the late seventies and early eighties, see Lewin (1980), Lewin \\& Joss (1983), Van den Heuvel (1983) and Verbunt \\&\\ Hut (1987). At that time there was no evidence for a substantial population of binaries in globular clusters; e.g.\\ Gunn \\& Griffin (1979) did not find a single binary in a spectroscopic search for radial velocity variations of 111 bright stars in M\\,3.  Clark (1975) suggested that the luminous cluster sources are binaries formed by capture from the remnants of massive stars.  Fabian, Pringle \\& Rees (1975) specified that they are formed via tidal capture of neutron stars in close encounters with main-sequence stars. Sutantyo (1975) suggested direct collisions between giants and neutron stars as a formation mechanism. Hills (1976) examined the formation of binary systems through star-exchange interactions between neutron stars and primordial binaries of low-mass stars. Hut \\& Verbunt (1983) compared the relative efficiencies of tidal capture and exchange encounters for neutron stars and for white dwarfs; and showed that the distribution of X-ray sources among globular clusters with different central densities and core sizes is compatible with the formation by close encounters (Verbunt \\&\\ Hut 1987). The importance of mass segregation, which drives the neutron stars to the core, thereby enhancing the capture rate, was demonstrated by Verbunt \\&\\ Meylan (1988).  As can be seen from the discovery references in Table 8.1, five luminous globular cluster X-ray sources were known by 1975, eight by 1980, ten by 1982, and thirteen to date. Twelve of these have shown type~I X-ray bursts. Measurements of the black-body radii of the burst sources indicated that they are neutron stars (Swank et al.\\ 1977; Hoffman, Lewin \\& Doty 1977a \\& 1977b; Van Paradijs 1978). For a review, see Lewin, Van Paradijs \\& Taam (1995); see also \\S3\\mk. The absence of luminous accreting black holes in clusters of our galaxy is presumably a consequence of the small total number of sources, as discussed in \\S\\ref{vlsecbh}. There is \\mk growing evidence that black-hole binaries may exist in globular clusters in several elliptical galaxies (see \\S8.3).  Because of the observed correlation between the occurrence of a luminous X-ray source in a globular cluster and a \\mk high central density, it was expected already early on that these luminous sources would be located close to the cluster centers. These expectations were confirmed by measurements, carried out with the SAS-3 X-ray Observatory, which showed that the positional error circles with radii of 20--30 arcsec (90\\% confidence) included the optical centers of the clusters (Jernigan \\& Clark 1979). Later work with the Einstein observatory greatly refined the positional measurements (Grindlay et al.\\ 1984). Bahcall \\& Wolf (1976) have shown that under certain assumptions, the average mass of the X-ray sources can be derived from their positions with respect to the cluster center. Even if one accepts the assumptions made, the average mass derived this way for the luminous X-ray sources in \\gcs\\ was not sufficiently accurate to classify these sources, but the result was consistent with the earlier conclusions (see e.g., Lewin 1980; Lewin \\& Joss 1983) that these are accreting neutron stars (Grindlay et al.\\ 1984).  \\begin{table*} \\centerline{ \\begin{tabular}{l@{ }ll@{ }lllll} cluster   & position  & discovery         & 1st burst    & $M_{\\lambda}$   &   $P_b$ & TOXB \\\\ NGC\\,$1851$ & 0512$-$40\\,\\cite{ham+01} & O7\\cite{cml75} & UH\\cite{fj76} & 5.6B\\cite{damd96} & & \\phantom{T}UUU\\\\ NGC\\,$6440$ & 1745$-$20\\,\\cite{plv+02} & O7\\cite{mbc+75} & BS\\cite{zvs+99} & 3.7B\\cite{vkzh00} & & T\\phantom{}$-$N$-$& \\\\ NGC\\,$6441$ & 1746$-$37\\,\\cite{ham+02} & UH\\cite{gmg+74} & EX\\cite{sfp+87} & 2.4B\\cite{damd98} & 5.7hr\\cite{sdal93} &  \\phantom{T}$-$NN \\\\ NGC\\,$6624$ & 1820$-$30\\,\\cite{ksa+93} & UH\\cite{gmg+74} & ANS\\cite{ggs+76} & 3.0B\\cite{amd+97} & 11.4m\\cite{spw87} & \\phantom{T}UUU \\\\ NGC\\,$6652$ & 1836$-$33\\,\\cite{heg01} & H2\\cite{hw85} & BS\\cite{zvh+98} & 5.6B\\cite{heg01} & c & \\phantom{T}UUU\\\\ NGC\\,$6712$ & 1850$-$09\\,\\cite{ghs+84} & AV\\cite{sptp76} & S3\\cite{hcl80} & 4.5B\\cite{hcn+96} &  20.6m$^d$\\cite{hcn+96} & \\phantom{T}UUU \\\\ NGC\\,$7078$-1$^a$ & 2127+12\\,\\cite{wa01} & Ch\\cite{wa01} & & 0.7B\\cite{aft84} & 17.1hr\\cite{iak+93} & \\phantom{T}$-$$-$$-$\\\\ NGC\\,$7078$-2$^a$ & 2127+12\\,\\cite{wa01} & Ch\\cite{wa01} & Gi\\cite{dim+90} & 3.1U\\cite{wa01} & & \\phantom{T}$-$$-$U \\\\ Terzan\\,1  & 1732$-$30\\,\\cite{jvh95} & Ha\\cite{moi+81} & Ha\\cite{moi+81} &               & & T\\phantom{}$-$$-$$-$   \\\\ Terzan\\,2   & 1724$-$31\\,\\cite{ghs+84} & O8\\cite{sbb+77} & O8\\cite{sbb+77} & & & \\phantom{T}$-$NU\\\\ Terzan\\,5  & 1745$-$25\\,\\cite{heg+03} & Ha\\cite{moi+81} & Ha\\cite{moi+81} & 1.7J\\cite{heg+03} & & T\\phantom{}$-$U$-$ \\\\ Terzan\\,6  & 1751$-$31\\,\\cite{zhm+03} & RO\\cite{phv91} & BS\\cite{zhm+03}& & 12.36h\\cite{zbc+00} & T\\phantom{}$-$N$-$ \\\\ Liller\\,1  & 1730$-$33\\,\\cite{hdam01} & S3$^b$\\cite{ldc+76} & S3\\cite{hml78} & & & T\\phantom{}$-$$-$$-$ \\\\ \\end{tabular} } $^a$A luminous X-ray source in NGC\\,7078 was already found with Uhuru\\cite{gmg+74}, the Chandra observations resolved this source into two sources\\\\ $^b$X-ray source (the Rapid Burster) discovered before the globular cluster!\\\\ $^c$43.6\\,m period originally assigned to this source is period of fainter X-ray source \\cite{heg01}\\\\ $^d$or the alias period of 13.2\\,m   \\caption[o]{\\it Some information on the luminous X-ray sources in globular clusters of our galaxy. Columns from left to right (1) cluster, (2) rough position (B1950, often used as source name) with reference to the currently most accurate position, (3) the satellite with which the source was discovered as a cluster source, (4) the satellite which detected the first type~I X-ray burst from the source, (5) absolute magnitude with filter of optical counterpart, (6) orbital period, (7) indication (with a ``T'') whether the source is a transient. The last three columns indicate whether a normal (N) or ultrashort (U) orbital period is suggested by the comparison of optical with X-ray luminosity (column 8, under ``O''), the X-ray spectrum (column 9, under ``X'') and the maximum flux reached during bursts (column 10, under ``B''). A ``$-$'' in columns 8-10 indicates that no information is available. Satellite names are abbreviated as O(SO-)7, O(SO-)8, UH(URU), H(EAO-)2, A(riel-)V, Ch(andra), Ha(kucho), RO(SAT), S(AS-)3, B(eppo)S(ax), EX(OSAT), Gi(nga). Note that the absolute magnitudes are subject to uncertainties in distance and reddening; also most sources are variable (see Deutsch et al.\\ 2000). \\nocite{dma00} \\label{tbright} } \\end{table*}  \\begin{figure} \\centerline{ \\parbox[b]{7.0cm}{\\psfig{figure=verbfb.ps,width=7.0cm,clip=t}} \\parbox[b]{5.0cm}{\\caption[o]{\\it Various types of X-ray sources in globular clusters; sg, ms, wd, and ns stand for subgiant, main-sequence star, white dwarf, and neutron star, respectively. From top to bottom: luminous low-mass X-ray binary, low-luminosity low-mass X-ray binary, recycled radio pulsar (here with a white dwarf companion), cataclysmic variable, and magnetically active binary. $L_{sd}$ stands for spin-down luminosity. Approximate maximum luminosities (in the 0.5-4.5 keV range) are indicated on the right. The low-mass X-ray binaries harboring a neutron star are referred to as \\lmns; when they harbor a back hole, we refer to them as \\lmbh, and we refer to both groups together as \\lmxb. \\label{verbfb}} } } \\end{figure}  Sources with $L_x\\ltap10^{35}$\\ergsec\\ were first found in globular clusters with Einstein (Hertz \\&\\ Grindlay 1983). More were found with ROSAT, by a variety of authors (Table\\,\\ref{tdim}); a final, homogeneous analysis of the complete ROSAT data was made by Verbunt (2001). On the basis of these Einstein and ROSAT results, it has gradually become clear that these sources are a mix of various types (see Figure\\,\\ref{verbfb}). Hertz \\&\\ Grindlay (1983) suggested that they were mainly cataclysmic variables, and noted that the low-luminosity source in NGC\\,6440 could be the quiescent counterpart of the luminous transient source in that cluster. Verbunt et al.\\ (1984) argued that the more luminous of the low-luminosity sources are all quiescent low-mass X-ray binaries. The first radio \\mk pulsar detected as a low-luminosity source in a globular cluster is the pulsar in NGC\\,6626 (M\\,28, Saito et al.\\ 1997). Finally, Bailyn et al.\\ (1990) pointed out that magnetically active binaries also reach X-ray luminosities in the range of the less luminous sources detected with ROSAT.  It was also realized that some of the sources could be unresolved multiple sources; and unresolved emission was found e.g.\\ by Fox et al.\\ (1996) in NGC\\,6341 and NGC\\,6205. However, it is fair to say that the actual plethora of sources shown by the Chandra observations in virtually every cluster that it observed (Tables\\,\\ref{tdim}, \\ref{verbta}) was unpredicted. These observations confirmed that quiescent low-mass X-ray binaries, cataclysmic variables, pulsars, and magnetically active binaries are all X-ray sources in globular clusters, as is discussed in \\S\\ref{vlsecfs}. Whereas some of the Einstein and ROSAT sources are confirmed with Chandra as single sources, others have been resolved into multiple sources; details are given in Table 8.2.  The positions obtained with ROSAT were sufficiently accurate to find plausible optical counterparts in Hubble Space Telescope (HST) \\mk observations in a number of cases. This work was pioneered in NGC\\,6397 with a search for $H\\alpha$ emitting objects by Cool et al.\\ (1993, 1995), and spectroscopic follow-up \\mk confirming the classification as cataclysmic variables by Grindlay et al.\\ (1995), Cool et al.\\ (1998) and Edmonds et al.\\ (1999). Plausible candidate counterparts were also found for two X-ray sources in the core of $\\omega$\\,Cen (Carson et al.\\ 2000). All of these suggested counterparts were confirmed with the more accurate positions obtained with Chandra. In 47\\,Tuc, of the candidate counterparts suggested by a variety of authors, Verbunt \\&\\ Hasinger (1998) only retain three, on the basis of more accurate positions of the X-ray sources; these were also confirmed with Chandra. Ferraro et al.\\ (1997) suggested ultraviolet stars as counterparts for two sources found by Fox et al.\\ (1996) in NGC\\,6205 (see also Verbunt 2001). An ultraviolet counterpart suggested by Ferraro et al.\\ (2000) for a source in NGC\\,6341 is incompatible with the position of that source (Geffert 1998, Verbunt 2001). Another approach is to look for X-rays from an already known special object. A dwarf nova known since 1941 well outside the central region of NGC\\,5904 (Oosterhoff 1941) was detected with ROSAT (Hakala et al.\\ 1997), and a pulsar in M28 (Lyne et al.\\ 1987) was detected with ASCA (Saito et al.\\ 1997). Before Chandra, no magnetically active binary was suggested as an optical counterpart for a specific X-ray source.  \\begin{table*} \\centerline{ \\begin{tabular}{lllll} cluster   & E & R & C/X & comments \\\\ NGC\\,104/47Tuc & 1 &   5+4 & 39+66 & E=R9=C42 (CV) R5=C58 R7=C46\\\\ & & & &                  R6=C56 R10=C27 R11=C25 R13=C2\\\\ & & & &                  R19=C30 R4 outside C-frame\\\\ NGC\\,288       & & 1 & & \\cite{sir+99} \\\\ NGC\\,362 &       & 2  \\\\ Pal 2   &     & 0+1 & & \\cite{rdlm94} \\\\ NGC\\,1904/M79 & 1 & 0+1 & &  E=R \\\\ NGC\\,5139/$\\omega$Cen  &  1+4 & 3+3 & 3+97 & \\cite{vj00} core: EC$>$(R9a=C6/R9b=C4) \\\\ & & &                  & both CVs; out-of-core: EB=R7=C3 qLMXB; \\\\ & & &                  & EA=R3, ED=R4, EE=R5 foreground stars\\\\ NGC\\,5272/M3 & 1 &   1  & & \\cite{dag99} E=R, CV/SSS? also \\cite{hgb93}, opt.id.\\ \\cite{ekh04} \\\\ NGC\\,5824    & 1 &   0 & & R limit just below E detection level \\\\ NGC\\,5904/M5    & & 0+1 & 10$^a$ &  \\cite{hcjv97} \\\\ NGC\\,6093/M80 & & 1  & 9+10 & R$>$(C1/C2/C4/C7..) \\\\ NGC\\,6121/M4    & & 1 & 12+19 & R=C1 \\\\ NGC\\,6139      &  & 1  \\\\ NGC\\,6205/M13  & & 2+1 & 2+1 & core: RGa=X3 qLMXB; RGb$\\neq$X \\\\ & & & &                X2$\\neq$R out-of-core: RF=X6 \\\\ NGC\\,6266/M62 & &     1 & 45$^a$  \\\\ NGC\\,6341/M92 & &     1  & & \\cite{jvh94}, \\cite{flm+96} \\\\ NGC\\,6352     & &     0+1 & & \\cite{jvh96} \\\\ NGC\\,6366     & &     1 & 1$^a$ & \\cite{jvh96}  \\\\ NGC\\,6388   & & 0+1  \\\\ NGC\\,6397   & & 5+1 & 9+11 & R4a/b/c/d/e=C19/17/23/22/18 R13=C24\\\\ NGC\\,6440 & 1 & 2 & 24 & E$>$(R1$>$C2/C4/C5..,R2$>$C1/C3...) \\\\ NGC\\,6541  & 1 & 1   \\\\ NGC\\,6626/M28 & & 3+1  & 12+34 & core: (R2a+2b)=C26 R2c=C19 \\\\ & & & &               out-of-core: R7=C17\\\\ NGC\\,6656/M22 & 1+3 & 1 & 3+24 & core: E=R=X16, opt id.\\cite{ack03} X18$\\neq$R \\\\ & & & &             out-of-core: E prob.\\ not related to cluster\\\\ NGC\\,6752 & & 4+2 & 9+8 & core: R7a$>$C4/7/9 R21$>$C11/12/18 R7b=C1 \\\\ & & &  &             R22=C6 out-of-core: R6=C3 R14=C2\\\\ NGC\\,6809 & & 1  & & \\cite{jvh96} \\\\ NGC\\,7099 & & 0+1  & 5$^a$ & \\cite{jvh94} \\\\ \\\\ total:  &    8+7 & 37+18 \\\\ \\end{tabular} } $^a$number within half-mass radius from \\cite{pla+03}, detailed analysis not yet published  \\caption[o]{\\it Observations of low-luminosity sources in globular clusters. We list the number of sources found with Einstein (under E), ROSAT (R) and Chandra or XMM (C/X). Numbers following the + sign indicate sources outside the cluster core. Note that the detection limits are very different between clusters. References are Hertz \\&\\ Grindlay (1983) for Einstein sources, Verbunt (2001) and references therein for ROSAT sources. References for Chandra and XMM-Newton are listed in Table\\,\\ref{verbta}. Under comments we provide additional references for ROSAT, give occasional source types quiescent (i.e.\\ low-luminosity) low-mass X-ray binary (qLMXB) and cataclysmic variable (CV), and indicate the relation between sources observed by subsequent satellites. $=$ identical; $>$ resolved into multiple sources; A$\\neq$B source A not detected by satellite B, due to significant variability. Source numbers under comments are those in the references given. \\label{tdim} } \\end{table*}  The luminous \\mbox{X-ray} sources in globular clusters are binary systems, and most (if not all) of the low-luminosity X-ray sources are also binary systems or have evolved from them.  The presence of binaries is a very important factor in the evolution of a globular cluster (Hut et al.\\ 1992). Theoretical considerations and numerical calculations show that a cluster of single stars is unstable against collapse of its core (H\\'enon 1961). If binaries are present, however, close binary-single star encounters can increase the velocity of the single stars by shrinking the binary orbits.  Binaries can therefore become a substantial source of energy for the cluster, sufficient even to reverse the core collapse.  Even a handful of very close binaries can significantly modify the evolution of a globular cluster (Goodman \\& Hut, 1989).  With a million stars in the cluster as a whole, the number of stars in the core of a collapsed cluster may be only a few thousand.  A close binary system, such as an \\mbox{X-ray} binary, will have a binding energy that can easily be a few hundred times larger than the kinetic energy of a single star.  A dozen such systems, as they were formed, released an amount of energy that is comparable to the kinetic energy of the core as a whole.  Encounters between such binaries and other single stars or binaries have the potential to change the state of the core dramatically by increasing or decreasing the core size, and by kicking stars and binaries into the cluster halo or even out of the cluster altogether. The study of the binaries, and \\mbox{X-ray} binaries in particular, is therefore of great importance as they play a key role in the cluster's dynamical evolution.  It has been suggested that globular clusters are responsible for the formation of all or some of the low-mass X-ray binaries in our Galaxy, also those outside clusters now (e.g.\\ Grindlay \\&\\ Hertz 1985). Specifically, such an origin was suggested by Mirabel et al.\\ (2001) for the black-hole X-ray binary XTE\\,J1118+480, and by Mirabel \\&\\ Rodrigues (2003) for Sco X-1, on the basis of their orbits in the Galaxy. The discovery of very large populations of cluster X-ray sources in other galaxies has rekindled the question of cluster origin for non-cluster sources. In \\S\\ref{secfield} we will give the \\mk reasons why we believe that most X-ray binaries in the disk of our Galaxy were formed there, and not in globular clusters. \\nocite{gh85}\\nocite{mdm+01}\\nocite{mr03} ",
        "conclusions": ""
    },
    "0404/astro-ph0404300_arXiv.txt": {
        "abstract": "{ We study the population of compact X-ray sources in the Large Magellanic Cloud using the archival data of the XMM-Newton observatory. The total area of the survey is $\\approx 3.8$ square degrees with a limiting sensitivity of $\\approx 10^{-14}$ erg/s/cm$^2$, corresponding to a luminosity of $\\approx 3\\cdot 10^{33}$ erg/s at the LMC distance. Out of $\\sim 460$ point sources detected in the 2--8 keV energy band, the vast majority are background CXB sources, observed through the LMC. Based on the properties of the optical and near-infrared counterparts of the detected  sources we identified 9 likely HMXB candidates and 19 sources, whose nature is uncertain, thus providing lower and upper limits on the luminosity distribution of HMXBs in the observed part of LMC. When considered globally,  the bright end of this distribution is consistent within statistical and systematic uncertainties with extrapolation of the universal luminosity function of HMXBs. However, there seems to be fewer low luminosity sources, $\\log(L_X)\\la 35.5$, than predicted. We consider the impact of the  ``propeller effect'' on the HMXB luminosity distribution and show that it can qualitatively explain the observed deficit of  low luminosity sources.  We found significant field-to-field variations in the number of HMXBs across the LMC, which appear to be uncorrelated with the star formation rates inferred by the FIR and $H_\\alpha$ emission. We suggest that these  variations are caused by the dependence of the HMXB number on the age of the underlying stellar population. Using the existence of large coeval stellar aggregates in the LMC, we constrain the number of HMXBs  as a function of time $\\tau$ elapsed since the star formation event in the range of $\\tau$ from $\\sim  1-2$ Myr to $\\sim 10-12$ Myr. ",
        "introduction": "\\label{sec:intro}    First imaging X-ray observations of the \\object{Large Magellanic Cloud} with Einstein \\citep{wang91} and ROSAT  \\citep{hp99} observatories revealed a moderate population of several tens of X-ray sources associated with this closest neighbor of the Milky Way. While in the soft X-ray band a large fraction of these sources are supernovae remnants, at higher energies, above $\\sim$ few keV, X-ray binaries provide the dominant contribution. For example, based on the spectral hardness of bright ROSAT sources, \\citet{kahabka} identified in the entire LMC $\\sim 20-30$ probable X-ray binary candidates with luminosity exceeding $L_X\\ga {\\rm few}\\times 10^{35}$ erg/s.  The relative numbers of high and low mass X-ray binaries (H or LMXBs) are defined by the specific star formation rate  SFR/$M_*$ of the host galaxy \\citep{grimm03,gilfanov04}. Owing to rather low mass, $M_*\\sim2\\cdot10^{9}$ M$_{\\sun}$ (section \\ref{sec:lmxb}), and moderate  star formation rate (SFR) of $\\sim 0.5~M_{\\sun}$ yr$^{-1}$ (section \\ref{sec:sfr}), this quantity is rather high for the LMC, with a SFR/$M_*\\sim (2-3)\\cdot 10^{-10} {\\rm ~yr}^{-1}$,  exceeding by a factor of $\\sim 5-10$ that of our Galaxy. Correspondingly, the ratio of the expected numbers of bright, $\\log(L_X)\\ga 35$, X-ray binaries in the LMC equals $N_{LMXB}:N_{HMXB}\\sim 1:7$ \\citep[][]{gilfanov04}, i.e. is nearly opposite  to that observed in the Milky Way \\citep{grimm02}. The absolute number of HMXBs in the LMC is expected to be $\\sim 1/4$ of the number observed in our Galaxy, corresponding to the ratio of star formation rates in these two galaxies.     As has been shown by \\citet{grimm03}, the X-ray luminosity function (XLF) of HMXBs obeys, to the first approximation, a universal power law distribution with a differential slope of $\\approx 1.6$, whose normalization is proportional to the star formation rate of the host galaxy. The validity of this universal HMXB XLF has been established in a broad range of the star formation rates and regimes and in the luminosity range   $\\log(L_X)\\ga 35.5-36$. Based on the ASCA observations of the Small Magellanic Cloud and on the behavior of the integrated X-ray luminosity of distant galaxies located at redshifts $z\\sim 0.3-1.3$ observed by Chandra in the Hubble Deep Field North, \\citet{grimm03} tentatively suggested that the HMXB XLF is not dramatically affected by metallicity variations.     Study of the population of high mass X-ray binaries in the LMC is of importance for several reasons: \\begin{enumerate} \\item The Magellanic Clouds are known to have a significant under-abundance of metals \\citep{mc_book97}. For the interstellar medium and young stellar population in the LMC the metal abundance is approximately $\\sim 1/3$ of the solar value \\citep{mc_book97, garnett1999, korn2002}. The effects of the metallicity variations on the population of X-ray binaries are poorly understood. Study of the LMC, a galaxy with a relatively well-known chemical composition gives a unique opportunity to investigate such effects observationally. \\item Owing to the proximity of the LMC, the weakest sources become reachable with a moderate observing time. Indeed, the sensitivity of a typical Chandra or XMM-Newton observation, $\\sim 10^{-14}$ erg/s/cm$^2$, corresponds to the luminosity of $\\sim 3\\cdot 10^{33}$ erg/s at the LMC distance. This opens the possibility to study the low luminosity end of the HMXB XLF, extending by $\\sim 2-3$ orders of magnitude towards low luminosities the luminosity range studied by \\citet{grimm03}. Among other applications, this will allow study of the impact of propeller effect \\citep{propeller} on the luminosity distribution of HMXBs. \\item The proximity and moderate inclination angle of the Magellanic Clouds (especially that of the LMC) allows one to construct Herzsprung-Russel diagrams and accurately determine ages,  star formation history and initial mass function of various components of the stellar population. In particular, large aggregates of young and coeval stellar populations of ages ranging from $\\la 1-2$ Myr to $\\ga 50-100$ Myr were discovered in the LMC. The most extreme examples of these are the LMC 4 supergiant shell \\citep{braun97} and R136 (\\object{HD 38268}) stellar cluster near the center of the nebula 30 Dor \\citep{massey98}. Combined with HMXB number counts in the optically studied regions, these results provide a unique possibility to directly determine the efficiency of HMXB formation as a function of time elapsed since the star formation event.  \\end{enumerate}  The proximity of the LMC, on the other hand, creates several difficulties in its studies. Due to its large angular extent on the sky, $\\sim 10\\degr\\times 10\\degr$, multiple mosaiced observations are required to achieve a meaningful coverage of the galaxy. The low surface density of X-ray binaries results in a large fraction of foreground stars and background AGNs among the detected sources, thus creating the problem of separating true galactic members from interlopers. This was one of the main obstacles in the earlier studies of  X-ray binaries in the LMC with moderate angular resolution, insufficient for reliable identification of X-ray sources with optical catalogs. With the advent of Chandra and XMM-Newton observatories, the latter problem is to a large extent alleviated.   \\begin{table*} \\renewcommand{\\arraystretch}{1.2} \\caption{List of XMM-Newton observations used for analysis} \\label{tb:pnt} \\begin{centering} \\begin{tabular}{lcccccc} \\hline Obs. ID   & Target & R.A.  & Dec.  & Instrument & Exposure & SFR$^{(1)}$\\\\ &        & J2000 & J2000 &            & ksec & $M_{\\sun}$/yr\\\\ \\hline 0113000501& 0519-69.0      &05 19 23 & -69 00 52 & MOS1+MOS2  & 25+47&$4.13\\cdot10^{-3}$\\\\ 0111130201& 2E 0509.5--6734&05 09 16 & -67 31 41 & MOS1+MOS2 & 34+34&$7.77\\cdot10^{-4}$\\\\ 0062340101& 2E 0453.8--6834&04 53 41 & -68 31 07 & MOS1+MOS2 & 16+16&$1.49\\cdot10^{-3}$\\\\ 0111130301& 2E 0525.3--6601&05 25 13 & -65 59 38 & MOS1+MOS2 & 9.5+9.5&$1.83\\cdot10^{-3}$\\\\ 0134520701& AB Dor         &05 28 35 & -65 28 30 & MOS1+MOS2 & 48+48&$2.92\\cdot10^{-4}$\\\\ 0123510101& CAL 83         &05 43 33 & -68 24 07 & MOS1+MOS2 & 10+10&$1.37\\cdot10^{-3}$\\\\ 0104060101& LMC Deep field &05 31 17 & -65 55 58 & MOS1+MOS2 & 38+38&$6.35\\cdot10^{-4}$\\\\ 0109990201& DEM L71        &05 05 49 & -67 54 11 & MOS1+MOS2 & 22+23&$1.30\\cdot10^{-3}$\\\\ 0109990101& LHA 120--N63A  &05 35 51 & -66 00 32 & MOS1+MOS2 & 9.6+9.6&$1.38\\cdot10^{-3}$\\\\ 0094410101& LMC-2(f1)      &05 42 35 & -69 03 20 & PN & 8.8&$9.74\\cdot10^{-3}$\\\\ 0094410201& LMC-2(f2)      &05 42 58 & -69 28 22 & PN & 8.8&$5.51\\cdot10^{-3}$\\\\ 0094410401& LMC-2(f4)      &05 46 48 & -69 33 39 & MOS1+MOS2 & 5.2+5.2&$4.60\\cdot10^{-3}$\\\\ 0126500101& LMC X-3        &05 38 44 & -64 06 11 & MOS1+MOS2  & 20+20&$7.82\\cdot10^{-5}$\\\\ 0113000301& N103B          &05 08 41 & -68 43 53 & MOS2 & 23ks&$4.52\\cdot10^{-3}$\\\\ 0089210701& N120           &05 18 37 & -69 37 56 & MOS1+MOS2 & 36+36&$5.76\\cdot10^{-3}$\\\\ 0089210901& N206           &05 31 51 & -70 58 39 & MOS1+MOS2 & 24+24&$3.35\\cdot10^{-3}$\\\\ 0071940101& N51D           &05 26 05 & -67 27 21 & MOS1+MOS2 & 32+32&$3.37\\cdot10^{-3}$\\\\ 0127720201& Nova LMC 2000  &05 24 41 & -70 14 00 & MOS1+MOS2 & 22+22&$1.34\\cdot10^{-3}$\\\\ 0113000401& PSR 0540--69.3 &05 40 02 & -69 21 19 & MOS2 & 35&$2.29\\cdot10^{-2}$\\\\ 0113020201& PSR J0537--6909&05 37 57 & -69 08 52 & MOS1 & 37&$5.08\\cdot10^{-2}$\\\\ 0008020101& RXJ0439.8--6809&04 39 57 & -68 07 27 & MOS1+MOS2 & 15+15&$1.83\\cdot10^{-4}$\\\\ 0089210601& SNR0450-709    &04 50 26 & -70 48 48 & MOS1+MOS2 & 56+57&$6.19\\cdot10^{-4}$\\\\ 0137160201& YYMen          &04 58 12 & -75 15 01 & MOS1+MOS2 & 87+87&$3.50\\cdot10^{-4}$\\\\ \\hline \\end{tabular}\\\\ \\end{centering} 1 -- star formation rate within the XMM-Newton FOV determined form the FIR flux as described in sections \\ref{sec:sfr} and \\ref{sec:hmxb}. \\end{table*}    In the present paper we study the population of X-ray binaries in the LMC based on archival XMM-Newton data. The distance modulus of the LMC is  $m-M=18.45\\pm0.10$ \\citep{mc_book97}, corresponding to the distance of $D\\approx 50$ kpc. The interstellar reddening varies across the LMC, with typical values in the range $E_{B-V}\\approx 0.05-0.2$ \\citep{mc_book97}. A value of $E_{B-V}=0.075$ corresponding to the direction towards the nominal center of the LMC is used throughout the paper.   The paper is structured as follows. XMM-Newton observations and data analysis are described  in section ~\\ref{sec:data}. In section ~\\ref{sec:nature} we discuss the nature of detected X-ray sources. In section ~\\ref{sec:ident} we describe the HMXB search procedure and its results. Resulting HMXB luminosity function and the  CXB log(N)--log(S) are presented in section ~\\ref{sec:hmxbcxb}. The impact of the propeller effect on the HMXB luminosity function is considered in section ~\\ref{sec:propeller}. In section ~\\ref{sec:age_effects} we discuss observed spatial non-uniformity of the $N_{HMXB}/SFR$ ratio and its dependence on the age of the underlying stellar population.  Our results are summarized in section \\ref{sec:summary}.    \\begin{figure} \\includegraphics[width=0.5\\textwidth,clip=true]{shtykfig1.ps} \\caption{Far-infrared (IRAS, $100\\mu$) map of the $10\\degr\\times 10\\degr$ region around the Large Magellanic Cloud. Circles show fields of view of the XMM--Newton observations used for the analysis. } \\label{fig:fir_map} \\end{figure} ",
        "conclusions": "\\label{sec:summary}  Based on the archival data of XMM--Newton observations, we studied the population of compact sources in the field of the LMC. The total area of the survey is $\\approx 3.8$ sq.degr. with a limiting sensitivity of $\\sim 10^{-14}$ erg/s/cm$^2$ (Fig.\\ref{fig:area}), corresponding to the luminosity of $\\sim 3\\cdot 10^{33}$ erg/s at the LMC distance. \\begin{enumerate} \\item Out of the 460 compact sources detected in the 2--8 keV energy band the vast majority, $\\ga 94\\%$, are CXB sources, observed through the LMC (section \\ref{sec:cxb}, Figs.\\ref{fig:lf0},\\ref{fig:lfxx}). \\item Based on the stellar mass and the star formation rate of the LMC we demonstrate that the majority, $\\ga 75-90\\%$, of the intrinsic LMC sources detected in the 2--8 keV band are high mass X-ray binaries (section \\ref{sec:nature}). \\item The proximity of the LMC and the adequate angular resolution of XMM--Newton make it possible to reliably filter out the sources whose properties are inconsistent with being an HMXB. Based on the optical and infrared magnitudes and colors of the optical counterparts of the X-ray sources we identify 9 likely HMXB candidates (6 of which were previously known HMXBs or HMXB candidates) and 19 sources of uncertain nature (section \\ref{sec:ident}, Table \\ref{tb:hmxbcand}). The remaining $\\sim 440$ sources, with a few exceptions, are background objects, constituting the resolved part of CXB. Their flux distribution is consistent with other determinations of the  CXB log(N)--log(S) (section \\ref{sec:cxbsrc}, Fig.\\ref{fig:cxb}). \\item With these results we constrain lower and upper bounds of the luminosity distribution of HMXBs in the observed part of the LMC. We compare these with the extrapolation towards low luminosities of the universal luminosity function of HMXBs, derived by \\citet{grimm03}. At the high luminosity end, the observed distribution is consistent with the extrapolation of the universal XLF. The number of bright, $L_X\\ge 10^{38}$ erg/s, ``historical'' HMXBs in the LMC also agrees with the value predicted from its global star formation rate. At lower luminosities, a deficit of sources is observed --  the luminosity distribution seems to be somewhat flatter than the universal XLF (section \\ref{sec:hmxbcand}, Fig.\\ref{fig:lfhmxb}). \\item We consider the impact of the ``propeller effect'' on the luminosity distribution of HMXBs (section \\ref{sec:propeller}, Fig.\\ref{fig:proplum}) and demonstrate that it can explain qualitatively the observed flattening of the XLF at low luminosities (Fig.\\ref{fig:lfhmxb}). We note, however, that due to the relatively small number of HMXB candidates in the observed part of the LMC and the lack of distinct signatures of the ``propeller effect'' in the $\\lg(L_X)\\ga 33$ luminosity domain (Fig.\\ref{fig:proplum}), it is premature to draw a definite conclusion about the relevance of the ``propeller regime'' to the accretion onto a strongly magnetized neutron star. \\item We found significant field-to-field variations in the number of HMXBs, which appear to be uncorrelated with the star formation rates inferred by FIR and $H_\\alpha$ emission (Fig.\\ref{fig:halpha_src}). We suggest that these variations are related to different ages of the underlying stellar population in different XMM-Newton fields. Using the existence of large coeval stellar aggregates in the LMC we constrain the number of HMXBs per unit stellar mass as a function of time elapsed since the star formation event (section \\ref{sec:age_effects}, Fig.\\ref{fig:eta_hmxb}). Based on a simple ad hoc model, we obtain the theoretical dependence $\\eta_{HMXB}(\\tau)$ and show that its normalization, as constrained by the LMC data, is consistent with the calibration of the $N_{HMXB}$--SFR relation, derived by \\citet{grimm03}. \\end{enumerate}"
    },
    "0404/hep-ph0404015_arXiv.txt": {
        "abstract": "If the sterile neutrino mass matrix in an otherwise conventional seesaw model has a rank less than the number of flavors, it is possible to produce pseudo-Dirac neutrinos. In a two-flavor, sterile rank 1 case, we demonstrate analytic conditions for large active mixing induced by the existence of (and coupling to) the sterile neutrino components. For the three-flavor, rank 1 case, ``3+2'' scenarios with large mixing also devolve naturally as we show by numerical examples.  We observe that, in this approach, small mass differences can develop naturally without any requirement that masses themselves are small.  Additionally, we show that significant three channel mixing and limited experimental resolution can combine to produce extracted two channel mixing parameters at variance with the actual values. ",
        "introduction": "Conventional wisdom holds that neutrinos ought to be Majorana particles with very small masses, due to the action of a ``seesaw'' mechanism\\cite{see-saw}, which is built on the concept of quark-lepton symmetry\\cite{GSMcK}.  Alternatively, there have been theoretical suggestions regarding pseudo-Dirac neutrinos in the past\\cite{wolfm2,oldref}, and again more recently\\cite{GSMcK,CHM,othr}, i.e., that neutrinos may well be Majorana particles occuring in nearly degenerate pairs. These can occur within the framework of the so-called ``singular'' see-saw where the rank of the mass matrix describing the (presumed to be) heavy neutrinos, which have no interactions (often referred to as ``sterile'' neutrinos) in the standard model~(SM), is less than maximal.  Recent results from Kamiokande\\cite{superK} on atmospheric neutrinos, from Sudbury\\cite{SNO} on solar neutrinos, and from KamLand\\cite{KamL} on long baseline reactor neutrinos, appear to require oscillations between nearly maximally mixed (active neutrino) mass eigenstates.  Each of these analyses, however, argues that this mixing cannot be dominantly to sterile states such as are found in pseudo-Dirac pairs. On the other hand, the concatenation of the data from these experiments with that from LSND\\cite{LSND} and other short baseline data does not appear to fit into a theoretical structure which only includes mixing among three active Majorana neutrinos.  Many have therefore been motivated to consider the effects of additional (sterile, Majorana) neutrino states, the existence of which is accepted in the conventional ``see-saw'' extension of the SM, although there the actual states are generally precluded from appearing directly in experiments by an assumption that the masses of the sterile states are very large.  We investigate here how small flavor mixing effects in the sterile sector can lead to large mixing among active neutrinos in the presence of a singular see-saw. (In Ref.\\cite{CHM}, large mixing  was achieved by means of a mass hierarchy in the Dirac mass sector.) Paralleling a convention in the quark sector, we assume the mass and flavor bases for the charged leptons are simultaneously diagonal, so that all flavor violations and oscillation phenomena are described as arising from the neutrino mixing angles alone.  It should be noted that there is no accepted principle that specifies the flavor space structure of the mass matrix assumed for the sterile sector. Some early discussions\\cite{see-saw,wolfm2} implicitly assume that a mass term in the sterile sector should be proportional to the unit matrix.  This has the pleasant prospect, in terms of the initial argument for the see-saw, that all active neutrino flavors have small masses on the scale of other fermions.  However, since there is no obvious requirement that Dirac masses in the neutral lepton sector are the same as Dirac masses in any other fermionic sector, this result is not compelling. Indeed, Goldhaber has argued for a view of family structure and self-energy based masses that naturally produces small neutrino masses\\cite{Maurice}. We discuss here a more conventional possibility which arises from a minimal modification of the standard see-saw, namely that the rank of the mass matrix for the sterile sector is less than the number of flavors. Note that this does not conflict with quark-lepton symmetry which applies only to the number and character of states.  In this paper, which is an extension of reference\\cite{hep}, we shall concentrate on the case of a rank $1$ sterile matrix, relegating the rank $2$ case to some remarks at the end. (The analysis of short baseline data by Sorel, Conrad and Shaevitz\\cite{Sorel} suggests that the rank $2$ case may not actually occur in Nature.) We note in passing that some Grand Unified Theories include more than $3$ fermions that are neutral under all of the interactions in the SM; a $4\\times4$ or larger, rank $1$ sterile mass matrix could lead to $3$ pseudo-Dirac pairs of neutrinos involving all of the active neutrinos of the SM.  Concentrating on a $3$-dimensional sterile space, we consider rank $1$ to be a natural case because whatever spontaneous symmetry breaking produces mass in that flavor space necessarily defines a specific direction. Before including the effects of the sterile mass, we assume three non-degenerate Dirac neutrinos, with Dirac masses, $m_1~<~m_2~<~m_3$, (although this is not essential,) which are each constructed from one Weyl spinor which is active under the $SU(2)_W$ of the SM and one Weyl spinor which is sterile under that interaction. (Being neutrinos, both Weyl fields have no interactions under the $SU(3)_C$ or the $U(1)$ of the SM.) There is then an MNS matrix\\cite{MNS} which relates these Dirac mass eigenstates to the flavor eigenstates in a manner completely parallel to that of the CKM matrix\\cite{ckm} for quarks.  Note, however, that these matrix elements are not the ones extracted directly from experiment, as the mass matrix in the sterile sector induces additional mixing.  We next use the Dirac mass ($m_{D}$) eigenstates to define bases in both the $3$-dimensional active flavor space and the $3$-dimensional sterile flavor space\\cite{fn1}. Following the spirit of the original see-saw, we exclude any initial Majorana mass term in the active space. If the Majorana mass matrix in the sterile space were to vanish also, the three flavors of Dirac neutrinos would be a mixture of (Dirac) mass eigenstates in a structure entirely parallel to that of the quarks.  A rank $1$ sterile mass matrix may be represented as a vector of length $M$ oriented in some direction in the $3$-dimensional sterile space. If that vector lies along one of the axes, then the Dirac neutrino that would have been formed from it and its active neutrino partner partake of the usual see-saw structure (one nearly sterile Majorana neutrino with mass approximately $M$ and one nearly purely active neutrino with mass approximately $m_D^2/M$) and the other two mass eigenstates remain Dirac neutrinos. If that vector lies in a plane perpendicular to one axis, the eigenstate associated with that axis will remain a pure Dirac neutrino, and the other two pairs of states form one pseudo-Dirac pair and a pair displaying the usual see-saw structure. Both of these pairs are mixtures of the $4$ Weyl fields associated with the two mixing Dirac neutrinos. In general, the structure consists of $2$ pseudo-Dirac pairs and one see-saw pair, all mixed.  As we implied above, the very large mixing required by the atmospheric neutrino measurements could have been taken to be evidence for a scheme involving pseudo-Dirac neutrinos. (This, after all, follows Pontecorvo's initial suggestion\\cite{BP}.) However, pure mixing into the sterile sector is now strongly disfavored\\cite{nsm}.  It is evident from the discussion above that there is a region of parameter space (directions of the vector) in which the two pseudo-Dirac pairs are very nearly degenerate, giving rise to the possibility of strong mixing in the active sector coupled with strong mixing into the sterile sector.  We explore this point here.  The organization of the remainder of the paper is as follows: In the next section we discuss a two flavor, $4\\times 4$ neutrino mass matrix analytically. In Sec.\\ref{genmass}, we present the general $6\\times 6$ mass matrix and discuss the parameterization of the sterile mass matrix and various limiting cases.  We show the spectrum for a general case.  In Sec.\\ref{2pds}, we apply our analysis to the case where the plane in question is perpendicular to the axis for the middle value ($m_2$) Dirac mass eigenstate, raising the possibility of near degeneracy between pseudo-Dirac pairs. Moving away from that plane produces large  mixing amongst the members of those pseudo-Dirac pairs. In Sec.\\ref{example}, we show an example of the oscillation patterns that are produced and how limited experimental resolution can lead to errors in the extraction of physical parameters if the data analysis assumes only two channel mixing.  Finally, we remark on the structures expected for a rank $2$ sterile matrix and then reiterate our conclusions. ",
        "conclusions": "We have considered here the effects on neutrinos in the SM of the recurrently successful and conventional constraint of quark-lepton symmetry, namely, the existence of six independent Weyl spinor fields of neutrinos, three corresponding to active and three corresponding to sterile neutrinos. In the now venerable see-saw approach, the latter three effectively disappear from the excitation spectrum, leaving small Majorana masses for the active states as a residuum. We have examined the effect on this system of a rank less than three character of the $3 \\times 3$ mass matrix in the sterile sector and studied the rank $1$ case, in particular.  In the rank $1$ case on which we have focused, we find that the neutrino fields naturally form into two pseudo-Dirac pairs, leaving only one almost pure Majorana active neutrino and one conventionally very heavy sterile Majorana neutrino. More importantly, we also find a naturally strong mixing between the active and sterile parts of the two pseudo-Dirac pairs.  Further, we find that this can easily affect the mixing between active neutrinos even if it is otherwise small. That is, even if the Dirac mass matrix induced mixing analogous to what is known to occur in the quark sector is small or absent, mixing between active neutrinos can develop with large values. In a two flavor case, we demonstrated analytically that this strong mixing can develop over a wide range of parameters.  We have chosen a limited relative value of the sterile neutrino mass scale, $M$, that allows for easy discernment of the nature of the effects. It should be noted, however, that the primary effect of increasing $M$ is to decrease the separation of the two members of each pseudo-Dirac pair while producing the usual see-saw behavior for the remaining pair. Thus, tiny differences in mass between masses that are not especially small themselves are, in the usual sense of the term, natural in this approach. This is contrary to the general expectation that the small mass differences responsible for the observed neutrino oscillation phenomena presage small absolute masses for all of the neutrinos. Furthermore, increasing the value of $M$ without altering the Dirac masses retains the features and scales of the oscillations essentially unchanged; the only significant change is that the appearance of the sterile components is delayed to even greater values of $L/E$.  The features described above are most easily discerned in the case when the Dirac mass terms for the neutrinos are well separated in value. It remains conceivable that, if their differences are small for some other reason, then the splitting between the pseudo-Dirac pairs may be larger than that between flavors. In this case, it is still true that large flavor mixing is naturally induced.  Finally, we presented a specific model which raises concerns about the conclusions drawn from analyses of neutrino oscillation experiments in terms of two channel mixing: Such analyses may be misleading regarding the true values of physical parameters. After the completion of this work, we learned of papers\\cite{ernst} which have independently suggested that such a concern may well be justified.  \\vspace*{-0.1in}"
    },
    "0404/astro-ph0404466_arXiv.txt": {
        "abstract": "We present some preliminary results from an on-going study of the evolution of stellar populations in rich clusters of galaxies. Our baseline sample contains core line-strength measurements from 183 galaxies with $b_{\\rm J} \\le 19.5$ from four clusters with $\\bar{z}\\sim0.04$, against which observations of higher-redshift clusters can be compared.  Using predictions from stellar population models to compare with our measured line strengths we can derive \\textit{relative} luminosity-weighted mean ages and metallicities for the stellar populations in each of our clusters. It must be stressed that these ages and metallicities are only accurate when used in an relative sense as the stellar population models, due to differing abundance ratios used in the models compared to those observed in elliptical galaxies, provide inaccurate absolute ages and metallicities. We also investigate the Mg\\textit{b}$^\\prime$-$\\sigma_0$ and H$\\beta_G^\\prime$-$\\sigma_0$ scaling relations. We find that Mg\\textit{b}$^\\prime$ is correlated with $\\sigma_0$, the likely explanation being that larger galaxies are better at retaining their heavier elements due to their larger potentials. H$\\beta_G^\\prime$, on the other hand, we find to be anti-correlated with $\\sigma_0$. This result implies that the stellar populations in larger galaxies are older than in smaller galaxies. ",
        "introduction": "\\label{sec:intro}  The formation and evolution of clusters of galaxies and of their early-type galaxy population is a topic of active debate. Studying the integrated light from stellar populations can help to discriminate between models of giant elliptical galaxy formation: the monolithic collapse model (\\cite[Eggen et al., 1962]{eggen62}) in which giant elliptical galaxies form rapidly in a process that can essentially be considered a single collapse and the hierarchical merging model (\\cite{searle78}) in which giant elliptical galaxies are built up over a long timescale by distinct merger events. If it was found that early-types were largely coeval in clusters then this would favour the monolithic collapse model, whereas a large spread in ages would favour hierarchical merging.  Increasing a galaxy's age by a factor of two while decreasing it's metallicity by a factor of three results in a spectrum almost indistinguishable from the original (\\cite{worthey94}). This age/metallicity degeneracy can be partially broken by comparing age-sensitive spectroscopic indices with metallicity-sensitive spectroscopic indices. By plotting two such indices against each other and comparing the results to stellar population models (e.g. \\cite{worthey94}; \\cite{thomas03}) we can determine the mean relative luminosity-weighted age and metallicity of a given stellar population.  Previous work in this area has concentrated on individual clusters such as Coma (\\cite{jorgensen99}; \\cite{moore02}) and Fornax (\\cite{kuntschner00}) with sometimes conflicting results coming from the same cluster, e.g. \\cite[Gonzalez et al. (1993)]{gonzalez93} and \\cite[Moore et al. (2002)]{moore02}. Coma and Fornax represent the extremes of the cluster richness scale and this project will study clusters that fill in this range. We will study, in detail, the stellar populations in several low-redshift ($z\\sim0.04$) clusters from the core out into the surrounding structures. These results will then be used as a comparison for observations of a range of higher-redshift clusters ($0.3\\leq z \\leq 0.55$), allowing us to draw conclusions about the formation and evolution of the stellar populations in rich clusters.  Line-strength measurements, redshifts and velocity dispersions were obtained for 183 galaxies from four clusters over three nights in April 2002 with the 2dF system on the 3.9m AAT. The sample was selected from galaxies with $b_{\\rm J}\\le19.5$, within a $2^{\\rm o}$ diameter FOV centred on each of the clusters and within a $3\\sigma$ redshift range centred on the cluster redshift. The 300B grating ($\\sim9$\\AA\\, FWHM) was used to obtain a wide spectral range covering as many indices as possible, while the 1200V grating ($\\sim2$\\AA\\, FWHM) was used to obtain higher resolution spectra yielding precise velocity dispersions.  \\begin{figure} \\epsfig{file=figure1.eps,width=\\textwidth} \\caption{H$\\beta_G$ equivalent width versus [MgFe] equivalent width. \\cite[Thomas et al. (2003)]{thomas03} stellar population model grids have been overlaid on data from the Coma cluster (\\textit{left}) and Abell 1139 (\\textit{right}).} \\label{fig:thomas} \\end{figure} ",
        "conclusions": ""
    },
    "0404/astro-ph0404185_arXiv.txt": {
        "abstract": "s{ We present the clustering properties of hard (2-8\\,keV) X-ray selected sources detected in a wide field ($\\approx \\rm 2\\,deg^{2}$) shallow [$f_X(\\rm 2-8\\,keV)\\approx 10^{-14}\\rm \\, erg \\, cm^{-2} \\, s^{-1}$] and contiguous  XMM-{\\it Newton} survey. We perform an angular correlation function analysis using a total of  171 sources to the above flux limit.  We detect a $\\sim 4\\sigma$ correlation signal out to 300\\,arcsec with $w(\\theta < 300^{''})\\simeq  0.13 \\pm 0.03$. Modeling the two point correlation function as a  power law of the form $w(\\theta)=(\\theta_{\\circ}/\\theta)^{\\gamma-1}$ we find: $\\theta_{\\circ}=48.9^{+15.8}_{-24.5}$ arcsec and $\\gamma=2.2\\pm {0.30}$. Fixing the correlation function slope to $\\gamma=1.8$ we obtain  $\\theta_{\\circ}=22.2^{+9.4}_{-8.6}$\\,arcsec. Using Limber's integral equation and a variety of possible luminosity functions of the hard X-ray population, we find a relatively large correlation length, ranging from $r_{\\circ}\\sim 9$ to 19 $h^{-1}$ Mpc (for $\\gamma=1.8$ and the {\\em concordance} cosmological model), with this range reflecting also different evolutionary models for the source luminosities and clustering characteristics.} ",
        "introduction": "The overall knowledge of the AGN clustering using X-ray data comes mostly from the soft X-ray band (Boyle \\& Mo 1993; Vikhlinin \\& Forman 1995; Carrera et al. 1998; Akylas, Georgantopoulos, Plionis, 2000; Mullis 2002), which is however biased against absorbed AGNs. Recently, Yang et al. (2003) performing a counts-in cells analysis of a deep ($f_{2-8 keV} \\sim 3 \\times 10^{-15}$ erg s$^{-1}$ cm$^{-2}$) {\\em Chandra} survey in the Lockman Hole North-West region, found that the hard band sources are highly clustered with $\\sim$ 60$\\%$ of them being distributed in overdense regions.  In this study we use a hard (2-8\\,keV) X-ray selected sample, compiled from a shallow (2-10\\,ksec per pointing) XMM-{\\em Newton} survey near the North and South Galactic Pole regions. A total of 18 pointings were observed out of which only 5 were discarded due to elevated particle background at the time of  the observation. A full description of the data  reduction, source detection and flux estimation are presented by Georgakakis et al. (2003). Here we just note that it comprises of 171 sources above the $5\\sigma$ detection threshold to the limiting flux of $f_X(\\rm 2-8\\,keV) \\approx 10^{-14}\\, erg \\, s^{-1}\\, cm^{-2}$. ",
        "conclusions": ""
    },
    "0404/astro-ph0404520_arXiv.txt": {
        "abstract": "The transport properties of dense stellar electron-proton plasma is studied following an exact relativistic formalism in presence of strong quantizing magnetic field. The variation of transport coefficients with magnetic fields are found to be insensitive for the field strengths $\\leq 10^{17}$G, beyond which all of them abruptly go to zero. As a consequence, the electron-proton plasma behaves like a superfluid insulator in presence of ultra-strong magnetic field. ",
        "introduction": "The study of the effect of strong magnetic field on dense stellar plasma is one of the oldest branches of physics. It has gotten a new life after the discovery of a few magnetars- the strange stellar objects, with unusually high surface magnetic fields \\cite{R1,R2,R3,R4}. These stellar objects are believed to be strongly magnetized young neutron stars. The surface magnetic fields are observed to be $\\sim 10^{15}$G. Then it is quite possible that the field strength at the core region may go up to $10^{18}$G.  The exact source of this strong magnetic field is of course yet to be known. These objects are also supposed to be the possible sources of anomalous X-ray and soft gamma emissions (AXP and SGR). If the magnetic fields are really so strong, in particular at the core region, they must affect significantly most of the important physical properties of such stellar objects and the physical processes, e.g., weak and electromagnetic interactions taking place at the core region. Which means, the presence of strong quantizing magnetic field at the core region should modify, both qualitatively and quantitatively the equation of state of dense neutron star matter, and as a consequence the gross-properties of neutron stars \\cite{R5,R6,R7,R8}, e.g., mass-radius relation, moment of inertia, rotational frequency etc. should also change significantly. In the case of compact neutron stars, the phase transition from neutron matter to quark matter at the core region, if any, will also be affected by strong quantizing magnetic field.  It has been shown that a first order phase transition initiated by the nucleation of quark matter droplets is absolutely forbidden if the magnetic field $\\sim 10^{15}$G at the core region \\cite{R9,R10}. However a second order phase transition is allowed, provided the magnetic field strength $<10^{20}$G. This is of course too high to achieve at the core region.  The elementary processes, in particular, the weak and the electromagnetic processes taking place at the core region of a neutron star are strongly affected by such ultra-strong magnetic field \\cite{R11,R12}. Since the cooling of neutron stars are mainly controlled by neutrino/anti-neutrino emissions, the presence of strong quantizing magnetic field should affect the thermal history of strongly magnetized neutron stars. Further, the electrical conductivity of neutron star matter which directly controls the evolution of neutron star magnetic field will also change significantly.  In another kind of work, the stability of such strongly magnetized rotating objects are studied. It has been observed from the detailed general relativistic calculation that there are possibility of some form of geometrical deformation of these objects from their usual spherical shapes \\cite{R13,R14,R15}. In the extreme case such objects may either become black strings or black disks. In the non-extreme case, however, it is also possible to detect gravity waves from these deformed rotating objects.  In a recent study on microscopic model of dense neutron star matter, we have observed that if most of the electrons occupy the zeroth Landau level, with spin anti-parallel to the direction of magnetic field and very few of them are with spin along the direction of magnetic field and Landau quantum number $>0$, then either such strongly magnetized system can not exist or such a strong magnetic field is just impossible at the core region of a neutron star \\cite{R16}.  Motivated by the problems as mentioned in preceding two sections, in this paper we shall study the effect of strong quantizing magnetic field on the transport coefficients of dense stellar electron-proton plasma. We shall follow an exact formalism \\cite{R17} which is applicable for a wide range of magnetic field strengths and obtain the transport coefficients from the relativistic version of Boltzmann kinetic equation by linearizing the distribution function and using relaxation time approximation. We shall obtain the relaxation time from the rates of standard electromagnetic processes taking place inside electron-proton plasma and make necessary modification in the rate calculation due to the presence of strong quantizing magnetic field.   We have noticed that the electrical conductivity of the medium becomes extremely small in presence of ultra-strong magnetic field ($\\geq 10^{17}$G). The magnetic field at the core region of a magnetar must therefore decay very rapidly (time scale $\\sim$ a few mins.) and becomes moderate or low enough. As a consequence there will be in principle no problem on the existence of magnetars (with very low or moderate core magnetic field). The formalism we have developed to obtain rates of electromagnetic processes or the relaxation time is also applicable to evaluate neutrino emissivity and mean free path in presence of strong quantizing magnetic field.  In presence of strong quantizing magnetic field, since the momentum component in the transverse plane with respect to the external magnetic field gets quantized, whereas the component along the field direction varies continuously (from $-\\infty$ to $+\\infty$), the momentum space volume element becomes \\begin{equation} \\frac{d^3p}{(2\\pi)^3}=\\frac{dp_xdp_ydp_z}{(2\\pi)^3} =\\frac{eB_m}{4\\pi^2}\\sum_{\\nu=0}^\\infty(2-\\delta_{\\nu 0}) dp_z \\end{equation} (we have assumed $\\hbar=c=k_b=1$) where we have chosen the gauge $A^\\mu\\equiv (0,0,xB_m,0)$, so that the constant magnetic field $B_m$ is along z-direction. We have considered the simplest possible picture of neutron star matter with $n-p-e$ out of thermodynamic equilibrium and the neutrinos are assumed to be non-degenerate.  The baryonic components are interacting via $\\sigma-\\omega-\\rho$ meson exchange type mean field and the electrons are assumed to be freely moving particles.  In this article, we shall first calculate the transport coefficients of electron gas. Then it is very easy to obtain the transport coefficients for proton matter just by replacing mass, chemical potential etc. of electrons by protons and taking into account the proper modification in presence of $\\sigma-\\omega-\\rho$ meson exchange type mean field \\cite{R6}. Spinor solutions for electrons in presence of strong quantizing magnetic fields are then given by \\begin{equation} \\Psi^{(\\uparrow)(e)}=\\frac{1}{\\sqrt {L_yL_z}} \\frac{\\exp(-i\\varepsilon_{\\nu}^{(e)}t+ip_y y+ip_z z)} {[2\\varepsilon_{\\nu}^{(e)}(\\varepsilon_{\\nu}^{(e)}+m_e)]^{1/2}} \\times\\left( \\begin{array}{c}(\\varepsilon_{\\nu}^{(e)}+m_e) I_{\\nu;p_y}(x)\\\\ 0\\\\p_z I_{\\nu;p_y}(x)\\\\-i(2\\nu e B_m)^{1/2}I_{\\nu-1;p_y}(x) \\end{array} \\right) \\end{equation} and \\begin{equation} \\Psi^{(\\downarrow)(e)}=\\frac{1}{\\sqrt {L_yL_z}} \\frac{\\exp(-i\\varepsilon_{\\nu}^{(e)}t+ip_y y+ip_z z)} {[2\\varepsilon_{\\nu}^{(e)}(\\varepsilon_{\\nu}^{(e)}+m_e)]^{1/2}} \\times\\left( \\begin{array}{c} 0\\\\(\\varepsilon_{\\nu}^{(e)}+m_e)I_{\\nu-1;p_y}(x)\\\\ i(2\\nu e B_m)^{1/2}I_{\\nu;p_y}(x)\\\\-p_z I_{\\nu-1;p_y}(x) \\end{array} \\right) \\end{equation} where $\\uparrow$ and $\\downarrow$ represent up and down spin states respectively, \\begin{equation} I_{\\nu;p_y}(x)=\\left( \\frac{e B_m}{\\pi}\\right)^{1/4} \\frac{1}{\\sqrt{\\nu!}~2^{\\nu/2}} \\times\\exp\\left[ -\\frac{1}{2}e B_m \\left( x- \\frac{p_y}{e B_m}\\right)^2 \\right] H_{\\nu} \\left[ \\sqrt{e B_m}\\left( x- \\frac{p_y}{e B_m} \\right) \\right], \\end{equation} with $H_{\\nu}$ the Hermite polynomial of order $\\nu$ and \\begin{equation} \\varepsilon_\\nu^e=(p_z^2+m_e^2+2e\\nu B_m)^{1/2} \\end{equation} the energy eigen value with $\\nu=0,1,2,.....,$ the Landau quantum numbers. For neutron we consider the usual spinor solutions. Since the temperature of the system is $\\ll$ electron chemical potential, we have not considered the negative energy spinor solutions, In the energy eigen value $p_\\perp= (2\\nu eB_m)^{1/2}$ is the transverse component of electron momentum. To overcome the serious problem on the mechanical stability and hence the existence of magnetars, we have studied the variation of transport coefficients, in particular the electrical conductivity of electron gas / proton matter in presence of strong quantizing magnetic field and try to show whether the electrical conductivity which is solely responsible for the evolution of neutron star magnetic field, becomes sufficiently small in presence of ultra strong magnetic field.  The paper is organized in the following manner. In section 2, we have developed the relativistic version of Boltzmann kinetic equation for fermions in presence of strong quantizing magnetic field and obtain the expressions for transport coefficients following de Groot \\cite{R18}. We shall use the relaxation time approximation for the collision term. In section $3$ we shall evaluate the the relaxation time from the rates of standard electromagnetic processes taking place at the core region. We shall incorporate the necessary changes in the rate calculation due to the presence of strong quantizing magnetic field (or the effect of quantized Landau levels). In the last section we shall discuss importance of the results. ",
        "conclusions": "From the figs.(1)-(6) it is quite obvious that the kinetic coefficients are almost independent of magnetic field for moderate strengths $(<10^{17} G)$,  but all of them go to almost zero for magnetic field strength beyond $10^{17}G$. The first conclusion is therefore, that at ultra strong magnetic field, the matter (electron or proton matter) behaves like a superfluid insulator. The mechanism of superfluidity is of course completely different from conventional neutron matter.  Now as we know from classical plasma physics in presence of strong magnetic field, that the charged particles can only travel along the lines of force, motions are almost forbidden across the field, as a consequence $\\sigma_{\\perp\\perp}$ or $\\sigma_{\\perp z}$ vanish for a classical plasma in presence of ultra-strong magnetic field, whereas $\\sigma_{zz}$ remains non-zero even at very high magnetic field strength. On the other hand, in the case of quantum mechanical plasma as the magnetic field become strong enough, the maximum value of Landau quantum number $\\nu_{\\rm{max}}\\rightarrow 0$, which further means the system becomes effectively one dimensional in the ultra-strong magnetic field limit. Then in the collision dominated scenario, since $p_z$ varies continuously from $-\\infty$ to $+\\infty$ and $p_\\perp\\rightarrow 0$, $\\sigma_{zz}\\rightarrow 0$. To elaborate this point a bit- since the current is flowing along the field lines only, then from the symmetric nature of $p_z$, we have \\begin{equation} j_z^{(+)}=-j_z^{(-)}\\Longrightarrow j_z=j_z^{(+)}+j_z^{(-)}=0 \\end{equation} This is specially true for an effectively one dimensional collision dominated quantum plasma. In the case one dimensional flow (free streaming) of charged particles along the field lines, this is not true; $\\sigma_{zz}$ has non-zero finite value.  Since all the three component of electrical conductivity vanish in presence of ultra-relativistic magnetic field, this must affect significantly the Ohmic decay of strong magnetic field. We know that the Ohmic decay time scale is given by \\begin{equation} \\tau_d=\\frac{4\\pi\\sigma l^2}{c^2} \\end{equation} where $l$ is the dimension of the system. Since $\\sigma\\rightarrow 0$ for extremely strong magnetic field, as a consequence, the field should decay quickly. We have noticed that a field of strength $10^{18}$G become $\\sim 10^{15}$G within a minute. This moderate field of course will take longer time to decay further (long decay time scale). Therefore the last conclusion is that stronger the magnetic field, shorter the decay time scale and as a result we do believe that the magnetic field strength at the core region of a magnetar can not be high enough and consequently there will be no problem on the mechanical stability or the existence of magnetars. \\begin{figure} \\psfig{figure=shr.eps,height=0.5\\linewidth} \\caption{ Variation of shear viscosity coefficient with magnetic field $B$. Curve (a): $T=15$MeV and $n_B=3n_0$, Curve (b): $T=15$MeV and $n_B=5n_0$, and Curve (c): $T=30$MeV and $n_B=3n_0$, } \\end{figure} \\begin{figure} \\psfig{figure=blk.eps,height=0.5\\linewidth} \\caption{ Variation of bulk viscosity coefficient with magnetic field $B$. Curve (a): $T=15$MeV and $n_B=3n_0$, Curve (b): $T=15$MeV and $n_B=5n_0$, and Curve (c): $T=30$MeV and $n_B=3n_0$, } \\end{figure} \\begin{figure} \\psfig{figure=con315.eps,height=0.5\\linewidth} \\caption{ Variation of heat conductivity coefficient  with magnetic field $B$. Curve (a): $T=15$MeV and $n_B=3n_0$, Curve (b): $T=15$MeV and $n_B=5n_0$, and Curve (c): $T=30$MeV and $n_B=3n_0$, } \\end{figure} \\begin{figure} \\psfig{figure=szz.eps,height=0.5\\linewidth} \\caption{ Variation of {\\bf{zz}} component of electrical conductivity  with magnetic field $B$. Curve (a): $T=15$MeV and $n_B=3n_0$, Curve (b): $T=15$MeV and $n_B=5n_0$, and Curve (c): $T=30$MeV and $n_B=3n_0$, } \\end{figure} \\begin{figure} \\psfig{figure=spp.eps,height=0.5\\linewidth} \\caption{ Variation of {\\bf{pp}} component of electrical conductivity  with magnetic field $B$. Curve (a): $T=15$MeV and $n_B=3n_0$, Curve (b): $T=15$MeV and $n_B=5n_0$, and Curve (c): $T=30$MeV and $n_B=3n_0$, } \\end{figure} \\begin{figure} \\psfig{figure=szp.eps,height=0.5\\linewidth} \\caption{ Variation of {\\bf{zp}} component of electrical conductivity  with magnetic field $B$. Curve (a): $T=15$MeV and $n_B=3n_0$, Curve (b): $T=15$MeV and $n_B=5n_0$, and Curve (c): $T=30$MeV and $n_B=3n_0$, } \\end{figure}"
    },
    "0404/astro-ph0404403_arXiv.txt": {
        "abstract": "We present new empirical calibrations of the Red Giant Branch (RGB) Bump and Tip based on a homogeneous near\\--Infrared database of 24 Galactic Globular Clusters. The luminosities of the RGB Bump and Tip in the J, H and K bands and their dependence on the cluster metallicity have been studied, yielding empirical relationships. By using recent transformations between the observational and theoretical planes, we also derived similar calibrations in terms of bolometric luminosity. Direct comparison between updated theoretical models and observations show an excellent agreement. The empirical calibration of the RGB Tip luminosity in the near\\--Infrared passbands presented here is a fundamental tool to derive distances to far galaxies beyond the Local Group, in view of using the new ground-based adaptive optics facilities and, in the next future, the James Webb Space Telescope. ",
        "introduction": "This is the second in a series of papers aimed at studying the Red Giant Branch (RGB) photometric properties and its major evolutionary features in the near\\--Infrared (IR) spectral domain. In the first paper \\citep[][ hereafter Paper~I]{pap1} an extensive database, collected by our group over the last 10 years, has been presented. By combining a new sample of 10 Galactic Globular Clusters (GGCs) belonging to different galactic populations (i.e. Halo and Bulge) with the data set published by \\citet[][ hereafter F00]{F00}, \\citet{V04} and \\citet[][ hereafter S04]{BB}, a homogeneous sample of 24 GGCs spanning a wide metallicity range ($-2.12{\\leq}$[Fe/H]${\\leq}$-0.49) has been obtained (see Table~\\ref{par}). In Paper~I, the entire database, calibrated in the {\\it Two Micron All Sky Survey} (2MASS) photometric system, has been used to measure a set of photometric indices describing the RGB: {\\it i)~} location (colours and magnitudes), and {\\it ii)~} morphology (slope). Updated calibrations of these indices in terms of the clusters metallicity have been also derived. More details on the dataset can be found in Paper~I.\\\\ In the present paper the major RGB evolutionary features, namely the Bump and the Tip, are measured in the J, H and K bands. Their calibrations as a function of the clusters metallicity, in both the observational and theoretical planes, are derived and discussed.\\\\ The RGB Tip luminosity is a bright standard candle and turns to be particularly useful to measure galaxy distances. The extension of the RGB Tip calibrations to the near\\--IR passbands is of fundamental importance since the recent advent of ground-based adaptive optics systems and the future availability of the James Webb Space Telescope are allowing to spatially resolve bright giants and to accurately measure their magnitudes up to a few Mpc distances.\\\\ In \\S 2 and \\$ we measure the near\\--IR magnitudes of the RGB Bump and Tip, respectively, of our sample of 24 GGCs and we investigate their behaviour with varying the cluster metallicity. The results of the transformations between the observational and theoretical planes for the RGB Bump and Tip luminosities are presented in \\S 4. Finally, in \\S 5 we briefly summarize our results.  \\begin{table} \\begin{center} \\label{par} \\caption{Adopted parameters for the program clusters.} \\begin{tabular}{lcccc} \\hline\\hline \\\\ Cluster& [Fe/H]$_{CG97}$ &[M/H] & E(B\\--V) & (m\\--M)$_0$ \\\\ \\\\ \\hline \\\\ M 92 &-2.16  &-1.95 & 0.02 & 14.78\\\\ M 15& -2.12   & -1.91 & 0.09 &15.15\\\\ M 68 &-1.99  &  -1.81 & 0.04 & 15.14\\\\ M 30 &-1.91  &  -1.71 & 0.03 & 14.71 \\\\ M 55 &-1.61  &  -1.41 & 0.07 & 13.82 \\\\ ${\\omega}$~Cen & -1.60 & -1.39 & 0.11 & 13.70 \\\\ NGC 6752&-1.42  &  -1.21 & 0.04 &13.18 \\\\ M 10  &-1.41 &-1.25 & 0.28 & 13.38 \\\\ M 13 &-1.39  &-1.18 & 0.02 & 14.43 \\\\ M 3  &-1.34  &-1.16 & 0.01 & 15.03 \\\\ M 4  &-1.19  & -0.94  & 0.36 & 11.68 \\\\ NGC 362 &-1.15  &  -0.99 & 0.05 & 14.68 \\\\ M 5  &-1.11  &-0.90 & 0.03 & 14.37 \\\\ NGC 288 &-1.07  &  -0.85 & 0.03 & 14.73 \\\\ M 107 &-0.87 & -0.70 & 0.33 & 13.95 \\\\ NGC 6380 &-0.87 &  -0.68 & 1.29 & 14.81 \\\\ NGC 6342 &-0.71 &  -0.53 & 0.57 & 14.63 \\\\ 47 Tuc&-0.70 & -0.59 & 0.04 & 13.32 \\\\ M 69 &-0.68  & -0.55 & 0.17 & 14.64 \\\\ NGC 6441 &-0.68 &  -0.52 & 0.52 & 15.65 \\\\ NGC 6624 &-0.63 &  -0.48 & 0.28 & 14.63 \\\\ NGC 6440&-0.49  &  -0.40 & 1.15 & 14.58 \\\\ NGC 6553&-0.44;-0.06$^*$ &  -0.36;-0.05$^*$ & 0.84 & 13.46 \\\\ NGC 6528 &-0.38;+0.07$^*$&  -0.31;+.05$^*$& 0.62 & 14.37\\\\ \\\\ \\hline \\end{tabular} \\end{center} (*) The most recent metallicity estimate by means of high resolution spectroscopy from \\citet{eugi01}. \\end{table} ",
        "conclusions": "New calibrations of the RGB Bump and Tip in the J, H and K bands as well as in bolometric , based on a global sample of 24 GGCs, have been presented. The behaviour of these evolutionary features have been investigated with varying the cluster metallicity in the CG97 and global scale, thus taking into account the effect of the ${\\alpha}$\\--enhancement.\\\\ Quadratic and linear best-fitting relations linking the RGB Bump and Tip magnitudes and metallicity, respectively, have been obtained (see eqs. 1-16). Comparisons between observations and theoretical models show a good agreement.\\\\ The RGB Bump and Tip represent powerful tools to obtain independent estimates of metallicity and distance, respectively, in old stellar systems within the Local Group. The new IR adaptive optics facilities available at ground-based 8m-class telescopes as well as the future imaging capabilities of the James Webb Space Telescope are allowing to use the RGB Tip distance indicator in galaxies well beyond the Local Group, up to a few Mpc away."
    },
    "0404/astro-ph0404545_arXiv.txt": {
        "abstract": "We use multi-frequency information in first year \\WMAP\\ data to search for the Sunyaev-Zeldovich (SZ) effect. WMAP has sufficiently broad frequency coverage to constrain SZ without the addition of higher frequency data: the SZ power spectrum amplitude is expected to increase 50\\% from W to Q frequency band. This, in combination with the low noise in \\WMAP, allows us to strongly constrain the SZ contribution. We derive an optimal frequency combination of \\WMAP\\ cross-spectra to extract SZ in the presence of noise, CMB, and radio point sources, which are marginalized over. We find that the SZ contribution is less than 2\\% (95\\% c.l.) at the first acoustic peak in W band. Under the assumption that the removed radio point sources are not correlated with SZ this limit implies $\\sigma_8<1.07$ at 95\\% c.l. We investigate the effect on the cosmological parameters of allowing an SZ component. We run Monte Carlo Markov Chains with and without an SZ component and find that the addition of SZ does not affect any of the cosmological conclusions. We conclude that SZ does not contaminate the \\WMAP\\ CMB or change cosmological parameters, refuting the recent claims that they may be corrupted. ",
        "introduction": "{\\em Wilkinson Microwave Anisotropy Probe} \\cite{2003ApJS..148....1B} (\\WMAP) observations of the Cosmic Microwave Background (CMB) have ushered in a new era of high precision observational cosmology. Such a tremendous increase in data quality requires a corresponding increase in the care that goes into the data analysis and interpretation. One of the lingering concerns surrounding the analysis is the residual effect from additional sources of anisotropies. A very prominent candidate among these is the Sunyaev-Zeldovich effect (SZ). Electrons in hot gas scatter photons and distort the blackbody spectrum of the CMB. Galaxy clusters, where gas is the hottest, contribute the bulk of the effect, shock-heating the gas in their potential wells. Radio and microwave band observations pointed at known clusters routinely yield SZ detections, so this effect is now well established observationally. The scattering preferentially raises the energy of the CMB photons, but the number of scatterings is low, so the process never achieves thermal equilibrium.  Therefore SZ appears as a CMB temperature decrement at low frequencies and as an increment at high frequencies. This frequency dependence is well known: it is constant in the Rayleigh-Jeans (RJ) regime of the blackbody spectrum, and is universal (independent of gas temperature or density) for non-relativistic electrons.  In the channels of \\WMAP, SZ is a temperature decrement.  Table \\ref{tab:wmapbands} provides \\WMAP's frequency bands, and the SZ frequency dependence in those bands.  The K and Ka bands are the most heavily polluted by galactic contamination, so the best prospects for identifying SZ in WMAP are in the differencing assemblies of the upper three bands: two assemblies in Q, two in V, and four in W. While it is usually argued that SZ is indistinguishable from the CMB in the \\WMAP\\ channels this is actually not so: in CMB temperature units the SZ power spectrum increases by 50\\% from W to Q channel and in W it is only 61\\% of the RJ power. We will show that this suffices to place strong constraints on SZ using \\WMAP\\ data alone, a consequence of the remarkably low noise in \\WMAP\\ data. \\begin{table} \\begin{center} \\begin{ruledtabular} \\begin{tabular}{lrrr} % Band& $\\nu$ (GHz) & ${\\Delta T^\\text{SZ}/(T^\\text{CMB} y)}$ & $C_l/C_l^\\text{RJ}$\\\\ \\hline % K    &   23             &   $-1.97$ &   $0.972$  \\\\ % Ka   &   33             &   $-1.94$ &   $0.945$  \\\\ % Q   &    41             &   $-1.91$ &   $0.915$  \\\\ % V   &   61              &   $-1.81$ &   $0.820$  \\\\ % W   &   94              &   $-1.56$ &   $0.611$  \\\\ % \\end{tabular} \\end{ruledtabular} \\end{center} \\caption{SZ contribution in the bands of \\WMAP.  We note the band name, frequency, temperature perturbation relative to the comptonization parameter $y$, which does not depend on frequency, and power spectrum relative to the power in the Rayleigh-Jeans regime.  In the \\WMAP\\ bands, SZ is a temperature decrement.  In the Rayleigh-Jeans regime, ${\\Delta T^\\text{SZ}/(T^\\text{CMB} y)}=-2$. }\\label{tab:wmapbands} \\end{table}  In the literature, several groups have attempted to identify SZ in \\WMAP\\ data using cross-correlations with other tracers of large scale structure (LSS). In \\citeN{2003ApJS..148...97B} they build an SZ template from the XBACs catalog of x-ray clusters \\cite{1996MNRAS.281..799E}, and fit for the amplitude of SZ, arguing for a $2.5\\sigma$ detection. In \\citeN{2003astro.ph..8260A} they compute the cross-power spectra  of the 2MASS extended source catalog \\cite{2000AJ....119.2498J} with \\WMAP's Q, V, and W bands, then fit for the amplitude of the SZ signal, arguing for a 3.1--3.7$\\sigma$ detection of SZ, depending on their mask. In \\citeN{2003MNRAS.346..940D} they compute the cross-power spectrum of the ROSAT diffuse x-ray background maps \\cite{1997ApJ...485..125S} and  a \\WMAP\\ (Q-W) difference map, finding no detection. In \\citeN{2003ApJ...597L..89F} they compute the cross-correlation function of SDSS DR1 \\cite{2003AJ....126.2081A} galaxy survey data with the V band and fit for the amplitude of SZ, finding a 2.7$\\sigma$ detection. In \\citeN{2004MNRAS.347..403H} they build templates for SZ from several optical and x-ray cluster surveys, then fit for the amplitude of these templates using maps from the W band and \\citeN{2003PhRvD..68l3523T}, finding no detection for optical clusters and 2--5$\\sigma$ detections for x-ray clusters.  Finally, in \\citeN{2004MNRAS.347L..67M} they compute the mean temperature in concentric rings about APM \\cite{1990MNRAS.243..692M} and ACO \\cite{1989ApJS...70....1A} groups and clusters, noting a decrement they attribute to SZ.  They interpret this decrement to extend to large scales, $\\sim 1$ degree, although the covariance of their bins is unclear.  Redshift $z<0.2$ clusters dominate their sample.  Their most extreme model assumes that extended SZ emission is representative of the temperature and spatial clustering of gas to $z=0.5$.  In this model the SZ power is 30\\% of the first acoustic peak of the CMB.  \\ Thus they conclude the integrity of the \\WMAP\\ cosmological parameter fits may be compromised.  In this paper we seek to place limits on SZ from \\WMAP\\ data alone, thus avoiding any uncertainties in connecting the results of \\WMAP-LSS cross-correlation to those of \\WMAP\\ auto-correlation.  The cross-correlation methods used before require a model for relating SZ temperature perturbations to some proxies for SZ, which are nonlinear structures such as clusters, and this connection can be quite uncertain.  Our method is less model dependent. The downside is that \\WMAP-only methods sacrifice signal to noise because they do not focus on matter concentrations, where SZ should be strongest, so our method is not optimized for obtaining an SZ detection. However, our main goal is to investigate the amount to which the CMB analysis may be compromised by the residual SZ component, for which we need analysis as model independent as possible.  Our principal method, described in section \\ref{sec:est}, constructs a linear combination of the \\WMAP\\ Q, V, and W band cross-power spectra which maximizes the contribution of SZ, while at the same time minimizing the radio point source and CMB contribution.  We use this linear combination to fit for the SZ power spectrum amplitude, using a spectrum shape from halo model calculations \\cite{2002MNRAS.336.1256K}.  In section \\ref{sec:markov} we supplement this method by fitting for cosmological parameters from the \\WMAP\\ temperature power spectrum using the Markov chain Monte Carlo method, allowing for an additional SZ component in the power spectrum with a free amplitude to be determined by the data. The two methods give consistent results, with the latter giving somewhat weaker constraints. Conclusions are presented in section \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We have analyzed the \\WMAP\\ power spectra in two ways and found them to be free of serious contamination from SZ. The more powerful and less model dependent of the two methods is to use multi-frequency information: by combining the \\WMAP\\ cross-spectra, we limit at 95\\% confidence the amplitude of SZ to be below 2\\% of the CMB at the position of first peak in W band. By searching for an SZ-shaped component in the \\WMAP\\ combined spectrum with a Markov chain Monte Carlo we also do not find any evidence of a signal, but we can only set a weaker limit. Combining the analyses, we show that the cosmological parameters are not affected by the SZ within the range allowed by the multi-frequency analysis.  There are several possible improvements in the analysis that could further tighten the limits and may be applied to the second year data once they become available. In addition to having lower noise, the two year data are also expected to have better control of systematics such as beam uncertainties  and noise correlations.  First, we would like to improve our method for obtaining the cross-spectrum covariance matrix.  The estimate for the covariance is noisy, which makes our weights for combining the power spectra noisier than they should be.  If we double the size of the bins we average over to obtain the covariance, our limit changes slightly ($B<4.3$ opposed to $B<3.3$ at 95\\%) and the point source estimate is virtually unchanged.  However, it is impossible to tell if this effect is due to a covariance matrix with less noise (which is better) or due to weights whose wide bins combine high- and low-signal-to-noise multipoles less optimally (which is worse). This can be improved if the \\WMAP\\ second year data release includes the cross-spectrum covariance matrix.  The second possible improvement of our analysis  is to improve the frequency dependence of point sources or the \\WMAP\\ beam deconvolution, both of which can bias our estimate.  Our estimator is only unbiased if the models one applies to it are faithful to the data.  The models we have used are the best available, and this information will improve in the second year data. It seems unlikely that these improvements will strongly modify the conclusions in this paper.  The third improvement is to perform the analysis on the maps which do not have point sources excised. As discussed above, masking of resolved point sources from the \\WMAP\\ data may reduce the SZ signal if the two are correlated. This would weaken our limit on $\\sigma_8$ from the absence of SZ signal, but does not change our conclusions about the effect of SZ on the CMB and cosmological parameters. Still, given that the current limits are reaching the levels of interest it is worth exploring this possibility further.  The fourth improvement is to verify the assumed shape for the SZ power spectrum. This is currently based on halo models, since hydrodynamic simulations cannot simulate sufficiently large volumes to make predictions reliable. If the power spectrum has a radically different shape, our limit is hard to interpret, but the level of contamination on primary CMB is likely to remain unchanged, as is clear from  Figure \\ref{fig:szestpow}.  As an example we estimated the SZ in bins of $\\Delta l=120$.  In the bin containing the first acoustic peak, we limit the SZ power to 4.9 times the halo model power predictions at 95\\% confidence or 4\\% contamination in W. It is clear that as long as the SZ spectrum is smooth the limits on the contamination are not going to change significantly over the range with best sensitivity at $100<l<300$.  Finally, our procedure to marginalize over CMB and point sources is conservative and the errors could be further reduced if a joint estimation is used instead. We know something more about the CMB power spectrum than just its frequency dependence, and in this case we can use this information. We also note that a relatively trivial modification to \\WMAP's method for estimating the power spectrum could eliminate any contaminant with an SZ-like shape and frequency dependence, much in the same way that they eliminate point sources. However, the fact that we find no residual SZ effects implies that this procedure is not really required and justifies the treatment of the \\WMAP\\ team in ignoring SZ.  Even with future data detecting the SZ power spectrum from \\WMAP\\ alone may be difficult.  If the bulk of the covariance is from noise, it would take \\WMAP\\ roughly 3 years to get a 1$\\sigma$ detection of SZ using the cross-spectrum combination method of this paper, if $\\sigma_8=0.9$. Elimination of non-noise sources of covariance, such as uncertainties remaining in the beams, may help. If $\\sigma_8=1$ then a 2$\\sigma$ detection should be possible with 4 year data and one should start seeing a hint of the signal already with 2 year data. Theoretical predictions remain somewhat uncertain, so it is possible that these predictions have a systematic uncertainty of around 10\\%, but it is clear that it will be worth looking for SZ in future WMAP releases."
    },
    "0404/physics0404114_arXiv.txt": {
        "abstract": "The dimensionless kinetic energy dissipation rate $C_{\\varepsilon }$ is estimated from numerical simulations of statistically stationary isotropic box turbulence that is slightly compressible. The Taylor microscale Reynolds number $(\\Rey_{\\lambda})$ range is $20\\lesssim \\Rey_{\\lambda }\\lesssim 220$ and the statistical stationarity is achieved with a random phase forcing method. The strong $\\Rey_{\\lambda }$ dependence of $C_\\varepsilon$ abates when $\\Rey_{\\lambda } \\approx 100$ after which $C_\\varepsilon$ slowly approaches $\\approx 0.5,$ a value slightly different to previously reported simulations but in good agreement with experimental results. If $C_\\varepsilon$ is estimated at a specific time step from the time series of the quantities involved it is necessary to account for the time lag between energy injection and energy dissipation. Also, the resulting value can differ from the ensemble averaged value by up to $\\pm 30\\%.$ This may explain the spread in results from previously published estimates of $C_\\varepsilon.$ ",
        "introduction": "The notion that the mean turbulent kinetic energy dissipation rate $\\varepsilon$ is finite and independent of viscosity $\\nu$ was originally proposed by G.~I. Taylor\\cite{t35}. Its importance is so recognized now that it is commonly referred to as the ``zeroth law'' of turbulence. Its existence was assumed by von K\\'{a}rm\\'{a}n and Howarth, Loitsianskii and also, significantly, Kolmogorov\\cite{k41} in establishing his celebrated similarity hypotheses for the structure of the inertial range of turbulence. Kolmogorov assumed the small scale structure of turbulence to be locally isotropic in space and locally stationary in time - which implies the equality of turbulent kinetic energy injection at the large scales with the rate of turbulent kinetic energy dissipation at the small scales\\cite{ccddt97}. Although this view should be strictly applied only to steady turbulence, the mechanism of the dissipation of turbulent kinetic energy can be considered the most fundamental aspect of turbulence not only from a theoretical viewpoint but also from a turbulence modeling viewpoint. Indeed, the mechanism that sets the level of turbulent dissipation in flows that are unsteady is a difficult, if not intractable, aspect of turbulence modeling.  The rate of turbulent kinetic energy dissipation is determined by the rate of energy passed from the large-scale eddies to the next smaller scale eddies via a forward cascade until the energy is eventually dissipated by viscosity. Thus, $C_\\varepsilon$ defined as, \\begin{equation} C_{\\varepsilon}= \\varepsilon  L/{u'}^3 , \\label{eq:ceps} \\end{equation} (here, $ \\varepsilon $ is the mean energy dissipation rate per unit mass, $L$ and $u'$ are characteristic large length and velocity scales respectively) should be independent of the Reynolds number and of order unity. An increase in Reynolds number should only result in an increase in the typical wave number where dissipation takes place\\cite{l92}. In the past few years there have been a number of numerical (see Ref.~\\cite{kiyiu03} and references therein) and experimental (see Refs.~\\cite{pkw02,w02,pkj03} for recent results) efforts to determine the value of $C_\\varepsilon$ and its dependence on the Reynolds number. Perhaps the most convincing of these are the numerical attempts since there is no re-course to one-dimensional surrogacy as there is for experiments. Notwithstanding this fact, there is good agreement, both numerically and experimentally, with the long held view that $C_\\varepsilon$ is $\\sim O(1)$ when the Reynolds number is sufficiently high. The collection of isotropic simulation results for $C_\\varepsilon$ shown in Ref. \\cite{kiyiu03} indicates that ``high enough'' Reynolds number appears to be $\\Rey_\\lambda \\sim O(100).$ Here, $\\Rey_\\lambda (={u'}^2[15/\\nu\\varepsilon ]^{1/2})$ is the Taylor microscale Reynolds number. At higher $\\Rey_\\lambda$, slow $\\Rey_\\lambda$ dependencies for $C_{\\varepsilon}$, such as that proposed by Lohse\\cite{l94}, cannot be ruled out. Measuring such $\\Rey_\\lambda$ dependencies, either numerically or experimentally, will be close to impossible.  One unresolved issue is that raised by Sreenivasan\\cite{s98}. After assembling all the then known experimental decaying grid turbulence data\\cite{s84} and numerical data for both decaying and stationary isotropic turbulence he concludes that ``the asymptotic value (of $C_\\varepsilon$) might depend on the nature of large-scale forcing, or, perhaps, on the structure of the large scale.'' He also demonstrates\\cite{s95b} in homogeneously sheared flows that the large structure does influence $C_\\varepsilon.$ However, it might be argued that these results were obtained at low Reynolds numbers and the issue of a universal asymptotic value for $C_\\varepsilon$ could still be considered open. Alternatively it could be argued that homogeneous shear flows and the like are strictly unsteady turbulent flows and the zeroth law, in its simplest guise, should not be expected to apply to such flows. The possibility of some characteristics of large-scale turbulence being universal should not be ruled out. The recent observation that input power fluctuations, when properly re-scaled, appear universal\\cite{bhp98} may be construed to suggest the possibility of universality for $C_{\\varepsilon}$. The aim of the present work is to estimate $C_\\varepsilon$ from direct numerical simulations (DNS) of statistically stationary isotropic turbulence and compare with previously reported DNS results (summarized in Fig.~3 of Ref.~\\cite{kiyiu03}) and experiments carried out in regions of low $\\left (dU/dy \\approx dU/dy|_{\\rm max}/2 \\right )$ or zero mean shear. The present DNS scheme differs from methods already reported in that a high-order finite difference method is used. To our knowledge, these are the first finite difference results for $C_\\varepsilon.$ Hence, it is worthwhile to test if different numerics and forcing at the large scales result in vastly different values for $C_\\varepsilon$ to those already reported.  \\begin{table*} \\begin{ruledtabular} \\begin{tabular}{lccccccccccccc} Run & $N$ & $\\Rey_\\lambda$ & $T_{\\rm tot}/T$ & $\\nu \\left(\\times 10^{4}\\right )$ & $ \\varepsilon \\left(\\times 10^{5}\\right )$&$\\Delta t/t_\\kappa\\left(\\times 10^{2}\\right )$  &$L$ & $\\lambda$ & $u'\\left(\\times 10^{2}\\right )$ & $\\tau_{\\rm max}/T$ & $C_\\varepsilon$ & $\\eta \\left(\\times 10^{3} \\right )$ & $k_{\\max} \\eta$ \\\\ \\hline A & 32 &  20 &  31 &  40  & 24 & 1.9  &1.9 &  1.2 &  7.1 &  0.15 &  1.2 & 128 & 2.1\\\\ % B & 64 &  42 &  30 &  15  & 22 & 1.5  &1.6 & 0.81 &  7.8 &  0.37 & 0.75 &  63 & 2.0\\\\ % C &128 &  90 &  11 &  4.0 & 24 & 1.5  &1.3 & 0.43 &  8.4 &  0.62 & 0.54 &  23 & 1.5\\\\ % D &256 &  92 &  19 &  4.0 & 21 & 0.71 &1.4 & 0.45 &  8.1 &  0.69 & 0.53 &  24 & 3.0\\\\ % E &256 & 152 &  20 &  1.6 & 21 & 1.1  &1.4 & 0.29 &  8.4 &  0.74 & 0.49 &  12 & 1.5\\\\ % F &512 & 219 &   7 &  0.80& 25 & 0.86 &1.3 & 0.20 &  8.9 &  0.67 & 0.47 &   7 & 1.7\\\\ % \\end{tabular} \\end{ruledtabular} \\caption{Examples of DNS parameters and average turbulence characteristics. $N$ is the number of grid points in each of the Cartesian directions, $\\Rey_\\lambda$ is the Taylor microscale Reynolds number $\\equiv  u' \\lambda/\\nu$, $T_{\\rm tot}$ is the total run time after the run became statistically stationary, $T$ is the eddy turnover time $\\equiv L/u'$, $\\Delta t$ is the run time increment, $t_\\kappa$ is the Kolmogorov time scale $\\equiv \\nu ^{1/2} \\varepsilon ^{-1/2},$ $\\lambda$ is the Taylor microscale $\\equiv u'\\sqrt{15 \\nu/\\varepsilon},$ $\\tau_{\\rm max}$ is the average time for the energy cascade from large to small scales, and $\\eta$ is the Kolmogorov length scale $\\equiv \\nu ^{3/4} \\varepsilon ^{-1/4}$.}  \\label{tab:average results} \\end{table*} ",
        "conclusions": "The present work has revisited the zeroth law of turbulence for both numerical simulations of statistically stationary isotropic turbulence and experiments. The numerical simulations are slightly compressible isotropic turbulence and the statistical stationarity is achieved with a random phase forcing applied at low wave numbers. The main result of the numerical simulations is the demonstration that $C_\\varepsilon$ should only be estimated with ensemble averaged quantities from the entire time series for which the statistics are stationary. If $C_\\varepsilon$ is to be estimated at each time snap shot it is necessary to correctly account for the time lag that occurs from the large scale energy injection to the fine scale energy dissipation. Even after correctly correlating the energy injection with the energy dissipation, the instantaneous value of $C_\\varepsilon$ can vary quite considerably (e.g.\\ $\\pm 30\\%$) over the extent of the simulation. Such a variation may account for the scatter in magnitude of $C_\\varepsilon$ in previously published results. Both the present numerical and experimental results suggest that the asymptotic value for $C_\\varepsilon$ is $\\approx 0.5.$ In light of this, the previously held view that the asymptotic value of $C_\\varepsilon$ may be dependent on the large scale energy injection could be suspect. Lastly, the results presented are strictly applicable only to isotropic turbulence that is stationary in time. However, it would be interesting to estimate $C_\\varepsilon$ for simulations of  turbulence unsteady in space and/or time e.g. anisotropic turbulence or anisotropic homogeneous turbulence with a mean shear because there is little known for these flows about how the turbulent kinetic energy is dissipated."
    },
    "0404/astro-ph0404259_arXiv.txt": {
        "abstract": "{By means of simple numerical models we discuss whether {\\it collisionless stellar evaporation} from cluster elliptical galaxies could be an effective mechanism for the production of intracluster stellar populations.  The effectiveness of this mechanism is due to the fact that, for realistic galaxy and cluster models, the galaxy oscillation periods near equilibrium configurations in the cluster tidal field are of the same order of stellar orbital times in the external parts of the galaxies themselves.  With the aid of Monte-Carlo simulations we explore the evolution of stellar orbits in oscillating galaxies placed near different equilibrium positions.  We found that, over an Hubble time, the main effect is a substantial expansion of the galactic outskirts, particularly affecting the galaxy at the cluster center and those orbiting near the cluster core radius: overall, approximately the 10\\% of the galaxy mass is affected.  Thus, the proposed mechanism could be of some importance in the shaping of the halo of cD galaxies and in making easier ``galaxy harassment'' in the formation of the intracluster stellar population. ",
        "introduction": "Observational evidences of an intracluster stellar population (hereafter ISP) are mainly based on the detection of planetary nebulae and red giant branch stars freely floating in the intergalactic space. For example, Theuns \\& Warren (1996) identified 10 intergalactic planetary nebulae in the Fornax cluster, while M\\'endez et al. (1997) detected 11 intergalactic objects in the Virgo cluster whose cumulative luminosity functions is in good agreement with planetary nebula luminosity function. In addition, Ferguson et al. (1998) identified an intergalactic red giant branch stars population in the Virgo cluster, while Feldmeier et al. (1998) (with observations of three blank fields in the Virgo cluster), confirmed that a significant fraction of Virgo's starlight is due to the ISP.  More recently Okamura et al. (2002) have identified 38 candidates of intracluster planetary nebulae in the core of the Virgo cluster.  Overall, the data suggest that approximately 10\\% (or even more) of the stellar mass of the cluster is in intergalactic stars (e.g., see Ferguson et al. 1998, Arnaboldi et al. 2002, Durrel et al. 2002, Arnaboldi et al. 2003, Totani 2003, see also Arnaboldi, Gerhard \\& Freeman 2003, and references therein).  The usual scenario assumed to explain the finding above is that the tidal interactions between galaxies (for example during a fast encounter, e.g., see Merritt 1983, 1984, 1985), and of galaxies with the cluster gravitational field, lead to a substantial stripping of stars from galaxies to the parent cluster (the so-called ``galaxy harassment'' scenario, see, e.g., Moore et al. 1996; Napolitano et al. 2003).  In this paper we explore an additional ``stripping'' mechanism, namely we quantitatively discuss the effect of resonant interaction between stellar orbits inside the galaxies and the cluster tidal field (hereafter, CTF). The present study is supported by the fact that the characteristic oscillation times of a galaxy near its equilibrium position in the CTF and the mean stellar orbital times in the galaxy external regions are of the same order of magnitude, as already recognized by Ciotti \\& Giampieri (1998, hereafter CG98).  In fact, Hawley \\& Peebles (1975) reported a possible indication (confirmed by Thompson 1976), that the galaxies are preferentially aligned along the radius vector to the center of the cluster for the Coma cluster; they also suggested the CTF as the possible cause of the alignment.  The best evidence for alignment is for brightest cluster galaxies: Adams, Strom \\& Strom (1980) found, in 7 very elongated clusters, a general trend for ellipticals to be aligned with the cluster major axis, and also Trevese, Cirmele \\& Flin (1992) found a strong alignment of the brightest galaxy major with the long axis of the parent cluster.  Numerical N-body simulations (Ciotti \\& Dutta 1994, hereafter CD94) indeed confirmed the hypothesis that the CTF could be at the origin of the observed alignment and, in particular, revealed that their model elliptical galaxies (Es) behaved with good approximation as rigid bodies. In a complementary analytical exploration of this problem CG98 determined the equilibrium positions of triaxial ellipsoids in the CTF for various cases, and showed that the oscillations period of the galaxies, when slightly displaced from their equilibrium configurations, are of the same order of magnitude of the stellar orbital periods in the galaxy outskirts. This curious finding naturally leads to ask what could be the effect of possible ``resonances'' between stellar orbital periods in cluster Es and their oscillation periods. Unfortunately, the $N$-body simulations of CD94 were characterized by a limited number of particles (i.e., $N\\simeq 3\\times 10^4$) and by the usual softening, so that evaporations rates could not be properly investigated.  In order to investigate the scenario described above, Muccione \\& Ciotti (2003a,b, hereafter MC03a,b) performed a few preliminary Monte-Carlo simulations of orbital exploration of oscillating galaxies, and the results were encouraging, with estimated escape rates summing up to 10 per cent of the initial galaxy mass over an Hubble time. In the present paper we considerably extend these two previous investigations, by assuming that the galaxies are triaxial ellipsoids and two cases are considered: when the center of mass of the ellipsoid is at rest at the equilibrium point of the field generated by the cluster, and when it is placed on a circular orbit around the center of a spherically symmetric cluster.  After deriving the equations of the motion of a star inside an oscillating galaxy, we integrate numerically these equations under different initial conditions realized by using a Monte-Carlo extraction.  The paper is organized as follows.  In Sect. 2 we briefly review the proper physical setting of the problem, and in Sect. 3 we describe the adopted galaxy and cluster models.  In Sect. 4 we describe in the numerical integration scheme, while in Sect. 5 we present the main results, that are finally summarized and discussed in Sect.  6. In the Appendix, the specific technique adopted to obtain the gravitational field inside the galaxy models and other useful dynamical quantities are briefly described. ",
        "conclusions": "In this paper we presented a representative selection of numerical models aimed at the study of collisionless evaporation of stars from cluster elliptical galaxies. From a theoretical point of view this possibility was pointed out in CG98, who showed that characteristic oscillation times of elliptical galaxies in near equilibrium configurations in the tidal field of the parent cluster are curiously comparable to the stellar orbital times in the outskirts of the galaxies themselves. From a numerical point of view the present work represents the natural follow-up of two recent preparatory works on the same subject (MC03ab), where we explored the simplest equilibrium configuration (i.e., a galaxy with the center of mass at rest at the center of a triaxial cluster), but we left unaddressed the more complicate (although more astrophysically representative) case of galaxies with their center of mass in rotation around the (spherical) cluster center. We recall that in all the presented simulations the {\\it direct} effect of the cluster gravitational field on stellar orbits was {\\it not} included, in order to study the effects of galaxy oscillations only; we also recall that its effect is {\\it compressive} inside the cluster core radius, while {\\it expanding} outside (see, e.g., Valluri 1993, CD94).  The main results can be summarized as follows:  \\begin{itemize}  \\item A general analytical framework suitable for the study of orbital evolution of stars in galaxies oscillating around their stable equilibrium positions in the tidal field of the parent cluster is presented. In particular, we analyzed the cases of triaxial galaxies at the center of a triaxial cluster and in circular orbit around the center of a spherical cluster.  \\item It is shown that, for realistic galaxy and cluster models, galaxy oscillation times can be comparable to characteristic stellar orbital times in the external galactic regions, and thus important effects on stellar orbits are expected there.  \\item Numerical simulations of galaxiesat the cluster center remarkably confirm the preliminary results presented in MC03ab (where different galaxy models were adopted). In particular, we found that orbits in the external parts of the parent galaxy are affected by galaxy oscillations, with stars outside $\\mi\\simeq 3\\simeq 2.25\\Reff$ nearly doubling the apocenters. {\\it Globally, a galaxy mass fraction of the order of 10\\% is affected by this process}.  \\item We speculate that the findings above could be related to the formation of cD galaxies. In fact, it is well known that numerical simulations of galactic cannibalism, although successful at reproducing BCGs, are in general unable to obtain the extended halos of cDs. In this respect, the fact that the external parts of the galaxies are the only ones affected by collisionless evaporation, is a nice property of the explored scenario. Unfortunately, our simulations cannot be used to study the resulting galactic profiles, and high-resolution N-body simulations should be performed to study further the matter.  \\item For galaxies in circular orbits, the effect of collisionless evaporation is enhanced with respect to the previous case, and stars can reach significantly larger distances. This is not due to the direct effect of the CTF, but to galaxy oscillations only. In accordance with our preparatory analysis, we found that the radial trend of galactic oscillation periods at increasing distance from the cluster center is reflected by the location in the galaxy of the expanded region, being the most affected galaxies those orbiting near the cluster core radius. Even in this cases, {\\it the affected galaxy mass fraction ranges is of the order of 5\\% - 10\\%, in interesting agreement with observational estimates}.  \\item The natural consequence of the results listed above is that, if collisionless evaporation is of importance in the production of the ISP (as direct contributor and/or help to the stripping phenomena due the direct CTF effect and to galaxy-galaxy encounters, as investigated in detail by Merritt [1983, 1984, 1985]), the density of ISP (when normalized to the local cluster luminosity in galaxies) should be enhanced around the cluster core radius, and rapidly declining at the cluster outskirts, due to the increase of the galaxy oscillation periods with the distance from the cluster center.  \\end{itemize}  We finally conclude by pointing out that both orbital cases explored in this paper are rather exceptional.  First of all, we only considered galaxy models that were initially not tumbling nor rotating. In addition, most cluster galaxies neither rest in the cluster center nor move on circular orbits, but they move on elongated orbits with very different pericentric and apocentric distances from the cluster's center; in a triaxial cluster many orbits are boxes and some orbits can be chaotic. These latter cases can be properly investigated only by direct numerical simulation of the stellar motions inside the galaxies, coupled with the numerical integration of the equations of the motion of the galaxies themselves. In addition, a statistically significant galaxy population made by the sum of galaxies with different masses, scale-lengths, and flattenings, should be considered."
    },
    "0404/astro-ph0404290_arXiv.txt": {
        "abstract": "A model of a hydrogenic content of atmosphere of the isolated neutron star 1E1207.4-5209 is proposed. It is based on the assumption that the main component in the atmosphere is the exotic molecular ion $H_3^{2+}$ and that there exists a magnetic field in the range of $(4 \\pm 2) \\times 10^{14}$\\,G. Photoionization $H_3^{2+} \\rar e + 3p$ and photodissociation $H_3^{2+} \\rar H + 2p$ correspond to two absorption features at 0.7 KeV and 1.4 KeV, respectively, discovered by {\\it Chandra} observatory (Sanwal et al, 2002). The model predicts one more absorption feature at 80-150 eV corresponding to photodissociation $H_3^{2+} \\rar H_2^+ + p$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404015_arXiv.txt": {
        "abstract": "The conditions that lead to self-regulated star formation, star bursts and the formation of massive stellar clusters are discussed. Massive stars have a strong impact on their environment, especially on the evolution of dwarf galaxies which are the building blocks of giant galaxies. Energy input by massive young clusters might help to solve some of the most important puzzles of galaxy formation: the cosmological substructure problem and the angular momentum problem. ",
        "introduction": "Stars and gas in galaxies interact continuously through the exchange of matter and energy. It is this process that drives galactic evolution and that determines many of their observed characteristic properties. Low-mass stars have a limited effect on their environment by generating wind-blown bubbles and planetary nebulae. High-mass stars, on the other hand, disturb the dynamics of the surrounding interstellar medium on larger, that is galactic scales through their ionizing photons which destroy molecular material and through supernova explosions which act as point sources of thermal and kinetic energy input and which drive galactic chemical evolution. The most violent events are multiple supernova explosions that erupt on timescales of order $10^6$ yrs in young massive clusters (YMCs) and that can completely change the morphology and evolution of a galaxy by blowing its gas and metals out into intergalactic space. ",
        "conclusions": ""
    },
    "0404/astro-ph0404223_arXiv.txt": {
        "abstract": "{We report the discovery of three optical open clusters in the Milky Way. Two clusters are in Scutum (Cluster\\,1 at $\\ell=18.44^{\\circ}$ and $b=-0.42^\\circ$, and Cluster\\,2 at $\\ell=19.60^{\\circ}$ and $b=-1.02^\\circ$), thus projected not far from the Galactic center direction, and the other is in Canis Major (Cluster\\,3 at $\\ell=235.61^{\\circ}$ and $b=-4.10^\\circ$), near the anti-center direction. Cluster\\,3 is less populous than Clusters\\,1 and 2, but presents evidence of being a physical system. The objects were found optically by inspecting maps obtained from the Guide Star Catalogue and images from the Digitized Sky Survey. No previous identification of cluster has been reported in each area so far. The analysis was carried out with 2MASS photometry in J and H. For Cluster\\,1 we derive an age of $t=25\\pm5$\\,Myr, a reddening $\\ebv=2.18\\pm0.03$ and a distance from the Sun $\\ds=1.64\\pm0.19$\\,kpc; for Cluster\\,2, $t=500\\pm100$\\,Myr, $\\ebv=0.91\\pm0.03$ and $\\ds=2.19\\pm0.21$\\,kpc; finally for Cluster\\,3, $t=32-100$\\,Myr, $\\ebv=0.94\\pm0.03$ and $\\ds=3.93\\pm0.35$\\,kpc. Luminosity and mass functions are derived for Clusters\\,1 and 2 which, in turn, allowed us to estimate their observed masses as $147\\,\\ms$ and $89\\,\\ms$, respectively. Estimated total masses, by extrapolating the mass functions to $0.08\\,\\ms$, amount to $382\\,\\ms$ and $614\\,\\ms$, for the two clusters. Cluster\\,3 has an observed mass of $\\sim55\\,\\ms$. The present results indicate that further searches in the optical might still reveal new open clusters, and more so in infrared bands. ",
        "introduction": "The discovery of new physical stellar systems in the Galaxy is essential for the task of completing their census. This, along with the determination of properties such as distance, age, mass and dynamical state, and spatial distribution and evolutionary processes in the disk (e.g. Janes \\& Adler \\cite{JA1982}, Twarog et al. \\cite{Twa1997}) will give a better statistical insight on their formation, evolution and eventual disruption. Open clusters span a wide range of ages from the very young to the intermediate age ones, whose ages and distances can be determined by fitting isochrones to their colour-magnitude diagrams (CMDs), with a precision depending on the depth of the photometry and field contamination. These objects are formed along the Galactic plane where there is an abundance of gas and dust and their orbits become perturbed by the cumulative effect of passages near interstellar clouds. Due to their spatial location, young open clusters can be considered as tracers of the Galaxy's spiral structure (Chen et al. \\cite{Chen2003}). Thus, the discovery and characterization of additional open clusters represent a step further in the understanding of Galaxy structure and Galaxy formation processes.  In the recent decades some new open clusters have been discovered in the optical domain, e.g. Pfleiderer et al. (\\cite{Pfleiderer1977}), Turner et al. (\\cite{Turner1986}) and Saurer et al. (\\cite{Saurer1994}). Recently, two optical open clusters in Cygnus OB2 have been found and their properties determined (Bica et al. \\cite{BBD2003}). In the infrared, the amount of new embedded clusters is striking, mostly due to the recent release of the Two Micron All Sky Survey (hereafter 2MASS, Skrutskie et al. \\cite{2mass1997}) catalogue and atlas. We point out the recent discovery of 346 embedded clusters and candidates along the spiral arms of the Galaxy (Dutra et al. \\cite{Dutra2003a}, Bica et al. \\cite{Bica2003}).  In the present study we report the finding of three new optical open clusters. Their properties will be derived by means of 2MASS photometry, since these optical objects show up clearly in the near-infrared. In addition, the uniform and essentially complete sky coverage provided by 2MASS allows one to properly take into account background regions with suitable star count statistics, which is fundamental in order to correctly identify and characterize the stellar content of clusters.  In Section~\\ref{The3NOC} we present the new open clusters and show the optical images where the clusters have been found. In Sect.~\\ref{2massPh} we obtain the 2MASS photometry and introduce the $\\jj\\times\\jh$ CMDs. In Sect.~\\ref{StructAnal} we discuss the radial density distribution of stars and derive structural parameters for the clusters. In Sect.~\\ref{Fund_par} we fit isochrones to the near-infrared CMDs and derive cluster parameters. In Sect.~\\ref{LumFunc} we derive the luminosity and mass functions (hereafter LF and MF) and estimate the stellar masses of each cluster. Concluding remarks are given in Sect.~\\ref{Conclu}. ",
        "conclusions": "\\label{Conclu}  The recent discoveries of star clusters in the Galactic disk have shown that the census of these objects is not complete, even in the optical domain. In this paper we report the discovery of three optical open clusters in the Galaxy. Two clusters are in Scutum, at very low Galactic latitudes, and the other in Canis Major. The present photometric and structural analyses make use mostly of J and H 2MASS All Sky data release photometry.  Arguments in favour of their nature as open clusters are: {\\it (i)} the visual contrast between the stellar concentrations in the cluster and background regions (Figs.~\\ref{fig1}, \\ref{fig2} and \\ref{fig3}); {\\it (ii)} the CMDs which present well-defined sequences and are different from the offset field ones (Fig.~\\ref{figCMD}); {\\it (iii)} the radial star density profiles which show important excesses over the background and can be fitted by a King's law (Fig.~\\ref{figRAD}); {\\it (iv)} the constrained solutions of the isochrone fits to the CMDs, allowing a narrow range of ages (Figs.~\\ref{figIsocC1}, \\ref{figIsocC2} and \\ref{figIsocC3}); and {\\it (v)} the background-subtracted LFs which present significant star excesses over the background down to A\\,2--A\\,3 stars for Clusters\\,1 and 2, and to F\\,0 stars for Cluster\\,3 (Fig.~\\ref{figLF}). Completeness effects and photometric errors do not allow inferences on lower masses. For the sampled magnitude range, the number of member stars turns out to be 103, 80 and 26, respectively for Clusters\\,1, 2 and 3.  Fundamental parameters for the objects have been derived by fitting solar-metallicity Padova isochrones to the 2MASS $\\mj\\times\\jh$\\ CMDs. Cluster\\,1 has an age $t=20 - 32$\\,Myr, a reddening $\\ebv=2.18\\pm0.03$ and a distance to the Sun $\\ds=1.64\\pm0.19$\\,kpc. For Cluster\\,2 we derive $t=400 - 630$\\,Myr, $\\ebv=0.91\\pm0.03$ and $\\ds=2.19\\pm0.21$\\,kpc. For Cluster\\,3, $t=32 - 100$\\,Myr, $\\ebv=0.94\\pm0.03$ and $\\ds=3.93\\pm0.35$\\,kpc. At the above distances, the core radii for Clusters\\,1, 2 and 3 turn out to be $0.32\\pm0.03$\\,pc, $0.68\\pm0.06$\\,pc and $0.21\\pm0.02$\\,pc, respectively. The linear limiting diameters are $2.6\\pm0.3$\\,pc, $4.4\\pm0.4$\\,pc and $3.8\\pm0.3$\\,pc, respectively, thus suggesting dynamical evolution effects for the older clusters.  For Clusters\\,1 and 2, a MF fit $\\phi(m)\\propto m^{-(1+\\chi)}$ resulted in a slope $\\chi=0.85$ and 1.04, respectively. Thus, the observed stellar mass in Cluster\\,1 is $147\\,\\ms$. In Cluster\\,2, including the evolved stars, it is $89\\,\\ms$. The stellar mass in Cluster\\,3 is $\\sim55\\,\\ms$. Extrapolating the MF fits down to the theoretical low-mass end $m_{low}=0.08\\,\\ms$, the masses of Clusters\\,1 and 2 turn out to be $382\\,\\ms$\\ and $614\\,\\ms$, respectively.  The present study suggests that more optical open clusters are yet to be found in the Galaxy."
    },
    "0404/astro-ph0404479_arXiv.txt": {
        "abstract": "Results of simultaneous {\\it INTEGRAL} and optical observations of galactic microquasar SS433 in May 2003 are presented. The analysis of the X-ray and optical eclipse duration and hard X-ray spectra obtained by {\\it INTEGRAL} together with optical spectroscopy obtained on the 6-m telescope allows us to construct a model of SS433 as a massive X-ray binary. X-ray eclipse in hard X-rays has a depth of $\\sim 80\\%$ and extended wings. The optical spectroscopy allows us to identify the optical companion as a A5-A7 supergiant and to measure its radial velocity semi-amplitude $K_v=132$ km/s. A strong heating effect in the optical star atmosphere is discovered spectroscopically. The observed broadband  X-ray spectrum 2-100 keV can be described by emission from optically thin thermal plasma with $kT\\sim 15-20 keV$ ",
        "introduction": "SS433 is a supercritically accreting microquasar with precessing accretion disk and collimated mildly relativistic ($v\\approx 0.26 c$ ) jets (see e.g. Cherepashchuk et al. 1996 and references therein). Since the discovery in 1978, this unique massive X-ray binary system has been investigated in optical, radio and X-rays by many authors (e.g. Cherepashchuk 1981, Margon 1984, Kawai et al. 1989, Kotani et al. 1996, Goranskii et al. 1998a,b, Marshall et al 2002, Fabrika 2004).  SS433 is highly variable system and shows different types of periodicities:  1. Precessional variability ($P_{prec}=162^d.5$) which is observed by periodical Doppler-shifts of H, He I, Fe XXV and other X-ray emission lines in optical and X-ray spectra and is clearly visible in optical and X-ray light curves.  2. Orbital periodicity ($P_{orb}=13^d.082$). The shape of the orbital optical light curve strongly changes with precession phase (Goranskii et al. 1998 a,b, Cherepashchuk and Yarikov 1991). The value of the orbital period remains stable over more than 30 years, which can be considered as an argument against the common envelope model for SS433.  3. Nutation periodicity ($P_{nut}=6^d.2877$) which is observed as periodic deviations from purely precessional Doppler motion of emission lines and can be recovered from photometric data (Goranskii et al. 1988 a,b, Cherepashchuk and Yarikov 1991). The period and phase of the nutational variability remain stable over at least 16 years (around 950 nutation periods). Nutation radial velocity variations are delayed from the nutation photometric variability  by 0.6 days, which corresponds to the travel time from the accretion disk center to the formation region of optical moving emission lines downstream the relativistic jets located at $l\\simeq 10^{14}-10^{15}$ cm away from the disk center.  The AO-1 {\\it INTEGRAL} observations of SS433 discovered hard (up to 100 keV) X-ray spectrum in this supercritically accreting microquasar (Cherepashchuk et al. 2003), suggesting the presence of an extended hot (with temperature up to $10^8$ K) region at the central parts of the accretion disk. These new data made it possible to compare the eclipse characteristics of SS433 at different energies: soft X-rays (2-10 keV, the {\\it ASCA} data), soft and medium X-ray (1-35 keV, the {\\it Ginga} data), hard X-ray (20-70 keV, the {\\it INTEGRAL} data), and optical. Such a comparison allows us to study the innermost structure of the supercritical accretion disk and to constrain the basic parameters of the binary system. This is especially important in view of recent controversial results on the optical spectroscopy of SS433 (Gies et al. 2002, Charles et al. 2004). ",
        "conclusions": "The {\\it INTEGRAL} and optical observations of SS433 provide support that SS433 is a massive black hole X-ray binary system. The presence of broad wings in the observed hard X-ray eclipse profile evidences for an extended central hot region in the SS433 accretion disk. This region is observed to be substantially eclipsed during precessional motion of the accretion disk. The detected broad-band (2-100 keV) X-ray spectrum is best-fitteed by thermal plasma with $kT\\sim 15-20$~keV. The origin of such an extended hot region in the cenral parts of supercritically accretion disk in SS433 needs more studies.  The measured semi-amplitude of the radial velocity of the optical star is $K_v=132\\pm 9$ km/s, which confirms the results by Gies et al. (2002). However, the actual semi-amplitude should be decreased down to $\\sim 70$ km/s to be in agreement with the observed duration of X-ray and optical eclipses. This correction seems indeed inevitable because of strong heating effect present in the optical star of SS433.  To measure more precisely the binary parameters of SS433, further hard X-ray observations of SS433 are required at different precession phases and high-resolution optical spectroscopy need to be carried out at the maximum disk open angle to accurately measure $K_v$ and $K_x$. Analysis of X-ray observations at minimum X-ray flux and the cross-over precession phase will allow more detailed reconstruction of X-ray emitting zone in the center of the supercritically accretion disk in SS433."
    },
    "0404/astro-ph0404147_arXiv.txt": {
        "abstract": "We search for microvariability in a sample of 485 Mira variables with high quality $I$-band light curves from the second generation Optical Gravitational Lensing Experiment (OGLE-II). Rapid variations with amplitudes in the $\\sim$0.2--1.1 mag range lasting hours to days were discovered in Hipparcos data by de Laverny et al. (1998). Our search is primarily sensitive to events with time-scales of $\\sim$1 day, but retains a few percent efficiency (per object) for detecting unresolved microvariability events as short as 2 hours. We do not detect any candidate events. Assuming that the distribution of the event time profiles is identical to that from the Hipparcos light curves we derive the 95\\% confidence level upper limit of 0.038 yr$^{-1}$ star$^{-1}$ for the rate of such events (1 per 26 years per average object of the ensemble). The high event rates of the order of $\\sim$1 yr$^{-1}$ star$^{-1}$ implied by the Hipparcos study in the $H_P$ band are excluded with high confidence by the OGLE-II data in the $I$ band. Our non-detection could still be explained by much redder spectral response of the $I$ filter compared to the $H_P$ band or by population differences between the bulge and the solar neighborhood. In any case, the OGLE-II $I$-band data provide the first limit on the rate of the postulated microvariability events in Mira stars and offer new quantitative constraints on their properties. Similar limits are obtained for other pulse shapes and a range of the assumed time-scales and size-frequency distributions. ",
        "introduction": "\\label{sec:intro}  Mira variables are a subclass of Red Variables, also known as Long Period Variables (LPVs). They are low mass AGB giants, mostly with M type spectra (e.g. Gautschy \\& Saio 1996). Their large amplitude light variations on time scales of $\\sim$1 year are thought to arise due to the pulsation instability (e.g. Keeley 1970). There are intriguing reports of rapid light variations in Mira stars on time scales of a few days (Maffei \\& Tosti 1995) or even hours and minutes (Schaefer 1991) with amplitudes in the range 0.1--1.4 mag. The most systematic study of microvariability in Mira stars is based on Hipparcos data (de Laverny et al. 1998, hereafter dL98). It reports the discovery of 51 events in 39 out of 239 observed Mira stars. The amplitudes of events detected by Hipparcos are between 0.2 and 1.1 mag. To best of our knowledge, so far there are no published estimates for the rate of such events. Other reports of rapid variations in broad band photometry of specific Mira variables include Smak \\& Wing (1979) or more recently de Laverny et al. (1997) and Teets \\& Henson (2003). From existing data very little can be inferred about the spectral properties of those rapid variations, but microvariability in several spectral features has been reported (Odell et al. 1970, Kovar et al. 1972). On the other hand, Smith et al. (2002) do not detect any rapid variations in the COBE DIRBE mid-infrared light curves of 38 Miras.  Willson \\& Struck (2002) explored possible interpretations and favored one involving hot flashes due to interaction of the extended Mira atmosphere with a Jovian planet or a brown dwarf. This process was studied numerically by Struck, Cohanim \\& Willson (2002, 2004). The scenario can explain the time scale of the observed microvariability and produces energies and luminosities roughly of the right order of magnitude, but it ignores the fact that dL98 have seen approximately equal numbers of positive and negative flux excursions.  Here we use a large set of high quality $I$-band photometry from the Optical Gravitational Lensing Experiment (OGLE) to constrain the rate of microvariability for Mira stars found in the Galactic Bulge. ",
        "conclusions": "\\label{sec:discussion}  Using a large sample of objects with high quality light curves we place stringent limits on the frequency of rapid variability events in the Galactic Bulge Mira variables in the $I$ band. To the best of our knowledge this is the first published limit for the rate of rapid variability in Miras. The study is most sensitive to events lasting about 1 day or longer, however it retains a few percent efficiency (per star) for detecting unresolved microvariability on time scales as short as 2 hours. With 95\\% probability we excluded events of the type found by dL98 occurring more frequently than 1 per 26 years in the average object of our sample (0.038 yr$^{-1}$ star$^{-1}$). The results are not sensitive to the details of the temporal profile of the flares. The most significant dependence of the detection efficiency in our ensemble is on the event time scale because of the 1--3 day light curve sampling. The limits on the frequency of flashes with a sudden flux rise and exponential fall-off with power-law distribution of energies (at fixed parent star luminosity) have been presented. For flare durations ranging from 0.1 to 1 day the corresponding limiting rates are between once every few years and once every few tens of years.  In our calculations we essentially extrapolate the properties of the rapid variability events observed by Hipparcos from $H_P$ band to the $I$ band. Under that assumption, the OGLE-II data are inconsistent with high rates implied by the number of Hipparcos detections. A possible explanation is that the microvariability is much less prominent or the events are shorter in the $I$ band. In that case we are still learning something about microvariability in Mira stars. The typical accuracy of the OGLE-II photometry is much better than 0.2 mag (Figure~\\ref{fig:examples}) and we could not find any indication of rapid variations in any of the 485 studied objects. All light curves have very smooth appearance.  There are at least three possible explanations for the difference between the two results:  \\begin{enumerate}  \\item Mira variables in the Galactic Bulge population differ from their counterparts in the solar neighborhood. In the scenario of a planet interacting with the Mira atmosphere, e.g., the Bulge population might not have any planets. It is plausible that the Galactic Bulge Miras have lower metallicity than a typical Hipparcos star. Planets are more prevalent around metal-rich stars (e.g. Santos, Israelian, \\& Mayor 2004). It has been suggested that dust may play an important role in microvariability of Miras (Stencel et al. 2003). The dust content of the stellar envelope could be related to metallicity.  \\item Microvariability in Miras is occurring in the part of the spectrum that is not covered by Cousins $I$ band but is contributing to the blue part of the $H_P$ band.  \\item Microvariability in Miras is less common than implied by results of dL98. The only sure way to verify this is to collect more data and carefully analyze it. Numerous groups are setting up fast cadence photometric monitoring experiments with the goal of detecting real time transients (Vestrand et al. 2002) or finding transiting Jovian planets (Horne 2003). Teets \\& Henson (2003) are conducting a special purpose monitoring campaign of $\\sim$5 sources previously reported to display rapid variability using $B, V, R, I$ filters. The current data set covers many nights spanning the time interval of over 2 years with up to 5 photometric points per hour for up to 8 hours a night. So far only a single event in RR Boo has been detected. With the available phase coverage this is far below expectations based on the Hipparcos data.  \\end{enumerate}  The $H_P$ filter of the Hipparcos mission is broader than $V$ and still sensitive in the wavelength range 650--900 nm, where the spectral energy distribution of a Mira variable rises steeply. We can calculate the amount of energy emitted by a Mira star in the $H_P$ band knowing that the ratio of luminosities in $H_P$ and $V$ is about 2.5. This is easily estimated assuming that the spectrum of a Mira star blue-ward of 1000 nm is a black body with effective temperature $T_{eff} = 2500$ K and by folding it with filter transmission curves (Johnson 1965, van Leeuwen et al. 1997). Then using the absolute magnitudes $M_V = -0.3$ of an M5 giant and the $V$-band Solar luminosity $L_{V, \\odot} = 4.64\\times10^{32}$ erg s$^{-1}$ (Binney \\& Merrifield 1998) we obtain $L_{H_P} = 1.3\\times10^{35}$ erg s$^{-1}$ for a Mira star. We adopt $M_{V, \\odot}$ = 4.83. Mira variables vary considerably less in the $H_P$ band compared to $V$ with amplitudes rarely exceeding 4 mag (Whitelock, Marang \\& Feast 2000). Even during minimum light and taking an extreme value for the amplitude (4 mag) the total energy of a flare with amplitude $\\Delta m$ and duration $\\tau$ is:  \\begin{equation} \\label{eq:energy} E = 2.8\\times10^{38} {\\rm erg} \\times C \\times \\left({{\\tau} \\over {1 ~\\rm day}}\\right)\\times \\left(10^{0.4 |\\Delta m|} - 1\\right), \\end{equation}  \\noindent where $C$ is the numerical factor depending on the profile of the light variation. For a top hat function $C = 1$. We can see that a flare of a third of a magnitude lasting 24 hours will require $10^{38}$ ergs of energy---an order of magnitude more than the energy required for the same contrast in the $V$ band calculated by Willson \\& Struck (2002). If in fact the events detected by Hipparcos are confirmed, one will be faced with finding a large energy source that is triggered approximately as often as an energy sink of the same size. The observational evidence is accumulating and it should not be long before we can formulate a more definitive answer to the exciting question of rapid high amplitude variability on Miras."
    },
    "0404/astro-ph0404237_arXiv.txt": {
        "abstract": "We improve the algorithm of Kosowsky, Milosavljevic, and Jimenez (2002) for computing power spectra of the cosmic microwave background. The present algorithm computes not only the temperature power spectrum but also the E-mode polarization and the temperature-polarization cross power spectra, providing the accuracy required for current cosmological parameter estimation. We refine the optimum set of cosmological parameters for computing the power spectra as perturbations around a fiducial model, leading to an accuracy better than 0.5\\% for the temperature power spectrum throughout the region of parameter space within {\\sl WMAP}'s first-year 3$\\sigma$ confidence region. This accuracy is comparable to the difference between the widely-used CMBFAST code (Seljak and Zaldarriaga 1996) and Boltzmann codes. Our algorithm (CMBwarp) makes possible a full exploration of the likelihood region for eight cosmological parameters {\\it in about one hour on a laptop computer}. We provide the code to compute power spectra as well as the Markov Chain Monte Carlo algorithm for cosmological parameters estimation at {\\tt http://www.physics.upenn.edu/$\\sim$raulj/CMBwarp}. ",
        "introduction": "The recent Wilkinson Microwave Anisotropy Probe ({\\sl WMAP}) first-year data release has shown that, despite advances in our ability to compute the power spectrum of microwave background anisotropies and increases in computer power, cosmological parameter estimation is still relatively slow and computationally expensive. For example, for the simplest parameter space of flat, Lambda Cold Dark Matter (LCDM) models, the {\\sl WMAP} team needed about a week of supercomputer time (32 CPUs on an SGI Origin) to determine the two parameters joint 2-$\\sigma$ confidence contours in the parameter space \\cite{SpergelWMAP03}.  Allowing a non-flat geometry slows down the analysis significantly and more than quadruples the time required to draw the 2-$\\sigma$ contours. The main bottleneck in the parameter estimation process is the theoretical computation of the power spectrum anisotropy for a very large number of models, scattered throughout the parameter space. This is mostly done using publicly available codes such as CMBFAST (\\cite{sel96}) or CAMB (\\cite{CAMB}).  To accelerate the calculations, (DASH)\\cite{KKS02} introduced a shortcut in the calculation of the angular power spectra, based on analytical and semi-analytical approximations.  Rigorous analysis of future microwave background and other complementary data sets require recomputing the likelihood contours for cosmological parameters each time new data becomes available. Additionally for each new experiment, not one but many likelihood analyses are required to explore the impacts of potential systematic errors.  Not every researcher has access to months of supercomputer time for these computations; a fast and reasonably accurate method to compute confidence regions for cosmological parameters is needed. Motivated by the enormous computational burden that a proper 3 to 4-$\\sigma$ exploration of the likelihood hypersurface represents and by the related need to understand more clearly the physical parameters that control the shape of the power spectrum, Kosowsky, Milosavljevic and Jimenez (KMJ) \\cite{KMJ02} presented a new set of ``physical'' parameters formed from the usual cosmological parameters which have nearly-orthogonal effects on the microwave background temperature power spectrum. Thus, with this set of parameters, the efficiency of Monte Carlo techniques for evaluating the likelihood region increases significantly; these parameters have been used in the WMAP analysis \\cite{SpergelWMAP03,VerdeWMAP03} and then incorporated into Lewis' suite of microwave background analysis codes \\cite{lew02}. Subsequent work in a similar vein constructed a set of normal parameters that locally transformed the likelihood in parameter space into Gaussian form\\cite{CKK03}.  In the new KJM ``basis'', it is also possible to obtain an efficient approximation of the temperature power spectrum in a region around the fiducial model as a simple linear extrapolation from a fiducial model power spectrum. This new technique makes possible to compute approximate power spectra, several orders of magnitude faster than any previous methods, and was extensively used by the {\\sl WMAP} collaboration \\cite{VerdeWMAP03} in the preliminary analysis of the data.  Despite a significant increase in computational speed and physical insight afforded by the KMJ parameters, some remaining technical issues prevented a direct application of the KMJ power spectrum fits to real data.  In particular: (a) the parameters and the linear fits were only computed for the temperature power spectrum and not for the polarization; (b) the fiducial model used by KMJ had an optical depth to the last scattering surface which was too small ($\\tau=0.06$) compared with the {\\sl WMAP} best-fit model ($\\tau=0.17$); (c) for a fiducial model with higher $\\tau$ some covariance remained between the slope of the power spectrum $n$, the amplitude of the fluctuations $\\cal S$ and the optical depth parameter ${\\cal Z}= \\exp (-2 \\tau)$ which suggested that the KMJ parameters ${\\cal Z}$, ${\\cal S}$ and $n$ were not completely orthogonal; (d) errors in the power spectrum estimation became larger than 1\\% in some relevant regions of the parameter space.  Here we present a refined version of the KMJ algorithm, CMBwarp, that solves these problems. We extend the method to include polarization spectra, make all parameters nearly orthogonal, use a fiducial model given by the best fit model to {\\sl WMAPext} data\\cite{SpergelWMAP03}, and obtain the power spectra using polynomial fits for improved accuracy over linear extrapolations.  Throughout the region of parameter space within {\\sl WMAP}'s first-year data 3-$\\sigma$ confidence region\\footnote{3-$\\sigma$ for each parameter, marginalized over the other n-1 parameters.}, the accuracy is now comparable to the difference between CMBFAST and other independent Boltzmann codes.  We make CMBwarp publicly available from {\\tt http:/www.physics.upenn.edu/$\\sim$raulj/CMBwarp}; in addition, we also release accompanying Markov Chain code that can be used to perform parameter estimation. ",
        "conclusions": "\\label{sec:discussion}  The set of approximations described here, CMBwarp, comprise a highly accurate and extremely fast method of approximating the microwave background temperature and polarization power spectra in standard inflation-like cosmological models for the range of cosmological parameters which encompass the likelihood regions allowed by the {\\sl WMAP} first-year data. For 8 parameters, CMBwarp computes each multipole moment with $\\sim$ 50 floating-point operations; it is unlikely that any significantly more efficient computational scheme can be constructed.  Note that the computational effort is the same for all multipole moments independent of $\\ell$; conventional codes require large increases in computational cost as larger-$\\ell$ multipoles are computed. This fact will become more important in the coming years as experiments such as Planck, ACT \\cite{kosowsky03} and the South Pole Telescope map the microwave background at high resolution.  CMBwarp is particularly well-suited to explorations of cosmological parameter space using Markov Chain techniques. As cosmological data continues to improve, joint analysis combining many different data sets will continue. The technique allows easy incorporation of the true microwave background constraints into joint analysis with a minimum of computational expense. A useful extension of CMBwarp would be a comparably fast and accurate computation of the transfer function and matter power spectrum. The optimal set of ``physical'' parameters for describing the matter power spectrum will not necessarily be the same as for the microwave background; the crucial property of a set of parameters is that they must have orthogonal effects on the matter power spectrum. This will be addressed elsewhere. The speed and flexibility of CMBwarp make possible analysis with many more parameters which are not computationally feasible using standard techniques: one example is estimation of the primordial power spectrum using a number of parameters to describe it (see, e.g., \\cite{mukherjee03}), taking into account the effects of the standard cosmological parameters.  Markov Chain Monte Carlo explorations of cosmological parameter space have now become common. The need for very fast, accurate, and portable microwave background power spectrum computations in this context is readily apparent. We hope CMBwarp proves to be a useful tool for these analyses."
    },
    "0404/astro-ph0404284_arXiv.txt": {
        "abstract": "We investigate the relation between the asymptotic circular velocity, \\vcirc , and the central stellar velocity dispersion, \\sigc , in galaxies.  We consider a new sample of high surface brightness spiral galaxies (HSB), low surface brightness spiral galaxies (LSB), and elliptical galaxies with HI-based \\vcirc\\ measurements. We find that:\\\\ 1) elliptical galaxies with HI measurements fit well within the relation;\\\\ 2) a linear law can reproduce the data as well as a power law (used in previous works) even for  galaxies with \\sigc$< 70$ km/s;\\\\ 3) LSB galaxies, considered for the first time with this respect, seem to behave differently, showing either larger \\vcirc\\ values or smaller \\sigc\\ values. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404598_arXiv.txt": {
        "abstract": "Cluster mergings are expected to have an influence on the radio emission of the galaxy population. We present the results of a deep radio survey in the A3558 complex in the central region of the Shapley Concentration, in order to further explore our hypothesis of a dependence of the radio luminosity function on the age of the merging. ",
        "introduction": "During the huge energy release of a cluster merging, different mechanisms are at work on galaxies. On one side, the rapid changes in the tidal gravitational field can drive gas from outer to inner parts of galaxies, triggering star bursts and feeding the central AGN (\\cite{bekki}), while on the other side tidal stripping and ram pressure tend to devoid galaxies of their hydrogen.  \\cite{owen99} found that A2125, a merging cluster, presents a factor 4 more radiogalaxies than a relaxed one (A2645) at the same distance. In order to better study this phenomenon, we performed an extensive radio survey at 20cm with the ATCA telescope on the cluster complexes (\\cite{bardelli98a}, \\cite{bardelli98b}, \\cite{bardelli00}, \\cite{bardelli02}) located in the Shapley Concentration (\\cite{zucca93}).  The only way to estimate the radio activity of the clusters is to compute the bivariate radio-optical luminosity function (Figure 1), i.e. the number of radiogalaxies in the considered structure (as confirmed by their redshift) over the total number of optical galaxies. This takes into account the different richness of the clusters. We found that the A3558 complex, which is a merging in an advanced phase, presents a significant lack of radiosources (\\cite{venturi00}, \\cite{venturi01}, \\cite{venturi02}). On the contrary the A3528 complex, which is at an early phase of interaction, presents a number of radiosources comparable to that of normal clusters. Also the A3571 complex, a merging at the end of the process, does not have pecularities in the number of radio objects.  This result is only apparently in contraddiction with \\cite{owen99}, because one has to take into account the spiral richness of interacting clusters and/or the merging phase. Our proposed scenario is the following: at very early stages of merging, the outer parts of clusters (which are spiral rich) begin to interact, triggering star formation activity and therefore enhancing the number of radiosources. After this episode, the radiosource excess disappears and the clusters return quiescent (right panel of Figure 1), until the ellipticals (with their AGNs) in the inner part of the clusters are reached by the merging effect, leading to a lowering of the AGN emission: as a result, there is a deficit of bright radiosources (left panel of Figure 1). Finally, the virialization re-establishes the starting situation.  \\begin{figure} \\centering \\includegraphics[angle=0,width=0.5\\hsize]{zucca_iau_fig1.ps} \\includegraphics[angle=0,width=0.4\\hsize]{zucca_iau_fig2.ps} \\caption{Bivariate radio-optical luminosity functions of the A3558 complex (left) and of the A3528 complex (right). Triangles represent the reference luminosity function of normal clusters from \\cite{lo96}. The first complex is an advanced merger, while the second is a merging at early phases. For the A3528 complex, the radio luminosity function lies in the shadowed area.} \\end{figure}  \\begin{figure} \\centering \\includegraphics[angle=0,width=0.7\\hsize]{zucca_iau_fig3.ps} \\centering \\includegraphics[angle=0,width=0.5\\hsize]{zucca_iau_fig4.ps} \\caption{{\\it Upper panel:} circles represent the field of view of our VLA 20cm pointings, superimposed on the isodensities of the optical distribution of galaxies in the A3558 complex. {\\it Lower panel:} location of optical counterparts overplotted to the radiocontours. Triangles are late type galaxies, crosses are early type galaxies.} \\end{figure} ",
        "conclusions": ""
    },
    "0404/astro-ph0404551_arXiv.txt": {
        "abstract": "{We analyse spectroscopic observations of WR\\,20a revealing that this star is a massive early-type binary system with a most probable orbital period of $\\sim 3.675$\\,days. Our spectra indicate that both components are most likely of WN6ha or O3If$^*$/WN6ha spectral type. The orbital solution for a period of $3.675$\\,days yields extremely large minimum masses of $70.7 \\pm 4.0$ and $68.8 \\pm 3.8$\\,M$_{\\odot}$ for the two stars. These properties make WR\\,20a a cornerstone system for the study of massive star evolution. ",
        "introduction": "Although it is generally accepted that Wolf-Rayet (WR) stars are evolved descendants of massive O-stars, many open questions remain about the evolution of these objects. Among the most fundamental issues is the importance of mass transfer in the evolution of massive O + O binaries (e.g.\\ Vanbeveren et al.\\ \\cite{Vb}). It is therefore of prime interest to achieve as complete a census of WR binaries as possible and to investigate the fundamental properties of individual massive binary systems.  Among the 227 objects listed in the VIIth Catalogue of Galactic Wolf-Rayet Stars (van der Hucht \\cite{vdH}), there are only 52 known spectroscopic binaries and only 32 of them have reliable orbital solutions. In addition to these confirmed binaries, van der Hucht (\\cite{vdH}) provides a list of 30 rather faint and poorly studied objects that are potential WR + OB binaries. Most of these binary candidates were classified as such because of their comparatively weak emission lines that could be the result of the dilution by the continuum light of an OB-type companion. In order to search for a composite spectrum - and hence confirm multiplicity - we performed a snapshot spectroscopic survey of a sample of WR binary candidates. One of the most interesting objects that emerged from our survey was WR\\,20a.  \\object{WR\\,20a} (= THA\\,35-{\\sc ii}-036 = SMSP2) was catalogued by Th\\'e (\\cite{The}) in an H$\\alpha$ emission line survey, and first discovered to be a WR star by Shara et al.\\ (\\cite{Shara}) in the course of a deep narrow and broadband photometric survey for galactic Wolf-Rayet stars in the southern Milky Way. The latter authors also obtained a low-resolution (2 pixel resolution of 7\\,\\AA) spectrum of the star and proposed a WN7 spectral type, although they noted the lack of a He\\,{\\sc ii} $\\lambda$\\,5411 emission line. Shara et al.\\ (\\cite{Shara}) suggested that WR\\,20a had a rather high H/He abundance ratio. Later, the object was re-classified as WN7:h/WC by Shara et al.\\ (\\cite{Shara2}). The star is believed to be a member of the young open cluster Westerlund\\,2 in the core of the H\\,{\\sc ii} region RCW\\,49 (Moffat et al.\\ \\cite{Moffat}).  \\begin{figure*} \\begin{center} \\resizebox{15cm}{9.7cm}{\\includegraphics{Gc251_f1.ps}} \\end{center} \\caption{Spectrum of WR\\,20a as observed on HJD\\,2\\,453\\,034.715 (violet setting, top panel) and 2\\,453\\,034.741 (blue setting, lower panel). Note the double N\\,{\\sc iv} $\\lambda$\\,4058 emission line as well as the double N\\,{\\sc v} $\\lambda$\\,4604 and N\\,{\\sc v} $\\lambda$\\,4620 absorptions. The stellar lines are labelled above the spectrum, whereas the interstellar lines and diffuse interstellar bands (DIBs) are indicated below.\\label{spec}} \\end{figure*} In this letter, we first present our blue/violet observations of WR\\,20a (Sect.\\,\\ref{obs}). The spectral features as well as the nature of WR\\,20a are revisited in Sect.\\,\\ref{spect}, whilst Sect.\\,\\ref{solorb} deals with an analysis of the radial velocities. Finally, we briefly discuss the implications of our results in Sect.\\,\\ref{discuss}. ",
        "conclusions": "} The WN6ha or O3\\,If$^*$/WN6ha spectral type of the components of WR\\,20a is significantly earlier than those of the earliest members (O6:V -- O7:V) of Westerlund\\,2 analysed by Moffat et al.\\ (\\cite{Moffat}).  Several authors have estimated the distance of Westerlund\\,2: $7.9^{+1.2}_{-1.0}$\\,kpc (Moffat et al.\\ \\cite{Moffat}); $5.7 \\pm 0.3$\\,kpc (Piatti et al.\\ \\cite{Piatti});  $6.4 \\pm 0.4$\\,kpc (Carraro \\& Munari \\cite{CM}). Adopting a visual brightness ratio of 1 together with the photometry of Moffat et al.\\ (\\cite{Moffat}; $V = 13.58$, $B - V = 1.51$) the various distance estimates yield absolute magnitudes $M_{\\mathrm V}$ in the range $-5.0$ to $-6.1$. This is fainter than the average of three WN6ha stars in NGC\\,3603, namely $-6.9$ to $-7.7$\\,mag\\footnote{Moffat et al.\\ (\\cite{MDS}) derived an average $M_{\\mathrm V}$ of $-7.7$ for the three H-rich WN stars in NGC\\,3603, whilst Crowther \\& Dessart (\\cite{CD}) re-evaluated their distance and reddening and obtained an average of $-6.9$\\,mag.}, such that only the largest published distance of Westerlund\\,2 would be in marginal agreement with the photometry of WR\\,20a, allowing for a cosmic scatter of 0.5 -- 0.7\\,mag on $M_{\\mathrm V}$. This problem would be less severe if the absolute magnitudes of the components of WR\\,20a were more typical of O3\\,I stars ($M_{\\mathrm V} = -5.9$ according to Crowther \\& Dessart \\cite{CD}). Alternatively, WR\\,20a could actually be unrelated to Westerlund\\,2 (though the rarity of early-type stars and the angular proximity of WR\\,20a to the cluster strongly suggest membership) or the photometry of Moffat et al.\\ could be affected by photometric eclipses. Concerning the last possibility, we note that given the extremely large minimum masses of the components in WR\\,20a, it appears very likely that the system could display eclipses. A photometric monitoring of this binary is therefore of the utmost importance in order to help derive its absolute magnitude, to get rid of the aliasing problem of the RV data and to constrain the orbital inclination. We intend to organize such a campaign in the near future.  An important issue in the context of the formation and the evolution of very massive stars is that of the mass of the most massive ones. Masses up to $\\sim 200$\\,M$_{\\odot}$ have been inferred for the most luminous O2 stars from a comparison with evolutionary tracks (Walborn et al.\\ \\cite{WHL02}). However, to date, none of the O2 stars is known to belong to a binary system that would allow to determine its mass in a less model dependent way. The most massive main-sequence stars in binary systems known so far were found in R\\,136-38 (O3\\,V + O6\\,V) and R\\,136-42 (O3\\,V + O3\\,V). Massey et al.\\ (\\cite{MPV}) determined masses of 57 and 40\\,M$_{\\odot}$ respectively for the O3\\,V primary stars of these systems. Another extremely massive object is the WN7ha primary in the 80-day period binary WR\\,22 (Rauw et al.\\ \\cite{RV96}; Schweickhardt et al.\\ \\cite{SSS99}) with a minimum mass of 55\\,M$_{\\odot}$. WR\\,20a thus probably consists of two of the most massive stars with a direct mass determination known so far. Therefore, WR\\,20a is certainly a cornerstone system for future investigations of massive star evolution.  Finally, we note that the emission lines in the spectrum of WR\\,20a display strong phase-locked profile variability. This aspect will be discussed along with an analysis of the spectrum with a model atmosphere code in a forthcoming paper.  \\acknowledgement{The Li\\`ege team is greatly indebted to the Fonds National de la Recherche Scientifique (Belgium) for multiple assistance. This research is also supported in part by contract P5/36 ``P\\^ole d'Attraction Interuniversitaire'' (Belgian Federal Science Policy Office) and through the PRODEX XMM-OM and INTEGRAL projects. The SIMBAD database has been consulted for the bibliography.}"
    },
    "0404/astro-ph0404417_arXiv.txt": {
        "abstract": "{We  further develop  the Tayler--Spruit dynamo theory,  based on the most efficient instability for  generating magnetic fields in  radiative layers of differentially rotating stars. We avoid the simplifying assumptions that either the $\\mu$-- or the $T$--gradient  dominates, but we treat the general case and we also account for the nonadiabatic effects, which favour the growth of the magnetic field. The general equation leads to the same  analytical solutions  in the limiting  cases considered by Spruit (\\cite{Spruit02}). Numerical models of a 15 M$_{\\odot}$ star with a magnetic field are performed. The differences between the asymptotic solutions and the general solution demonstrate the need to use the general solution. Stars with a  magnetic field rotate almost as a solid body. Several of their properties (size of the core,  MS lifetimes, tracks, abundances) are closer to those of models without rotation  than with rotation only. In particular, the observed  N/C or N/H excesses in OB stars are better explained by our previous models with rotation only than by the present models with magnetic fields that predict no nitrogen excesses.  We show that there is a complex feedback loop between the magnetic instability and the thermal instability driving meridional circulation. Equilibrium of the loop, with a small amount of differential rotation, can be  reached when the velocity $U_{\\mathrm{magn}}$ of the growth of the magnetic instability is of the same order as the velocity $U_{\\mathrm{circ}}$ of the meridional circulation. This opens the possibility for further magnetic models, but at this stage we do not know the  relative importance of the magnetic fields due to the Tayler instability in stellar interiors.      \\keywords stars: rotation -- stars: magnetic field -- stars: evolution } ",
        "introduction": "A major concern is to know whether magnetic fields are important for stellar evolution or not (cf.  Roxburgh \\cite{Rox03}). Some initial numerical tests have suggested that the magnetic fields created by the  Tayler--Spruit  dynamo (Spruit \\cite{Spruit99}, \\cite{Spruit02}) in rotating stars may be large enough  to transport angular momentum in stellar interiors as well as chemical elements in upper MS stars (Maeder \\& Meynet \\cite{Magn1}). In this work, we perform further theoretical developments of the dynamo process, as well as some numerical models,  to study the consequences of stellar evolution with a magnetic field.    In the nice theory of  magnetic instabilities recently developed by Spruit (\\cite{Spruit02}), only some important limiting cases have been considered. In particular, the so--called Cases 0 and 1 have been studied. In Case 0, the $\\mu$--gradient dominates over the thermal gradient in stellar interiors, i.e., in terms of the partial  Brunt--V\\\"{a}is\\\"{a}l\\\"{a} frequency we have $N_{\\mu} > N_{\\mathrm{T}}$. Case 1 applies when the thermal gradient is the main restoring force, i.e., when $N_{\\mu} < N_{\\mathrm{T}}$. In this  case, the radiative losses from the magnetic instability are also accounted for; these losses reduce the restoring force of buoyancy  which acts against the magnetic instability.  The problem with the above description is that, in the  interiors of models no longer on the zero--age sequence, there is  at some moderate distance of the core a significant zone with $N_{\\mu} < N_{\\mathrm{T}}$ (covering about 25 \\% of the stellar radius), which is in a situation where the current  simplifications made in Case 1 are not appropriate. In this zone (we called it zone 1P), the magnetic instability is adiabatic, while the hypothesis of large non-adiabaticity is made (see Sect.~2 below). It also happens in all stars   that there are zones where  $N_{\\mu}$  and $N_{\\mathrm{T}}$ are of the same order and where it is better to account for both gradients, even if in general such zones are rather small. We consider it preferable to  consistently treat the general case and to account  properly everywhere for the $T$-- and $\\mu$--gradients, as well as for the non--adiabatic effects.  Also, the two  approximations mentioned above have very different functional dependences with respect to stellar parameters, such as  the angular velocity $\\Omega$, the differential rotation $q= -\\frac{\\partial \\ln \\Omega}{\\partial \\ln r}$ or the thermal and chemical gradient. The use of asymptotic solutions such as Spruit's (\\cite{Spruit02}) introduces  artificial jumps in the   behaviour of the diffusion coefficients.  Our purpose here is to express  the theory in a unique and consistent formulation for all  possible cases of  $N_{\\mu}$ and  $N_{\\mathrm{T}}$ and of departures from adiabaticity. In Sect. 2, we discuss the non--adiabaticity as a function of radiative and magnetic diffusivities $K$ and $\\eta$ respectively. In Sect. 3, we establish the general equations for the magnetic field  and  we check for  consistency with the expressions derived by Spruit (\\cite{Spruit02}). In Sect. 4 we examine the general expressions for the transport of the angular momentum and of the chemical elements. Sect. 5 shows some initial numerical applications  as an example. In Sect.~6 we examine the reciprocal feedback between meridional circulation and Tayler instability and the fact that these two processes lead to very different amounts of differential rotation. The conclusions are given in Section 7. ",
        "conclusions": "The equations for the general case of the Tayler--Spruit dynamo with $\\mu$--gradients and nonadiabatic effects have been developed. The general expressions  consistently allow us to recover Spruit's asymptotic relations (\\cite{Spruit02}), when  $N_{\\mu}$ or $N_{\\mathrm{T}}$ are zero.  Numerically the general and asymptotic results are consistent with some   quantitative differences. This demonstrates the need to use the general solutions, in particular if we want to perform observational tests.  Some initial results on the role  of possible magnetic field in stellar models have been found. The Tayler--Spruit dynamo imposes that the stars to rotate nearly as a solid body. This leads to a different evolution of rotational velocities as stars move away from the zero--age sequence. In general, the rotation velocities are  higher when a magnetic field is present. A second general result is that internal mixing and surface enrichments in products of the CNO burning  are smaller when a magnetic field is present than in rotating models without a magnetic field. Since spectroscopic observations of OB stars show significant N-- and He--enrichments, this  cast some doubt as to whether magnetic fields have a dominant effect, especially more than the models which have only rotation and no magnetic field show a good agreement with observations.  If magnetic fields are present, we may expect, as shown in Sect.~5, that the magnetic instability and the thermal instability which drives meridional circulation reach some balance in their complex feedback loop. For reasons explained in Sect.~5, this balance should be characterized by less differential rotation than in current models with only rotation  and by more differential rotation than in the present  models with a magnetic field. Such models may lead to more mixing than the present ones.   Finally, we emphasize that  at this stage we really do not know how the Tayler--Spruit instability interacts with the meridional circulation and how  the energy constraints may modify the results. But we have explored some consequences of the magnetic field in the absence of  meridional circulation and examined the possibilities  for its existence."
    },
    "0404/hep-th0404216_arXiv.txt": {
        "abstract": "The intent of this letter is to point out that the accretion of a ghost condensate by black holes could be extremely efficient. We analyze steady-state spherically symmetric flows of the ghost fluid in the gravitational field of a Schwarzschild black hole and calculate the accretion rate. Unlike minimally coupled scalar field or quintessence, the accretion rate is set not by the cosmological energy density of the field, but by the energy scale of the ghost condensate theory. If hydrodynamical flow is established, it could be as high as tenth of a solar mass per second for 10MeV-scale ghost condensate accreting onto a stellar-sized black hole, which puts serious constraints on the parameters of the ghost condensate model. ",
        "introduction": "Prompted by an increasingly precise experimental measurements of cosmological parameters, and in particular detection of acceleration of the universe due to an unknown source which looks like a cosmological constant, in the recent years there has been a wide discussion in the literature about modifications of Einstein gravity on cosmological scales as a possible alternative to dark matter and/or energy. However, finding a self-consistent and well-motivated theory which agrees with all the observations is proving to be quite a challenge.  Recently, Arkani-Hamed {\\it et.\\ al.}\\ proposed a model \\cite{Arkani-Hamed:2003uy}, dubbed {\\it a ghost condensation}, which they argued to be consistent with all experimental observations and provide an interesting modification of gravity in the infrared, with potential applications to inflation \\cite{Arkani-Hamed:2003uz}, dark matter and cosmological constant problems. It involves an introduction of a scalar field which develops a non-zero expectation value of its (timelike) gradient in vacuum, due to non-trivial kinetic term in the action. Such modifications of the scalar field kinetic term were considered earlier on phenomenological grounds in the model known as $k$-inflation \\cite{Armendariz-Picon:1999rj, Garriga:1999vw}.  The ghost condensate model has been studied on perturbative level in effective field theory \\cite{Arkani-Hamed:2003uy}, which already leads to interesting consequences such as star trails \\cite{Dubovsky:2004qe, Peloso:2004ut}. We look at the ghost condensate from a slightly different perspective, namely we would like to investigate its behaviour in the strong gravitational field, for instance, near a Schwarzschild black hole \\begin{equation}\\label{eq:metric} ds^2 = - f(r) dt^2 + \\frac{dr^2}{f(r)} + r^2 d\\Omega_n^2, \\end{equation} where $f(r) = 1 - r_g/r$, and $r_g = 2Gm$ is a gravitational radius of Schwarzschild black hole of mass $m$. The problem is similar to interaction of a cosmological scalar field with a black hole \\cite{Frolov:2002va}, so one would expect ghost condensate to be accreted by a black hole.  Accretion of fluid onto a black hole has long been an important problem in astrophysics \\cite{Bondi:1952ni}. Spherically symmetric steady-state fluid accretion onto a Schwarzschild black hole was derived in Ref.~\\cite{Michel:1972}. Minimally coupled scalar field \\cite{Jacobson:1999vr} and quintessence \\cite{Frolov:2002va, Bean:2002kx} accrete onto black holes as well, although the accretion rate is limited by the cosmological density of the field \\cite{Frolov:2002va}. Accretion of exotic matter fields can lead to unusual results. For instance, accretion of a phantom energy (which violates energy dominance conditions) {\\it decreases} the black hole size \\cite{Babichev:2004yx}.  In this letter, we calculate the steady-state accretion rate of the ghost condensate by a black hole, and point out that it could be extremely efficient. This puts serious constraints on the parameters of the ghost condensate model. ",
        "conclusions": "In the last section, we calculated the steady-state accretion rate of the ghost condensate by a black hole for spherically symmetric flows. The most important result of the calculation is that the dimensionless coefficient $\\alpha$, which determines the accretion rate of the flow, is bounded below by $1$ even as density of the ghost condensate far from a black hole becomes vanishingly small. This means that it is the energy scale $M$ of the ghost condensate theory that sets the accretion rate, and not the cosmological abundance of the ghost condensate field as one might have naively expected.  The physical reason for this is that unlike the usual dark matter, the homogeneous ghost condensate does not have angular momentum and can be prevented from falling onto a gravitating body only by building up pressure support, which requires significant displacement from the minimum and densities of order $M^4$. This consideration alone might pose significant problems for ghost condensate as a realistic dark matter candidate, as dark matter galaxy halos are thought to be {\\em virialized} and not pressure-supported. This seems impossible to achieve in the framework of the ghost condensate model, unless the condensate becomes strongly non-homogeneous and breaks up into ``particles'', in which case it is hard to view it in any sense as a modification of gravity.  Up to a numerical coefficient of order one, the steady-state accretion rate is equal to the energy density $M^4$ falling down through the horizon area $4\\pi r_g^2$ at the speed of light. The top value of 10MeV for the ghost energy scale quoted in \\cite{Arkani-Hamed:2003uy} corresponds to a rather high density \\begin{equation} (10\\text{MeV})^4 = \\frac{(10\\text{MV} \\cdot e)^4}{\\hbar^3 c^5} = 2.32 \\cdot 10^{12} \\frac{\\text{kg}}{\\text{m}^3}. \\end{equation} If the steady-state flow of the kind we considered is established, the accretion rate of a 10MeV-scale ghost condensate by an astrophysical black hole would be enormous \\begin{equation} \\dot{m} = 0.08\\,\\alpha\\, \\frac{M_\\odot}{\\text{s}} \\left(\\frac{r_g}{3\\text{km}}\\right)^2 \\left(\\frac{M}{10\\text{MeV}}\\right)^4. \\end{equation} To avoid rapid black hole growth and its astrophysical consequences, energy scale $M$ of the ghost condensate should be significantly less than 10MeV. Stellar-size black hole would double in size over the lifetime of the universe (roughly 14Gyrs, or $4\\cdot10^{17}$s) for the ghost energy scale of order 1keV. This estimate goes down to 10eV for super-massive ($10^9 M_\\odot$) black holes.  In our calculations, we ignored the backreaction of the accreting ghost condensate matter onto a black hole metric. This is a good approximation for low accretion rates and ghost energy scales. However, if the accretion rate becomes as large as the above estimate, backreaction can no longer be ignored, both because the density of the ghost condensate near black hole is high and a large supply of ghost matter far from the black hole is needed to sustain the flow. Although proper treatment of backreaction is unlikely to alleviate the problem of excessive accretion rates, as it is caused by the largeness of the effect in the first place, sustainability of the steady-state flow with such enormous accretion rates is questionable. It would seem more likely that the black holes would completely evacuate all ghost matter from their gravitational potential wells, growing in the process.  Spherically symmetric steady-state accretion is an idealized situation, of course. We have not considered time-scale required to establish such flows, the role of initial conditions, motion of the black hole with respect to the condensate, or what happens if the ghost field becomes highly inhomogeneous. All of these are much harder problems, and it might turn out that some factors prevent the accretion from settling into an efficient steady-state regime. Still, having ghost condensate capable of such high accretion rates is alarming, and the issue should be further addressed by the ghost condensate scenario."
    },
    "0404/astro-ph0404485_arXiv.txt": {
        "abstract": "We reanalyze the formation and evolution of galaxies in the hierarchical clustering scenario.  Using a semi-analytic model (SAM) of galaxy formation described in this paper, which we hereafter call the Mitaka model, we extensively investigate the observed scaling relations of galaxies among photometric, kinematic, structural and chemical characteristics.  In such a scenario, spheroidal galaxies are assumed to be formed by major merger and subsequent starburst, in contrast to the traditional scenario of monolithic cloud collapse.  As a new ingredient of SAMs, we introduce the effects of dynamical response to supernova-induced gas removal on size and velocity dispersion, which play an important role on dwarf galaxy formation.  In previous theoretical studies of dwarf galaxies based on the monolithic cloud collapse given by Yoshii \\& Arimoto and Dekel \\& Silk, the dynamical response was treated in the extremes of a purely baryonic cloud and a baryonic cloud fully supported by surrounding dark matter.  To improve this simple treatment, in our previous paper, we formulated the dynamical response in more realistic, intermediate situations between the above extremes.  While the effects of dynamical response depend on the mass fraction of removed gas from a galaxy, how much amount of the gas remains just after major merger depends on the star formation history.  A variety of star formation histories are generated through the Monte Carlo realization of merging histories of dark halos, and it is found that our SAM naturally makes a wide variety of dwarf galaxies and their dispersed characteristics as observed.  It is also found that our result strongly depends on the adopted redshift dependence of star formation timescale, because it determines the gas fraction in high-redshift galaxies for which major mergers frequently occur.  We test four star formation models. The first model has a constant timescale of star formation independent of redshift.  The last model has a timescale proportional to the dynamical timescale of the galactic disk.  The other models have timescales intermediates of these two. The last model fails to reproduce observations, because it predicts only little amount of the leftover gas at major mergers, therefore giving too weak dynamical response on size and velocity dispersion of dwarf spheroidals.  The models, having a constant timescale of star formation or a timescale very weakly dependent on redshift, associated with our SAM, succeed to reproduce most observations from giant to dwarf galaxies, except that the adopted strong supernova feedback in this paper does not fully explain the color--magnitude relation under the cluster environment and the Tully-Fisher relation.  A direction of overcoming this remaining problem is also discussed. ",
        "introduction": "Dwarf galaxies give us many useful insights on the galaxy formation. They occupy a dominant fraction among galaxies in number and have a wide variety of structures and formation histories.  Furthermore it would be easier to understand how physical processes such as heating by supernova explosions affect their evolution, because of shallowness of their gravitational potential well, compared to massive galaxies that have more complex formation histories and physical processes. Equally important is that massive galaxies have evolved via continuous mergers and accretion of less massive galaxies, according to the recent standard scenario of large-scale structure formation in the universe with the cold dark matter (CDM). This indicates that it is essential to understand the formation of dwarf galaxies even in understanding the formation of massive galaxies, because they would be formed by mergers of their {\\it building blocks} like dwarf galaxies.  Before the CDM model gained its popularity, traditional models of monolithic cloud collapse such as the galactic wind model for elliptical galaxies \\citep{els62,l69,i77,s79,ys79,ay86,ay87,ka97} and the infall model for spiral galaxies \\citep{ayt91} had been widely used in analyses of galaxy evolution.  These are simple but strong tools to explain various observations and then have contributed to construction of a basic picture of galaxy evolution.  Based on such traditional models, \\citet{ds86} systematically investigated many aspects of dwarf galaxies. Especially they focused on the role of supernova (SN) feedback, that is, heating up and sweeping out of the galactic gas by multiple SN explosions, in the processes of formation and evolution of dwarf galaxies.  Since a large amount of heated gas is expelled by the galactic wind, the self-gravitating system expands as a result of dynamical response to the gas removal \\citep{h80, m83, v86}.  They considered two limiting cases for the gas removal.  One is a purely self-gravitating gas cloud, and another is a gas cloud embedded in a dominant dark halo.  They showed that many observed scaling relations are well explained by their model.  Combined with the evolutionary population synthesis code, \\citet{ya87} extended their analysis to estimating directly observable photometric properties, while only considering the purely self-gravitating gas cloud.  Since the hierarchical clustering of dark halos predicted by the CDM model becomes a standard structure formation scenario, the galaxy formation scenario must be modified so as to be consistent with the hierarchical merging of dark halos.  Explicitly taking into account the Monte Carlo realization of merging histories of dark halos based on the distribution function of initial density fluctuations, the so-called semi-analytic models (SAMs) of galaxy formation have been developed \\citep[e.g.,][]{kwg93, c94, bcf96, ngs99, sp99, spf01, ntgy01, nytg02}. SAMs include several important physical processes such as star formation, supernova feedback, galaxy merger, population synthesis and so on.  While SAMs well reproduce many observed properties of galaxies, most of analyses have been limited to those of massive galaxies.  Thus we focus on the formation of dwarf galaxies in the framework of the SAMs in this paper.  In the analyses of dwarf galaxies, the dynamical response to gas removal, which is known to play an important role as shown by \\citet{ds86} and \\citet{ya87}, must be taken into account. While self-consistent equilibrium models taking into account baryon and dark matter have been analyzed by \\citet{ys87} and \\citet{cp92}, we focus on the dynamical response of baryons within a dark matter halo. Since we have derived the mathematical formula of dynamical response for the galaxies consisting of baryon and dark matter in \\citet{ny03}, it is possible to incorporate them into our SAM.   The main purpose of this paper is a reanalysis of \\citet{ds86} and \\citet{ya87} in the framework of SAMs.  Thus we mainly focus on the formation of elliptical galaxies, which are assumed to be formed by major merger and subsequent starburst.  In order to do this, we construct the Mitaka model, which is a SAM including the effects of dynamical response.  So far many scaling relations among photometric, structural, and kinematical parameters of elliptical galaxies have been observed, such as the color-magnitude relation \\citep[e.g.,][]{b59}, the velocity dispersion-magnitude relation \\citep{fj76}, and the surface brightness-size relation \\citep{k77, kow83}.  To understand how these relations are originated, especially on scales of dwarf galaxies, helps to clarify the processes of galaxy formation and evolution in the context of the cosmological structure formation scenario. Although some authors found by principal component analysis that elliptical galaxies are distributed over the so-called fundamental plane in the three-dimensional space among photometric, structural and kinematical parameters \\citep{wko85,d87,dd87}, we focus on the direct observables in this paper rather than their principal component projection.   This paper is outlined as follows.  In \\S2 we describe our SAM.  In \\S3 we constrain model parameters in our SAM using local observations.  In \\S4 we compare the theoretical predictions of SAM galaxies with various observations.  In \\S5 we examine consistency check of our SAM with other observations.  In \\S6 we show the cosmic star formation history.  In \\S7 we discuss the nature of galaxies on the cooling diagram that has been traditionally used for understanding the formation of galaxies.  In \\S8 we provide summary and conclusion. ",
        "conclusions": "We have investigated the formation and evolution of galaxies in the context of the hierarchical clustering scenario by using the Mitaka model, or our SAM in which the effects of dynamical response on size and velocity dispersion of galaxies are explicitly taken into account, according to the formula by \\citet{ny03} for galaxies consisting of baryon and dark matter.  This paper is therefore an extension of previous analyses in the context of monolithic cloud collapse scenario \\citep{ds86, ya87}.  A $\\Lambda$-dominated flat universe, which is recently recognized as a standard, is exclusively used here.   The investigation mainly focuses on elliptical galaxies which are assumed to be formed by major merger and starburst-induced gas removal followed by dynamical response of the systems.  While a mass fraction of removed cold gas, after heated, is determined by the circular velocity of galaxies similar to the traditional collapse models, the total amount of cold gas depends on the formation history of galaxies which is realized by the Monte Carlo method based on the power spectrum of density fluctuation predicted by the CDM model.  The Mitaka model, which we have constructed in this paper, is found to reproduce a wide variety of observed characteristics of galaxies, particularly their scaling relations among various observables such as magnitude, surface brightness, size, velocity dispersion, mass-to-light ratio and metallicity.  This strongly supports the CDM cosmology and merger hypothesis of elliptical galaxy formation even on scales of dwarf galaxies.  Most of model parameters related to star formation, SN feedback and galaxy merger are constrained by the local luminosity function and mass fraction of cold gas in spiral galaxies.  As an extra parameter to be furthermore constrained, we have examined redshift dependence of SF timescale through comparison among the four SF models denoted by CSF, DSF0, DSF1, and DSF2.  The CSF refers to a constant SF timescale against redshift and has been suggested to be consistent with observations of galaxy number counts \\citep{ntgy01, nytg02}, quasar luminosity function \\citep{kh00, eng03}, and number and metallicity evolution of damped Ly-$\\alpha$ systems \\citep{spf01, ongy02}.  The DSF0 refers to a SF timescale simply proportional to dynamical timescale. The DSF1 and DSF2 are intermediate between the CSF and DSF0 (see Figure \\ref{fig:sfr}). It is found that the DSF0 model fails to reproduce observed properties of local dwarf spheroidals and galaxy counts.  Among the above four SF models, the CSF model has the longest SF timescale at high redshift and therefore the largest amount of cold gas at that epoch.  Accordingly, the fraction of removed gas during starburst is the largest in the CSF model, giving many extended galaxies in agreement with observation.  In this sense, the model of constant star formation (CFS) or at most mildly evolving star formation (DSF2) is favorable.  Note that {\\it dominant dark halo}, in which size and velocity dispersion do not change during gas removal as considered by \\citet{ds86}, apparently explains the observed scatter of size.  However, arguments from physical ground indicate that the dynamical response should play a significant role in forming galaxies with low velocity dispersion of $\\sigma_{0}(\\rm 1D)\\la 10$km s$^{-1}$, as actually observed in the Local Group.  There are some areas of further improvements in the current Mitaka model.  Because of strong dependence of SN feedback on circular velocity assuming $\\alpha_{\\rm hot}=4$, dwarf galaxies in our SF models are too faint to agree with expected magnitude from observed TFRs.  Furthermore, because of the same reason, dwarf galaxies are too blue to agree with expected colors from observed CMRs for cluster elliptical galaxies. On the other hand, such large value of $\\alpha_{\\rm hot}$ stops chemical and photometric evolution of elliptical galaxies at early epochs and is required to explain their observed low stellar metallicity as well as the galaxy number counts in the HDF and SDF.  In order to settle these contradictions, some new ingredients need to be introduced in the SAM analysis.  At least, for example, we must know how efficient the gas cooling is in massive dark halos and whether SN feedback at starburst works in the same way as that in disks.  We adopted a fitting mass function of dark halos by \\citet{yny03} for given redshift instead of the often used PS mass function.  This YNY mass function is slightly different from those by \\citet{st99} and \\citet{j01}, but is confirmed to provide a better fit to recent $N$-body results given by \\citet{yy01} and \\citet{y02}, which is discussed in more detail in a separate paper \\citep{yny03}.  Since the number density of dark halos affects SN feedback-related parameters to be chosen, it is very important to establish the mass function of dark halos, particularly on small mass scales, by high resolution $N$-body simulation.  Although such SAM analyses have just begun \\citep[e.g.,][]{bpfbj01, h03a, h03b}, further investigation will obviously be required.  We are upgrading the Mitaka model accommodated with full high resolution $N$-body simulations, and the results will be given elsewhere in the near future.  \\begin{figure} \\plotone{f29.eps} \\caption{Dynamical response of galaxy size to removed mass from the system.  The approximate results given by equation (A9) are shown for $D=0.05$ and 0.2 (solid lines).  For each value of $D$, the exact results given by equation (\\ref{eqn:full}) are shown for $z_{i}=0.5$ (dashed lines) and 0.05 (dot-dashed lines). } \\label{fig:approx} \\end{figure}"
    },
    "0404/astro-ph0404166_arXiv.txt": {
        "abstract": "We analyze the possible relationship between the current orbital elements fits of known exoplanets in the 2/1 mean-motion resonance and the expected orbital configuration due to migration. We find that, as long as the orbital decay was sufficiently slow to be approximated by an adiabatic process, all captured planets should be in apsidal corotations. In other words, they should show a simultaneous libration of both the resonant angle and the difference in longitudes of pericenter.  We present a complete set of corotational solutions for the 2/1 commensurability, including previously known solutions and new results. Comparisons with observed exoplanets show that current orbital fits of three known planetary systems in this resonance are either consistent with apsidal corotations ({\\it GJ876} and {\\it HD82943}) or correspond to bodies with uncertain orbits ({\\it HD160691}).  Finally, we discuss the applicability of these results as a test for the planetary migration hypothesis itself. If all future systems in this commensurability are found to be consistent with corotational solutions, then resonance capture of these bodies through planetary migration is a working hypothesis. Conversely, If any planetary pair is found in a different configuration, then either migration did not occur for those bodies, or it took a different form than currently believed. ",
        "introduction": "It is well known that extrasolar planets are not where we imagined. Classical planetary formation theories based on planetesimal accretion and core-instability for the giant planets predict bodies in quasi-circular orbits and semimajor axes $a$ far from the star. For solar type stars, the minimum semimajor axis is about 4 AU, which is the distance where non-rocky volatile elements can condense and accrete. However, many exoplanets do not follow this rule, but are found in highly eccentric orbits with $a < 1$ AU.  Two options have been proposed to explain this dilemma. In the first, it is assumed that present cosmogonic theories are in fault, or at least they followed different routes in practically all other planetary systems (thus making our own Solar System a very particular case). In the second, exoplanets really did form far from the central star, but suffered a posterior decay in their semimajor axes towards their present sites. Thus was born what is usually referred to as the ``Hypothesis of Planetary Migration''. However, in order for migration to be a real theory and not just a simplistic escapade, two conditions must be met: {\\it (i) } the existence of a plausible driving mechanism to explain the alleged decay in orbital energy, and {\\it (ii)} concrete evidence that exoplanets did undergo such an evolution.  Two different driving mechanism have been presented in the last few years. The first (Murray et al. 1998) is based on the interaction of the planets with a reminiscent planetesimal disk, and works in the same manner as migration of the giant planets in our Solar System (Fernandez and Ip 1984, Hahn and Malhotra 1999). However, this mechanism does not seem to be sufficiently efficient. First, it requires a very large disk mass, of the order of 0.1 $M_{\\odot}$, to explain an orbital decay of several astronomical units. Second, it is not completely evident that multi-planet systems should in fact undergo a simultaneous decrease in semimajor axis. Recall that in our system, Jupiter is believed to have suffered decrease in orbital energy, while Saturn, Uranus and Neptune have increased their value of $a$.  The second proposed mechanism is the interaction of the planets with the gaseous disk. The existence of disk torques cause a transfer of energy and angular momentum from the planet to the gas and, if the disk parameters are chosen correctly, then the solid body should undergo a negative migration (i.e. ${\\dot a}<0$). Several simulations have been performed in recent years (e.g. Snellgrove et al. 2001, Nelson and Papaloizou 2002, Kley 2003, Papaloizou 2003) and these seem to indicate that the mechanism works reasonably well. However, certain aspects of this process are also problematical. For some disk parameters the migration can be positive, leading to an increase in semimajor axis, which is just the opposite desired result. In some other cases, the orbital energy may even exhibit random-walks with no secular variation. Nevertheless, it seems that this mechanism is the most probable process to explain migration.  Having found a plausible process for the orbital decay, we must now search for evidence that this really occurred in the exoplanets. This question is particularly important, since such large-scale {\\it inward} migration did not happen in our case. A possible solution is to find a particular orbital characteristic of the extrasolar bodies, intimately related to migration, which can be used as (at least) indirect evidence of this process. Our Solar System presents two cases of confirmed migration (be it outward or inward): the giant planets and planetary satellites. As mentioned before, the outer planets migrated due to interaction with a remnant planetesimal disk, while many of the regular satellites of these same planets evolved due to tidal effects of the central mass. In this latter case, we know that an important consequence of the migration was capture of the satellites in exact mean-motion resonance (e.g. Colombo et al. 1974, Yoder 1979). A well known example is given by the Galilean satellites of Jupiter. It is known that these configurations cannot be explained solely with gravitational perturbations, but only through resonance trapping under the effects of an exterior non-conservative force. The case of our outer planets is different. It seems that they are not exactly in resonance due to certain random-walk characteristics of the driving mechanism itself (Hahn and Malhotra 1999). Thus, we can conclude that although migration does not always lead to resonance trapping, the existence of massive bodies in exact mean-motion resonance can only be explained via a migration mechanism.  The fact that our planets did not suffer a significant smooth inward migration is consistent with the fact that none of them are trapped in resonance. What about the exoplanets? A good evidence in favor of migration would then be to analyze whether planetary systems do show mean-motion resonant relations. Of the 13 presently known planetary systems, including both confirmed and un-confirmed cases, we restrict ourselves to those systems where the ratio in semimajor axes is sufficiently small to assure significant gravitational interaction between the bodies. Choosing this limit to be $a_2/a_1 = 3$, we find that 6 systems satisfy this condition. These are: {\\it GJ876}, {\\it 55Cnc}, {\\it 47Uma}, {\\it HD82943}, {\\it HD160691} and {\\it Ups And}. Of these, five are believed to be in the near vicinity of mean-motion resonances, while {\\it Ups And} is in a secular resonance.  Although this proportion is very significant, it must be considered with care. Recent data analysis (Mayor et al., private communication) shows that the published orbit of {\\it 47 Uma} may be questionable, and there is no general agreement if the system is near the 7/3 or 5/2 mean-motion resonances. Doubts also exist for {\\it HD160691}, and the second planet is not yet confirmed. In view of this debate, and considering the intrinsic errors in orbital fits, the mere proximity of these systems to mean-motion resonances is not evidence enough for migration, especially considering that there may be a natural tendency of researchers to place planets in commensurabilities even though they may not be near enough. For these reasons, we feel a more detail analysis is necessary.  This manuscript undertakes such a analysis. In Section 2 we present new results on the location and characteristics of corotations in the 2/1 resonance. Section 3 discusses the problem of planetary migration from the point of view of Adiabatic Invariant theory and we show that only corotational type configurations can apparently be expected in trapped planets. A comparison between these solutions and the current orbits of three exoplanetary systems are detailed in Section 4. Finally, conclusions close the paper in Section 5. ",
        "conclusions": "In this paper we have presented a new catalog of general corotational solutions for the 2/1 mean-motion resonance. Apart from the well-known aligned and anti-aligned solutions, we have also extended our knowledge of asymmetric configurations and have mapped the recently discovered $(\\pi,\\pi)$-corotations. Since these periodic orbits depend on the planetary masses only through the ratio $m_2/m_1$, and on the semimajor axes only by $a_1/a_2$, they are very general in nature and should be valid for any exoplanetary system showing two planets trapped in this commensurability.  The determination of the period of oscillation $\\tau_c$ around these fixed points of the averaged problem shows that any migration mechanism with characteristic timescale satisfying \\be \\label{eq13} {1 \\over a_i} {da_i \\over dt} \\gg {m_1 m_2 \\over a_i} \\tau_c \\ee should be adiabatic. Thus, starting from quasi-circular orbits, the evolutionary track of the planets inside the resonance should be well reproduced by the families of apsidal corotations for that particular mass ratio. This, together with the fitted values of the orbital eccentricities, allow us to stipulate whether present orbits are consistent with apsidal corotations and, consequently, with a planetary migration of the system from cosmogonic locations far from the star. In other words, we are able to suggest a simple test which relates the present orbits with restrictions on the properties of the formation process of resonant planets.  Application to the 2/1 resonance shows that, of the three published systems, {\\it GJ876} satisfies all our conditions satisfactorily. The old orbits for the {\\it HD82943} planets are not consistent with our hypothetical relation between corotation and migration. However, new orbital determinations for this system have yielded values very compatible with orbital decay, thus indicating that the problem is due to uncertainties in the fits themselves. Finally another problematic case, {\\it HD160691}, also seems to be biased by errors in the orbital fits. In fact, the exterior planet of this system is probably not existent at all.  In view of these examples, it seems that there is not indication of the existence of any planetary system in the 2/1 commensurability which is in specific contradiction with the hypothesis of planetary migration. Recent results by Zhou et al. (2004) for the {\\it 55 Cnc} system trapped in the 3/1 mean-motion resonance show similar results. It will be very interesting to extend this analysis to all future resonant systems and see whether an adiabatic migration mechanism continues to pass all the tests."
    },
    "0404/astro-ph0404350_arXiv.txt": {
        "abstract": "We study the effects of radiative cooling, star formation and stellar feedback on the properties and evolution of galaxy clusters using high-resolution Adaptive Mesh Refinement N-body+gasdynamics simulations of clusters forming in the $\\Lambda$CDM universe. Cooling leads to the condensation of gas in the inner regions of clusters, which in turn leads to steepening of the dark matter profile. The cooling gas is replaced by the higher-entropy gas from the outer regions, which raises the entropy and temperature of gas in the cluster core.  The magnitude of these effects is likely overestimated in the current simulations because they suffer from the overcooling problem: a much larger fraction of baryons is in the form of cold gas and stars than is observed.  We find that the thermal stellar feedback alone does not remedy this problem. Additional ad-hoc preheating can lower the amount of cold gas but a simple uniform preheating results in incorrect star formation history, as it delays the bulk of star formation until $z<1$.  Our analysis shows that the overcooling in a cluster as a whole is really the overcooling in the central galaxy and its progenitors at high redshifts.  This indicates that an additional heating mechanism that can continuously heat the gas in the cluster core is required to reproduce the observed cluster properties. Energy injection by the Active Galactic Nuclei, which may provide such heating, may thus be an important missing ingredient in the current theoretical models of cluster formation. ",
        "introduction": "Below, we highlight some of the important effects of gas cooling and star formation on the properties of galaxy clusters: \\vspace{2mm}  \\begin{itemize}  \\item Gas cooling induces condensation of gas in the cluster core region ($\\lesssim 100$kpc), and the baryon mass in the cluster core at $z=0$ is dominated by the stars associated with the central galaxy.  \\item The condensation of baryons causes contraction of dark matter and steepens the dark matter profile in the central region from $r^{-1.5}$ in the adiabatic runs to $r^{-2.0}$ in the simulations with cooling.  \\item Gas cooling induces inflow of high-entropy gas from the outside and raises the level of entropy and gas temperature in the cluster core region.  \\item Stellar feedback and/or additional preheating suppress the effects described above by regulating gas cooling. The magnitude of the effects on the distribution of dark matter, entropy and temperature of gas is, therefore, sensitive to the total fraction of baryons which cooled to form stars and cold gas clouds.  \\end{itemize}  \\begin{figure} \\centerline{ \\epsfysize=2.7in \\epsfbox{Figs/pro_fcold_2.ps} } \\vspace{-0.2cm} \\caption{ The cold gas + stellar mass fraction of the simulated cluster that includes cooling, star formation and stellar feedback.  Three lines indicate the CSF runs with varying star formation : the star formation is allowed to proceed only in the region where the local gas density is larger than 1.0 (Red), 0.1 (Blue) and 0.01cm$^{-3}$ (Green). The figure shows that 35-55\\% of baryons within the virial radius (indicated by the dashed line) are in the form of cold gas and stars at $z=0$, which is significantly larger than the observed value of 10-20\\%. } \\label{fcold} \\end{figure} ",
        "conclusions": "The overcooling problem is a serious challenge for the cluster formation simulations. Our analysis shows that the overcooling in a cluster as a whole is really the overcooling in its core. In other words, the central cluster galaxies formed in simulations are too massive and bright compared to observations.  Analysis of cluster evolution indicates that the problem sets in early at high redshifts, and is not limited to a single object or region. This indicates that an additional heating mechanism is necessary for reproducing the observed cluster properties. These results lend support to the idea that feedback from Active Galactic Nuclei is an important missing ingredient in the current simulations of cluster formation."
    },
    "0404/gr-qc0404067_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "At present there are only hints, but no compelling evidence for Lorentz violation from quantum gravity. Moreover, even if LV is present, the use of EFT for its low energy parametrization is not necessarily valid. Nevertheless, we believe that the constraints derived from the simple ideas discussed here are still important.  They allow tremendous advances in observational reach to be applied in a straightforward manner to limit reasonable possibilities that might arise from fundamental Planck scale physics. Such guidance is especially welcome for the field of quantum gravity, which until the past few years has had little connection with observed phenomena."
    },
    "0404/astro-ph0404216_arXiv.txt": {
        "abstract": "{ Pulsars with drifting subpulses are thought to be an important key to unlocking the mystery of how radio pulsars work.  We present new results from high sensitivity GMRT observations of \\OurPsr\\/ -- a wide profile pulsar that exhibits an interesting but complicated drifting pattern.  We provide a model to explain the observed subpulse drift properties of this pulsar, including the apparent reversals of the drift direction. In this model, \\OurPsr\\/ is close to being an aligned rotator.  Using information about the polarization and frequency evolution of the pulse profile, we solve for the emission geometry of this pulsar and show that the angle between the rotation axis and the dipole magnetic axis is less than 5\\degr. As a result, our line of sight samples a circular path that is entirely within the emission beam. We see evidence for as many as 6 to 7 drifting bands in the main pulse at 318 MHz, which are all part of a circulating system of about 15 spark-associated subpulse emission beams that form, upon averaging, one conal ring of the mean emission.  We also see evidence for a second ring of emission, which becomes dominant at higher frequencies (above 1 GHz) due to the nature of the emission geometry. We model the subpulse drift behaviour of this pulsar in detail, providing quantitative treatments of the aliasing problem and various effects of geometry which play an important role.  The observed drift rate is an aliased version of the true drift rate which is such that a subpulse drifts to the location of the adjacent subpulse (or a multiple thereof) in about one pulsar period. We show that small variations, of the order of 3-8\\%, in the mean drift rate are then enough to explain the apparent reversals of drift direction seen in the data. We find the mean circulation time of the drift pattern to be significantly longer than the predictions of the original Ruderman and Sutherland (1975) model and propose an explanation for this, which relates to modified models with temperature regulated partial ion flow in the polar vacuum gap. The small variations in drift rate are then explained by very small heating and cooling effects -- less than 3500 K change in the $\\sim 2.5\\times10^6$ K surface temperature of the neutron star polar cap.  From a detailed consideration of the variation of the mean subpulse separation across the main pulse window, we show that the circulating spark pattern is {\\it not}  centred around the dipole axis, but around a point much closer (within a degree or so) to the rotation axis. This is an indicator of the presence of a ``local pole'' corresponding to the non-dipolar magnetic fields that are expected to be present close to the neutron star surface.  \\OurPsr\\/ thus provides a very rich and powerful system in which to explore important aspects of the physics of pulsar radio emission and neutron star magnetospheres. } ",
        "introduction": "\\label{sec:intro}  Understanding how radio pulsars work remains a challenging unsolved puzzle in astrophysics. The phenomenon of drifting subpulses seen in some pulsars is widely thought to be an important key to understanding this mystery.  This phenomenon, first reported by \\cite{DC68}, manifests itself in the most basic form as an organised phase change of subpulses occurring in successive pulses, resulting in the formation of one or more drift bands spanning the pulse envelope \\citep{Backer73,JMR86}. The basic concept is illustrated in figure 8 of \\cite {Backer73}. In any single pulse, there are typically two to three approximately equispaced subpulses, with a separation of \\Ptwo\\/ degrees of longitude, whose typical values are $\\sim 10\\degr-20\\degr$.  The interval between recurrence of successive drift bands at a given pulse longitude is called \\Pthree\\/ and typical values range from 1 to 15 times \\Pone, the pulsar period. The intensity of the subpulses is systematically modulated as they move along the drift band.  Usually, the intensity peaks in the middle of the pulse envelope and decreases towards the edges; however, in a few cases, the opposite trend is also seen \\citep{TMH75}. The direction in which the subpulses drift -- either leading to trailing edge of the pulse window with pulse number, or the other way around -- is another important aspect.  Though both kinds of drifting pulsars are known, for a given pulsar, the direction in which the subpulses drift is usually fixed. However, there are a few exceptional pulsars where both directions of drift are seen (though at different times).  In the standard phenomenology \\citep[e.g.][]{JMR86}, cases of clear subpulse driftband patterns are usually seen in pulsars where the line of sight grazes along the edge of the radio emission beam. For pulsars where the line of sight has a more central traverse over the beam, the phenomenon can manifest itself as periodic intensity modulations of the subpulses at a given longitude, without the detection of any phase changes \\citep{Backer73}.  The observed periodicities related to patterns of drifting subpulses are usually found to be independent of radio frequency, thus excluding all frequency dependent plasma effects as plausible mechanisms of the drifting subpulse phenomenon. All the properties discussed above strongly suggest that this phenomenon is a manifestation of the circulation of subpulse-associated subbeams around the magnetic axis.  The drifting subpulse phenomenon finds a natural explanation in the model of \\citet[][RS75 hereafter]{RS75}. This model proposes the occurrence of discrete, localised spark discharges in a quasi-steady vacuum gap that forms just above the neutron star surface, as the seed activity that drives the radio emission mechanism.  Each spark generates a stream of electron-positron plasma that flows out along the dipolar field lines and, at heights that are estimated to be of the order of several tens of stellar radii, produces a beam of subpulse-associated radio emission due to some, as yet unidentified, plasma radiation mechanism. Further, this model proposes that the plasma of each spark discharge inside the vacuum gap undergoes a systematic motion around the magnetic axis, due to the \\ExB\\/ drift.  This, coupled with the rapidly fluctuating but rarely self-quenching nature of the sparks, leads naturally to the phenomenon of drifting subpulses. This drift rate is slower than the corotation rate and the sparks therefore lag behind with every rotation.  In this picture, each drifting spark would correspond to one band of drifting subpulses. Since the spark plasma can only lag the pulsar rotation, the model predicts drift in one direction only, depending on the relative orientation of the rotation and magnetic axes and the observer's line of sight for the given pulsar. The RS75 model also makes quantitative predictions of the total circulation time of the drifting pattern for a pulsar, which, in principle, can be checked against observations. Further refinements of the RS75 model include the models of Cheng \\& Ruderman (1980), Gil \\& Sendyk (2000, hereafter GS00) and Gil, Melikidze \\& Geppert (2003, hereafter GMG03) which attempt to overcome some of the deficiencies and problems identified with the RS75 model.  Other models that address the issue of subpulse drift are those of \\cite{am83,j84,ketal96,w03}.  However, there are some problems in the interpretation of the observations that make it difficult to compare with theoretical models.  Often, there are complications of the drift rates not being constant and stable with time.  Even when the drift rates (or value of \\Pthree) are stable and can be determined reliably, it is not clear whether they are the real values or not, as the measured values can be modified due to the effect of aliasing. This effect can occur if the sampling rate (once per pulsar period) is too slow to track the same subpulse-associated subbeam of radiation in a sequence of pulses.  Further, the total number of sparks, \\Nsp, needs to be known to ascertain the total circulation time of the pattern. In all this, the exact geometry of the observations needs to be known for correct interpretation of the results.  As a result, though there are about 40 pulsars with observations and measurements of the phenomenon of drifting subpulses \\citep[for a summary see][]{JMR86}, there are only a few for which quantitative results have been obtained which can be compared against the predictions of the models.  For PSR B0943+10, from a careful analysis of the fluctuation spectra, \\citet{DR99} were able to resolve the alias order for the observed periodicity (\\Pthree) and also infer the presence of 20 spark associated subpulse columns circulating with a time of $37 \\Pone~$. For PSR B0809+74, \\citet{vL_etal_2003} have recently shown that the drift rate we see is an unaliased version and hence $\\Pthree\\/ > 150\\Pone~$ and $\\Nsp\\/ \\ge 14$. For both cases, the final circulation time turns out to be significantly longer than the prediction of the RS75 model. Modified versions of this model (e.g. GMG03) attempt to account for these discrepancies.  Whereas observations of pulsars with drifting subpulses where the line of sight just grazes a cone of emission produces a clear signal for analysis and interpretation, a similar phenomenon when seen for pulsars where the line of sight samples a large fraction of the conal ring can produce additional insight into the phenomenon. This makes pulsars with wide profiles very interesting.  \\OurPsr\\/ is such a case. It has a very wide profile, showing emission over as much as 250\\degr\\/ of longitude at metre wavelengths \\citep{Durdin,Biggs,Reyes}.  However, it is not a widely studied pulsar, probably because of the fact that it nulls for as much as 70\\% of time \\citep{Durdin}, making it a difficult case to observe. The study reported by \\citet{Biggs} is the most exhaustive study published to date for this pulsar.  They found that there is subpulse structure at almost all longitudes in the wide profile of \\OurPsr, with remarkable patterns of drifting subpulses present over 200\\degr\\/ of longitudes.  They noted that the drift rate shows large variations, including reversals of the direction of the drift -- they concluded that this behaviour is inconsistent with the prevailing theoretical models for the phenomenon of drifting subpulses.  From the integrated pulse profiles and the polarization data at 408 and 610 MHz, \\citet{Biggs} concluded that \\OurPsr\\/ is likely to be a nearly aligned rotator, i.e. with rotation and magnetic axes almost parallel to each other, although their formal best fit values for the angle between the rotation and magnetic axes don't fully support this conclusion. However, in the morphological scheme of \\citet{JMR83,JMR93} it is classified as ``M'' type pulsar, i.e. having a central line of sight which goes close to the magnetic pole and produces distinct core and conal emission components in the average profile.  It is clear from the above description that \\OurPsr\\/ is an enigmatic pulsar that has the potential to reveal vital clues about the pulsar emission mechanism and therefore deserves a closer and more detailed study.  With this in aim, we have been carrying out high sensitivity observations of this pulsar with the Giant Metre-wave Radio Telescope (GMRT).  The results from our initial observations and the interpretations thereof are reported in this paper. In section \\ref{sec:obs}, we describe the observations and the basic data reduction.  Section \\ref{sec:RnI_main} presents the main results and our model that interprets these results, including determination of the emission geometry for this pulsar and detailed description of various effects that modify the true values of basic quantities related to the drifting phenomenon. In section \\ref{sec:discuss} we discuss the implications of different aspects of our model and present a final summary. In Appendix A we summarise the basic ideas relevant to the understanding of pulsar emission geometry, and in Appendix B we describe the details of our simulation studies. ",
        "conclusions": "\\label{sec:discuss}  \\paragraph{1. On the location and arrangement of sparks in PSR B0826$-$34 :} Our analysis and interpretation of the data for \\OurPsr\\/ reveals that we are looking at a pulsar that is very close to being an aligned rotator, with a value of \\A\\/ less than 5\\degr, and more likely to be around 1\\degr\\/ to 2\\degr.  Consequently, our line of sight traces out a circular track centred almost around the dipole magnetic axis.  This not only gives the exceptionally wide profile that we observe for this pulsar, but also allows us to sample the emission pattern along a ring around the magnetic axis.  This radiation pattern reveals the presence of multiple drifting bands of subpulses (up to 7 are visible), which we interpret as produced by subpulse producing plasma columns associated with multiple sparks circulating in the ring.  This is the first pulsar for which {\\it direct} evidence for such a large number of sparks has been obtained -- all previous cases show evidence for at most 2-3 drift bands visible within the pulse window.  Hence our results and interpretation can be considered as another direct confirmation of the idea of rings of circulating sparks on the polar cap, which on the average form conal beams of emission \\citep[][]{gks93,KG2002}.  The other interesting result is that the measured variation of \\Ptwo\\/ -- the longitude separation between the multiple drift-bands -- across the main pulse window is incommensurate with the values of \\A\\/ and \\B\\/ for the dipole geometry that we infer from the variation of main pulse width with radio frequency and the known polarization angle variation.  This leads us to a new and interesting conclusion : that the pattern of drifting subpulses does not circulate around the dipole magnetic axis. The actual point of circulation (which we refer to as the ``local pole''), is determined to be only $\\sim 3\\degr\\/$ away from the dipole axis; however, this small difference is enough to cause significant differences in the observed values of \\Ptwo.  What determines the location of this centre of circulation and how do we understand this in the context of the existing theoretical models?  In our interpretation, this centre of circulation of the drift pattern is a local pole produced by the more complicated magnetic field structure that could exist close to the neutron star surface.  Though we believe that the neutron star magnetic field is practically dipolar at the heights at which the radio emission is generated in the magnetosphere, there is now a growing set of evidence that the magnetic field close to the neutron star surface (up to a few stellar radii) may have higher order multi-pole structure \\citep[see ][ for review]{ug03}.  In fact, it could be a ``sun-spot'' like magnetic field (e.g. GS00). It is worth noting that the higher magnetic field strength and the much smaller radii of curvature that such a field configuration provides are actually essential for the working of the basic pair production mechanism inside the vacuum gap that drives the RS75 kind of models \\citep[][ and references therein]{gm02}. In such a scenario, it is possible to have on the polar cap, a local pole of the magnetic field configuration whose location is not aligned with that of the pole of the dipolar field.  In models with vacuum gaps close to the neutron star surface, it will be the structure of this non-dipolar magnetic field that will control the activities in the vacuum gap. In the model of GS00, the location of this local pole is taken as the location of the central spark (anchored to the local pole) that constitutes the core beam of emission. The rest of the polar cap is then populated by spark discharges that are arranged in concentric circles around the central spark, leading to a nested cone structure for the emission beam.  Within each circle, the sparks circulate due to the \\ExB\\ drift in a manner similar to the original RS75 model. Though the higher order field dominates close to the neutron star surface, with increasing height it is thought to make a smooth transition to the dipolar field that dominates at higher altitudes.  Thus we have a situation where the basic distribution of the seeds of emission (the sparks) on the polar cap is decided by the higher order surface magnetic field, while the actual height of the radio emission (and the evolution with frequency thereof) is dictated by the dipolar field at altitudes of typically about 50 stellar radii.  Our results and inferences thus lend strong support to this picture of strong non-dipolar magnetic fields close to the neutron star surface.  \\paragraph{2. The two cone model for PSR B0826$-$34 :} For most of the work presented here, we have concentrated on the emission properties of the main pulse region of this pulsar.  However, as we noted in section \\ref{sec:obs}, there is a weak interpulse seen in our 318 MHz average profile, which becomes gradually stronger as we go to higher frequencies, and even dominates over the main pulse emission at frequencies higher than $\\sim$ 1 GHz (see fig. \\ref{fig:freq_evln}).  The natural explanation for this is emission from a second, outer conal ring with the geometry contriving such that the line of sight samples a large fraction of the inner ring at the lower frequencies ($\\sim$ 300 MHz), while missing most of the outer ring.  Then, as the emission beams reduce in angular size at the higher frequencies, the line of sight first comes closer to the inner ring (in the main pulse region) -- at around 600 MHz -- and then actually crosses beyond the inner ring at frequencies above 1000 MHz, leading to relatively weak emission in the main pulse window.  This evolution naturally explains the increased ``filling-up'' of the bridge region in the main pulse window as we go from 318 MHz to 645 MHz. At the same time, the frequency evolution of the beams brings the outer conal ring closer to the line of sight in the interpulse region.  This evolution is such that at around 600 MHz, significant part of the outer ring begins to be seen by the line of sight and at frequencies around 1.4 GHz, the line of sight is actually grazing along the outer ring in the interpulse region.  We have modeled this configuration of two conal rings using a simulation program that takes a given distribution of sparks on the polar cap and translates them along dipolar field lines to frequency dependent heights of emission in the magnetosphere and generates the observed pulse profiles that would be seen along a specified line of sight (see Appendix~B for details). The dashed lines in figure \\ref{fig:freq_evln} show the simulated profiles obtained for a model having two conal rings of emission with frequency dependent radii determined by eqn.~(\\ref{eqn:r_em}) and centred around the local pole which is located at the point of the polar cap characterized by \\Ap$=0\\fdg436$, \\Bp$=3\\fdg064$ and $\\Phi_{off}=17\\fdg3$. As can be seen, the match with the observed profiles is quite good, providing direct support for the model of 2 conal rings.  One important point worth stressing is that once the amplitude ratio of main pulse to interpulse is fixed at one frequency, then the spectral evolution of these profile components is almost entirely due to dipolar spreading of magnetic field lines in the emission region.  Our two conal ring model for \\OurPsr\\/ makes clear predictions about how the average profile and single pulse characteristics should evolve with frequency and suitable experiments can be devised for testing these.  For example, one of the interesting things that it predicts is that the position angle curve should be shifted with respect to the centre of the main pulse window, because of the \\DelPhi\\/ rotational phase difference between the fiducial planes containing the local pole axis and the dipole axis.  In figure 6 of \\cite {Biggs}, there is some hint that the centre of the PA swing occurs at slightly later longitudes than the centroid of the main pulse, but the quality of the data is not good enough to ascertain this. Clearly, good quality polarimetric observations of \\OurPsr\\/ can help in testing this aspect of our model.  Further, since our model predicts the interpulse emission region to arise from a different conal ring, we expect that the subpulse properties here could be different from those for the main pulse region.  For example, depending on the number of sparks populating the outer ring, the measured value of \\Ptwo\\/ could be very different from the value we measure for the inner ring.  Perhaps the most interesting feature would be the value of drift rate and inferred values of \\Pthree\\/ and \\Pfour\\/ in the interpulse region at 600 to 1400 MHz.  \\paragraph{3. Pulse profile classification of PSR B0826$-$34 :} It is clear from our work that there is a need for a revision of the earlier classification of \\OurPsr as a ``M-type'' pulsar \\citep[e.g.][] {JMR93,LM88}, i.e. a pulsar where the line-of-sight and emission geometry are such that multiple cones of emission are visible along with the central core component.  In our model, all the emission seen in the main pulse window originates from a single conal ring of emission and that seen in the interpulse regions comes from the second, outer conal ring of emission; the line of sight thus does not sample the central core region at all.  Though the average profiles of this pulsar sometimes tend to show partial evidence for discrete emission components in the saddle region between the leading and trailing emission peaks in the main pulse region (as can be seen in figure \\ref{fig:freq_evln}), we believe these are artefacts of the quasi-longitude stationary drift pattern of the multiple drift bands, especially when average profiles are made over relatively short intervals of time, from limited number of pulses.  The classification of PSR B0826-34 as a ``M-type'' profile pulsar seemed strongly supported by the prominent circular polarization, reversing sense in the central part of the profile \\citep[] [their fig.~6]{Biggs}.  It therefore follows from our analysis of drifting subpulse patterns in this pulsar, that sense reversing circular polarization does not have to be associated with core emission only.  It can also appear in conal components, if the related subpulses (as in fig. \\ref{fig:stack_grey}) are quasi-stationary in phase \\citep[for more details see][]{gkms95}.  \\paragraph{4. Inferences from the drift behaviour of PSR B0826$-$34 :} We now discuss the results for the drifting behaviour of this pulsar in some detail.  Not only do we see subpulse emission from several spark-associated plasma columns, we are also able to detect the presence of correlated drift pattern from all of these subpulses.  The observed apparent reversals of drift direction, which are difficult to explain under any physical model of the subpulse drift phenomenon, are naturally explained as a combination of aliasing of the true drift rate coupled with small ($\\approx 3-8\\%$) variations of the true drift rate.  Thus, in our interpretation, the intrinsic drift rate does not show any reversals in sign, but only small changes in amplitude, making the behaviour of this pulsar compatible with these models.  The inferred values of both the mean drift rate and the proposed small variations are worthy of a detailed discussion, which we now present.  As we have shown, a unique solution for the spark and emission geometries for \\OurPsr\\/ is not possible with the current data, as the alias order and \\Nsp\\/ can not be unambiguously determined. Instead, for a given choice of the value of one of these, the value for the other that yields acceptable solutions can be found. The values for \\D, \\Pthree\\/ and \\Pfour\\/ can then be found for these cases and compared with expectations from models.  These results are summarised in Table \\ref{tbl2}, for three choices of the value of \\Nsp. The drift rates for most of the solutions are much lower than the values expected for the pure vacuum gap models (e.g. RS75, GMG03).  Correspondingly, the values of \\Pfour\\/ are substantially longer than the prediction.  This result for \\OurPsr\\/ is similar to the results obtained for PSR B0943+10 \\citep {DR99} and PSR B0809+74 \\citep {vL_etal_2003}. A natural explanation for this behaviour is provided by GMG03 who show that if the \\ExB\\/ drift is considered in a non-ideal vacuum gap, i.e. one in which there is a finite amount of charge flow \\citep [e.g.][]{CR80}, then the predicted value of the drift rate is reduced.  The proposed agency for the control of the ion flow in the gap is thermal regulation, i.e. the ratio of the surface temperature to the ion critical temperature controls the rate of ion flow, and hence the electrical potential drop across the gap and hence the value for the \\ExB\\/ drift rate.  This effect is characterised by a ``shielding factor'' by GMG03 who find typical values for it to be $\\sim$ 0.2.  For our possible solutions, the values of the shielding factor (SF) are shown in column 10 of Table \\ref{tbl2}.  For the slowest drift rate solution, this factor is 0.36 and it increases to 0.94 (i.e. close to vacuum gap conditions) for the case of $\\Nsp,k = 16,3$.  From this it is clear that solutions for $k>3$ are not viable as the implied drift rate would be too fast -- more than the maximum that is produced under pure vacuum gap conditions.  In the above scenario, the small variations around the mean drift rate would then be due to small changes in the neutron star's surface temperature in the polar cap region.  For example, to produce a 8\\% change in the drift rate, we find, using the formulation in GMG03, that only a 0.14\\% change in the surface temperature is needed.  This means a change of about $3500$~K out of a mean temperature of $\\sim 2.5\\times10^6$ K.  What could case such small and cyclic (though quasi-periodic at best) variations in the polar cap temperature is an interesting and open question. One possibility is a kind of ``self-regulating'' mechanism of heating and cooling of the polar cap by increased and reduced spark activities. Whatever it maybe, from our earlier deductions (at the end of section \\ref{sec:aliasing})  about the regions of slower and faster than normal drift rates, it would imply that the process of cooling below the mean temperature (which would lead to a faster than mean drift rate) is observed to be a relatively slower and smoother process than that the reverse process of heating (associated with slower than mean drift rates), which is found to be more abrupt and faster. Further, it is interesting that the observed drift rate fluctuations are found to be highly correlated for all the visible drift bands (section \\ref{sec:obs}), implying that the small temperature variations (if they are the cause of the drift fluctuations) occur simultaneously on a fairly large scale -- across a large part of the polar cap.  Here, detection and comparison of subpulse drift (and its variations) in the interpulse region (which we believe is due a second ring of sparks on the polar cap) with that in the main pulse region should produce valuable additional insight.  \\paragraph{5. General implications for the average pulsar population:} Finally, we turn to a discussion of the implications of slow variations of drift rate for the average pulsar population.  If the few percent variation of drift rate that we infer for \\OurPsr\\/ is a generic phenomenon affecting a large number of pulsars, then it could explain some of the mysteries of non-detection of periodic drift variations in the fluctuation spectra analysis of several pulsars.  Depending on the time scale of the drift rate fluctuations, compared to the pulsar period and the length of data taken for fluctuation spectra analysis, it is possible that drift rate variations could significantly, or even completely, blur the periodic signal.  Thus it is possible that several pulsars have drifting subpulses, but with some fluctuations in the drift rate, leading to reduced chances of detection of the phenomenon in a traditional fluctuation spectrum analysis.  Even amongst the ``classical'' drifting pulsars (e.g. PSR B0031$-$07, PSR B0943$+$10), variations in drift rate are a common observed feature.  As a specific example, we would like to point out the case of PSR B0540+23, which almost certainly shows stroboscopic phenomena due to aliasing, similar to \\OurPsr.  As reported by \\citet{now91}, the subpulse drift is irregular, without specified drift direction and with different speeds, including zero speed. As a result, there is no specified value of $P_3$, and there is no signature in the fluctuation spectrum.  Therefore, this pulsar is very similar to PSR 0826-34 \\citep[see Figs.~4, 5 and 6 in][]{now91}, except that the pulse profile is much narrower, implying a larger inclination angle.  \\paragraph{6. Summary:}  Using results from sensitive single pulse observations with the GMRT at 318 MHz, we have made significant progress in unraveling the complex drifting subpulse pattern of \\OurPsr.  In the process, we have shown that this pulsar is very close to being an aligned rotator.  This explains the unusually wide profile and multiple drift bands seen for this pulsar.  We have modeled the observed profiles and shown evidence for the presence of two conal rings of emission.  A detailed treatment of effects of aliasing and geometry on the drifting subpulse pattern has been provided.  We have given an explanation for the observed variability of the drift rate, including the apparent reversals of drift direction, within the framework of the currently available theoretical models.  As a result, we have shown that small variations in the surface temperature of the neutron star polar cap can explain the observed drift pattern. Finally, we have also demonstrated that the subpulse pattern shows evidence for the presence of non-dipolar magnetic fields close to the neutron star surface. Clearly, there is a lot to be learnt from \\OurPsr."
    },
    "0404/astro-ph0404093_arXiv.txt": {
        "abstract": "We have selected red MSX sources (RMS) that have the colours of massive young stellar objects (MYSOs). Our aim is to generate a large, systematically selected sample to address questions such as their luminosity function, lifetimes, clustering and triggering. Other objects such as UCHIIs, PN, PPN and AGB stars have similar IR colours and a large programme of ground-based follow-up observations is underway to identify and eliminate these from the sample of the red MSX sources. These include radio continuum observations, kinematic distances, ground-based mid-IR imaging, near-IR imaging and spectroscopy to distinguish. We report the progress of these campaigns on the 3000 candidates, with initial indications showing that a substantial fraction are indeed massive YSOs. ",
        "introduction": "Studies of massive star formation are currently hampered by the lack of large systematically selected samples of objects in the earliest phases. Any well-selected sample of massive YSOs must start from the IR where most of their bolometric luminosity emerges.  IRAS data has previously been used (e.g. Campbell et al. 1989; Chan et al. 1996; Sridharan et al. 2002), but its large beam size means that it is often confused in the galactic plane. The MSX survey of the plane in the mid-IR (Price et al. 2001) provides much better spatial resolution and virtually eliminates any bias against the dense and clustered environments where MYSOs are.  We have developed colour-cuts ($F_{21}>2F_{8}$ and $F_{8}<F_{14}<F_{21}$) to select MYSO candidates from the MSX PSC with additional colour-cuts ($F_{8}>5F_{K}$ and $F_{K}>2F_{J}$) using the 2MASS near-IR survey (Lumsden et al. 2002).  These colour-cuts and the elimination of known sources leaves a sample of about 1700 (leaving out the crowded galactic centre region). This sample still contains many other objects such as compact H II regions, PN, PPN, AGB stars, etc., which have the same IR colours. Our follow-up programme of observations is designed to identify and eliminate these. ",
        "conclusions": "The RMS survey is clearly finding large numbers of new MYSOs.  We currently estimate that we will deliver a well-selected sample of order 500 MYSOs with L$>10^{4}$L$_{\\odot}$. This is an order of magnitude larger than that currently known and will be close to being complete for the galaxy at the higher luminosities. This number is consistent with a simple estimate of the expected total number of MYSOs in the galaxy based on the accepted star formation rate, an assumed IMF and phase lifetime (Lumsden et al. 2002)."
    },
    "0404/astro-ph0404570_arXiv.txt": {
        "abstract": "Radio waves propagating from distant pulsars in the interstellar medium (ISM), are refracted by electron density inhomogeneities, so that the intensity of observed pulses fluctuates with time. The observed pulse shapes are used to diagnose electron-density distribution in the ISM. The theory relating the observed pulse time-shapes to the electron-density correlation function has developed for 30 years, however, two puzzles have remained. First, observational scaling of pulse broadening with the pulsar distance is anomalously strong; it is consistent with the standard model only when non-uniform statistics of electron fluctuations along the line of sight are assumed. Second, the observed pulse shapes are consistent with the standard model only when the scattering material is concentrated in a narrow slab between the pulsar and the Earth.  We propose that both paradoxes are resolved at once if one assumes stationary and uniform, but {\\em non-Gaussian} statistics of the electron-density distribution in the interstellar medium. Such statistics must be of L\\'evy type, and the propagating ray should exhibit a L\\'evy flight rather than the Gaussian random walk implied by the standard model. We propose that a natural  realization of such statistics may be provided by the interstellar medium with random electron-density discontinuities. A L\\'evy distribution has a divergent second moment, therefore, the standard approach based on the electron-density correlation function does not apply. We develop a theory of wave propagation in such a non-Gaussian random medium, and demonstrate its good agreement with observations. The qualitative introduction of the approach and the resolution of the anomalous-scaling paradox was presented earlier in [PRL~{\\bf 91}, 131101 (2003); ApJ~{\\bf 584}, 791 (2003)]. ",
        "introduction": "\\label{introduction} Observations of pulsar signals provide a valuable tool for investigating electron density distribution in the interstellar medium. The pulsar intrinsic signal is narrow in time, being about few percent of the pulsar period. The observed signal (averaged over many periods of pulsation) is broad and asymmetric with a sharp rise and a slow decay. The pulse broadening is attributed to the random refraction the waves experience while propagating in the interstellar medium.  \\subsection{Overview.} \\label{overview}  The shapes of the signals, and their scalings with wavelength~$\\lambda$ and with the pulsar distance have been investigated observationally and analytically for 30 years. The standard theory of interstellar scintillations (described in detail below) assumes that the propagating wave is refracted by random Gaussian gradients, thought to be a good approximation due to the central limit theorem~\\citep{tatarskii}. The symmetric Gaussian distribution is fully characterized by its second moment, and the standard theory of scintillations aimed at reconstructing this moment from observations. The contradictions of this theory with observations were noted in the early 1970's, although not many pulsars were investigated at that time to make a definitive conclusion.  The signal shape is characterized by its time-width,~$\\tau$, estimated at the $1/e$ level. In 1971, Sutton analyzed scalings of pulse broadening times, $\\tau$, with the radio wavelength,~$\\lambda$, and with dispersion measures, $DM$, corresponding to the pulsar's distance along the line of sight, $DM=\\int dz\\, n(z)$. Observational scaling for large dispersion measures, $DM > 30 \\mbox{pc cm}^{-3}$, is close to $\\tau\\propto \\lambda^4 DM^4$, while the theory gave $\\tau\\propto \\lambda^4 DM^2$. The recent observational results are shown in Fig.~(\\ref{elbow}). To overcome the difficulty with the anomalous DM-scaling, Sutton suggested that the interstellar turbulence was not statistically uniform along the lines of sight, so that the lines of sight for more distant pulsars intersected regions with stronger turbulence.  The second paradox was encountered by \\citet{williamson1,williamson2,williamson3}, who compared the observed shapes of the pulses with the shapes predicted by the standard theory. He obtained a surprising result that the model of continuous turbulent medium was not consistent with observations. Rather, the best agreement was given by the models where all the scattering material was concentrated either in a thin screen or in a slab covering only $1/4$ of the line of sight between the pulsar and the Earth,  with nearly equal quality of fit to the data [see the discussion and Fig.~4 in \\citep{williamson3}].  These two assumptions that are necessary to reconcile the standard theory with the observations may have physical grounds since the interstellar medium is not uniform and denser regions (with stronger turbulence) are encountered closer to the galactic center. However, in this paper we discuss a simple and perhaps more plausible physical explanation for the Sutton and Williamson paradoxes. According to this explanation, the Sutton and Williamson paradoxes are not the consequence of the non-uniform, large-scale galactic distribution of the electron density responsible for scattering. Rather, they reflect the universal properties of the microscopic structures of the density fluctuations.  \\subsection{L\\'evy flights. Mathematical background.} \\label{levyflights}  In our recent work~\\citep{boldyrev1,boldyrev2,boldyrev3} we proposed that the Sutton paradox is the evidence of statistically uniform but {\\em non-Gaussian} electron distribution in the interstellar medium. We noted that the time-broadened pulses belong to those observational objects that depend not on the moments of the electron-density distribution function, but on its full shape. The standard Gaussian approach did not recognize this fact because in the Gaussian case, knowledge of the second moment is equivalent to knowing the distribution function itself. However, for a general non-Gaussian electron-density distribution the second moment may formally diverge, and the electron-density correlation function may not characterize the pulse shape.  Physically, this divergence means that the second moment of ray-angle deviation is dominated not by the bulk of its distribution function, but by the extremely far cut-offs at the tails of this function [as we will see in the example of Section~\\ref{boundaries}]. In contrast, the observed signal shape is determined by the bulk of the distribution function and is not sensitive to its far-tail cut-offs. In other words, the electron-density correlation function and the shape of the observed pulsar signal provide different, complementary descriptions of the electron distribution in the ISM.  We also note that in the L\\'evy model, the pulse shapes may provide a unique opportunity to diagnose plasma fluctuations at very small scales, $>10^8$~cm. These scales can be smaller than the Coulomb mean-free path in HII regions, and can be close to the ion gyro radius. Observations of pulse shapes may thus provide diagnostics for magnetized plasma fluctuations, which may help to elucidate the role of magnetic fields in interstellar turbulence, see also~\\citep{armstrong,goldreich}.   In the theory of scintillation, the quantity of interest is the transverse electron-density difference, averaged along the line of sight, $\\Delta n_{\\perp}=\\int dz\\,[n({\\bf x}_1, z)-n({\\bf x}_2, z)]$ . A wave propagating in the interstellar medium is scattered (refracted) by density inhomogeneities, and the statistics of the refraction angle is related to the statistics of density differences across the line of sight.  In the approach proposed in~\\citep{boldyrev1,boldyrev2} we assumed that the distribution function of~$\\Delta n_{\\perp}$ was stationary and uniform, but non-Gaussian, and it had a power-law decay as $|\\Delta n_{\\perp}|\\to \\infty$. If this distribution does not have a second moment, then the sum of many density differences along the line of sight does not have the Gaussian distribution, i.e. the Central Limit Theorem does not hold. Instead, the limiting distribution, if it exists, must be the so-called L\\'evy distribution, which is common in various random systems~\\citep{klafter}. By construction, such a distribution is stable under convolution, i.e. the appropriately rescaled sum of two independent variables drawn from the same L\\'evy distribution, has again the L\\'evy distribution. The Gaussian distribution is a particular case of this more general stable distribution.  The Fourier transform (the characteristic function) of an isotropic L\\'evy probability density function, $P_{\\beta}(\\bsigma)$, has the simple form, \\begin{eqnarray} F(\\bmu)=\\int\\limits_{-\\infty}^{\\infty}d\\bsigma P_{\\beta}(\\bsigma) \\exp(i\\bmu\\cdot \\bsigma)=\\exp(-  |\\bmu|^{\\beta}), \\label{levy} \\end{eqnarray} where the L\\'evy index $\\beta$ is a free parameter, $0<\\beta \\leq 2$. This form is valid for an arbitrary dimensionality of the fluctuating vector~$\\bsigma $. As we explain below, in our case this vector is two-dimensional, $\\bsigma\\propto {\\hat {\\bf y}}\\int dz\\,[n(z,{\\bf x}_1)-n(z,{\\bf x}_2)]$, where ${\\bf x}$ is the coordinate in the plane perpendicular to the line of sight, $z$, and ${\\hat {\\bf y}}$ is the unit vector in the direction connecting the two points in this plane, ${\\hat {\\bf y}}={\\bf y}/|{\\bf y}|$, where ${\\bf y}={\\bf x}_1-{\\bf x}_2$.  Equation~(\\ref{levy}) can be taken as the definition of the {\\em isotropic} L\\'evy distribution, and in the present paper we will consider only isotropic distributions. However, it is important to note that a stable distribution need not be isotropic [see, e.g.,~\\citep{nolan}]. In the simplest case, the anisotropic distribution is a copy of the isotropic distribution, rescaled along one axis; but far less symmetric L\\'evy distributions are possible. Non-isotropic stable distributions might arise in the case of the magnetized interstellar medium, when wave scattering has one or several preferred directions. For example, anisotropic scattering may be characteristic of interstellar MHD turbulence, as described by~\\citet{goldreich,lithwick,cho,chandran,scalo}, where the preferred direction is given by the large-scale magnetic field.   The sum of $N$ L\\'evy distributed variables scales as $\\vert \\sum^{N}\\bsigma \\vert\\sim N^{1/\\beta}$, which becomes diffusion in the Gaussian limit $\\beta=2$. For $\\beta < 2$,  the probability distribution function has algebraic tails, $P_{\\beta}(\\bsigma)\\sigma \\,d\\sigma\\sim |\\bsigma|^{-1-\\beta}\\,d\\sigma$, for large $|\\bsigma|$, and its moments, $\\langle |\\bsigma|^{\\zeta} \\rangle$, of the order $\\zeta \\geq \\beta$ diverge. In \\citep{boldyrev1,boldyrev2}, we considered smooth density fluctuations, $\\sigma\\propto |{\\bf x}_1-{\\bf x}_2|$, and the proposed L\\'evy statistics of ray-angle deviation lead to the pulse-broadening scaling $\\tau\\propto \\lambda^4DM^{(2+\\beta)/\\beta}$. This agrees with observations when $\\beta\\approx 2/3$, and provides a natural resolution to the Sutton paradox.  \\subsection{The results of the paper.} \\label{overviewofthispaper}  In this paper we consider scintillations produced by a general, non-smooth density fluctuations, $\\sigma\\propto |{\\bf x}_1-{\\bf x}_2|^{\\alpha/2}$; the limit of smooth density fluctuations corresponds to~$\\alpha=2$.  More precisely, exponent $\\alpha$ denotes the scaling of the density-difference PDF defined in formula~(\\ref{levy}), i.e. this PDF should have the form $P_{\\beta}(\\sigma, y)\\sim P_{\\beta}(\\sigma/y^{\\alpha/2})/y^{\\alpha/2}$. In the non-smooth case, the effective L\\'evy index that enters the expressions for angle deviations is different from $\\beta$, and is given by $\\gamma=\\alpha\\beta/2$, as is explained in formulae~(\\ref{levyscreen}) and~(\\ref{tau}) below.  As we noted above, in the L\\'evy case the electron-density correlation function (or second order structure function) is contributed to by the far-tail cut-off of the electron-difference distribution function. We will denote the correlation-function scaling as $\\langle \\sigma^2\\rangle \\propto |{\\bf y}|^{\\delta}$, where the scaling exponent, $\\delta$, is in general different and independent of the L\\'evy distribution scaling exponents $\\alpha$ and $\\beta$. The exponent $\\delta$ is related to the Fourier spectrum of electron density fluctuations, $\\langle |n(k)|^2 \\rangle k^2 \\propto k^{-\\delta}$.   In the present paper we demonstrate that the proposed L\\'evy statistics of ray-angle deviations provide a resolution for the Williamson paradox as well. For this purpose we, first, develop a theory of wave propagation in a L\\'evy random medium by approximating the medium by a uniform series of scattering screens. Then, we provide a general method for constructing the pulse-broadening function for an arbitrary L\\'evy index~$\\gamma$. And finally, we compare our results with the observational signals of large-dispersion-measure pulsars, recently published by~\\citet{bhat} and by \\citet{ramach}. We obtain that the observational shapes agree well with the predictions of our theory for a uniform medium with the L\\'evy index $\\gamma\\sim2/3$ to~$1$, while they are inconsistent with the Gaussian theory corresponding to the L\\'evy index~$\\gamma=2$.  We also discuss the effect of ``over-scattering'' which is inherent for the L\\'evy scintillations, because the time-signals have long algebraic tails that do not decay to zero during pulse periods. This effect is crucial for comparing analytical and observational data, as we show in Section~\\ref{observation}.  For illustrative purposes, we also present a particular model of the density distribution that produces the L\\'evy statistics of scintillations. The proposed density distribution is strongly spatially intermittent,  and it can be visualized as follows. Imagine that the electron density is concentrated in separated regions that have sharp irregular boundaries.  We make the simplest assumption that these boundaries are random, similar to randomly folded two-dimensional sheets or shocks. If one further assumes that the line of sight intersects the boundary at a random angle whose distribution is uniform in all directions, such a picture corresponds to the particular case of a L\\'evy  distribution of $\\Delta n$, with $\\alpha=2$, $\\beta=1$, the so-called Cauchy distribution. This distribution is rather distinct from the Gaussian one, and is  close to the distribution that has been predicted in our theory.  \\subsection{Related models.} \\label{relatedmodels}  The importance of randomly oriented discontinuous objects that can be encountered across the line of sight has been emphasized in earlier theoretical work on scintillation, see e.g., \\citep{lambert} and \\citep{lazio}. As noted by \\citet{lambert}, sharp density discontinuities may characterize stellar wind boundaries, supernova shock fronts, boundaries of HII regions at the Str\\\"omgren radius, etc. They may also arise from supersonic turbulent motion. \\citet{rickett3} proposed that similar highly intermittent density structures may be responsible for observed low-frequency modulations (fringes) of dynamic pulsar spectra.  The approach of \\citet{lambert} utilized the non-Kolmogorov spectra of discontinuous density fluctuations. \\citet{lazio} considered wave scattering by confined or heavily modulated screens (such as disks, filaments, etc.), when the statistics of angular deviations were generally assumed Gaussian, with parameters varying along the screens. Both papers discussed important aspects of non-Kolmogorov and spatially intermittent electron-density distribution in the ISM, however, they implied the existence of the second moments of the scattering-angle distributions. Therefore, in the earlier considerations, the effects that we discuss in the present paper could not be discovered.  As we demonstrate in section~\\ref{boundaries}, sharp density discontinuities can, in fact, produce a non-Gaussian distribution of the scattering angle, whose second moment diverges. In our approach we exploit such intermittent density statistics in their full generality. To elucidate the universal role of L\\'evy distribution, we keep our consideration as simple as possible, assuming statistically uniform and isotropic scattering screens (although generalizations for the non-isotropic and non-uniform cases are possible). The presence of strongly-scattering structures is naturally represented in our model by slowly decaying, power-law tails of the scattering-angle distributions. Our model has a simple physical interpretation and provides a practical way of calculating pulse shapes and pulse scalings. Most importantly, it naturally resolves both the Sutton and the Williamson observational paradoxes, in a manner that is simpler and complementary to the standard Gaussian picture.    The L\\'evy-flight model may also be relevant to observations of extreme scattering events, such as those reported by \\citet{fiedler,wolszczan, stinebring}. Such events have been investigated theoretically, as by~\\citet{rickett3,romani}. We note that these theoretical studies invoke uncommon incidents of scattering much larger than typical values to explain these events.  In this sense, these events are consistent with the L\\'evy model we develop here, which includes rare, large events in a statistically-stationary way. However, we do not yet know whether these phenomena might fit into a single model for interstellar scattering along with the common phenomena we discuss in this work.  The answer might possibly depend on the details of the picture for turbulence in which the L\\'evy model is realized.  Accordingly, we do not discuss the extreme scattering events in the present work. ",
        "conclusions": "\\label{conclusions}  In the present paper, we have investigated the observational pulse shape. Another object directly related to the electron distribution function is the visibility function, $\\langle \\Phi({\\bf x}_1, t)\\Phi^*({\\bf x}_2,t)\\rangle$, measured at the locations ${\\bf x}_1$ and ${\\bf x}_2$ at the Earth. However, the electron fluctuations probed by this function are limited to small scales, $|{\\bf x}_1-{\\bf x}_2|\\lesssim 10^8$cm, which can be comparable to the inner scale of the turbulence~\\citep{gwinn2,kaspi,gupta} where the statistics of the ISM are different from the statistics responsible for pulse broadening. The visibility function can be easily constructed in our multi-screen model by solving Eqs.~(\\ref{vequation},\\ref{freeequation}) for $\\Delta k=0$, however, we leave its analysis for future communication. Here, we estimate what restrictions would be imposed on our theory by an inner scale of the order $y_{in}\\sim 10^{8} \\mbox{cm}$.  Due to the formulas of Section~\\ref{multi-screen}, large-angle refraction is provided by small-scale fluctuations of electron density.  In the case of L\\'evy statistics, a sum of random variables can be dominated by a single term, therefore, from formula~(\\ref{tau}) we can estimate $(\\Delta {\\tilde {\\xi}}_{max})^2\\sim 2\\pi \\tau c/(z_0 N \\lambda^2 B^2)$. From formula~(\\ref{intensity-non-dim}), the electron-density fluctuations producing such an angle deviation should be strong at the scale $y_{min}\\sim 2\\pi/(B \\Delta {\\tilde {\\xi}}_{max}) \\sim [2\\lambda^2 Nz_0/(\\tau c)]^{1/2}$, or \\begin{eqnarray} y_{min}\\sim\\sqrt{\\frac{2 c DM}{ \\langle n \\rangle \\tau \\nu^2}}, \\label{yin} \\end{eqnarray} where we introduced the dispersion measure $DM=\\langle n \\rangle N z_0$, and the signal frequency, $\\nu=c/\\lambda$. Quite conveniently, formula (\\ref{yin}) includes only the quantities averaged along the line of sight, but not the particular parameters $l$, $l_0$, $\\Delta n_0$, and $\\alpha$ of the scattering screens.  As an example, let us consider the parameters for the large-dispersion-measure pulsar J1852+0031 (DM=680.0) observed at 1175 MHz in \\citep{bhat}. Substituting $\\nu\\sim 10^9\\,\\mbox{s}^{-1}$, $DM\\sim 2\\cdot 10^{21} \\,\\mbox{cm}^{-2}$, $\\tau \\sim 0.5\\,\\mbox{s}$, and assuming that $\\langle n \\rangle \\sim 3\\cdot 10^{-2}\\,\\mbox{cm}^{-3}$, we obtain $y_{min}\\sim 10^8\\,\\mbox{cm}$. Therefore, the condition $y_{min}\\sim y_{in}$ is marginally satisfied. Due to the scaling $\\tau\\propto \\lambda^4 DM^4$, smaller-dispersion-measure pulsars satisfy this condition even better. For the cases when this condition is violated, a more appropriate description will be given by ``truncated'' L\\'evy distributions, see e.g.,~\\citep{nakao}.    In conclusion we would like to mention the observational correlation functions that are expressed through the higher-than-second moments of the wave amplitude, $\\Phi({\\bf x},t)$. These observational quantities require a special consideration that at present is not available to us.  The examples of such quantities are given by the frequency correlation function that is used to characterize pulse broadening for $DM<100 \\,\\mbox{pc cm}^{-3}$, see, e.g., \\citep{gwinn}, or by the intensity correlation functions that are used to characterize the density statistics on scales larger than the Earth size \\citep{shishov}. The scaling of such functions may be easily found in the Gaussian theory when these functions can be related to the second-order moment, $\\langle (\\Delta n)^2 \\rangle$. However, the meaning of these correlation functions in the L\\'evy case has yet to be determined. For instance, it seems possible that the deviation of the pulse-width scaling from $\\tau\\propto DM^4$ for the lower part of the elbow diagram, in  Fig.~(\\ref{elbow}), is due to the different method of reconstructing $\\tau$ in this region. Indeed, below $DM\\sim 100\\mbox{pc cm}^{-3}$, $\\tau$ is not measured directly, but rather is defined as the inverse frequency decorrelation bandwidth~\\citep{sutton,rickett1,rickett2,ramach,gwinn,bhat}.  To summarize, in the present paper we proposed a multi-screen model of the uniform, non-Gaussian interstellar medium, and we found its analytic solution. We demonstrate that this model explains both the shapes and the scalings of the observational pulse profiles, and it is free of the Sutton and Williamson paradoxes that could hardly be resolved in  the standard Gaussian theory of scintillations. Also, we provided a physical model of density distribution having random discontinuities, which produces L\\'evy scintillations and which may serve as a good approximation for the analysis of the observational data. We have demonstrated that the observational pulse shapes depend on the entire shape of the distribution of electron density, rather than simply on the second moment. Pulse profiles may therefore serve as a valuable tool for reconstructing this function from observations, and we have developed the corresponding method in the present paper."
    },
    "0404/astro-ph0404436_arXiv.txt": {
        "abstract": "We present the HIPASS Bright Galaxy Catalog (BGC) which contains the 1000 \\HI-brightest galaxies in the southern sky as obtained from the \\HI\\ Parkes All-Sky Survey (HIPASS). The selection of the brightest sources is based on their \\HI\\ peak flux density (\\Speak\\ $\\ga$ 116 mJy) as measured from the spatially integrated HIPASS spectrum. The derived \\HI\\ masses range from $\\sim10^7$ to $4 \\times 10^{10}$\\Msun. While the BGC ($z < 0.03$) is complete in \\Speak, only a subset of $\\sim$500 sources can be considered complete in integrated \\HI\\ flux density (\\FHI\\ $\\ga 25$ Jy\\kms).  The HIPASS BGC contains a total of 158 new redshifts. These belong to 91 new sources for which no optical or infrared counterparts have previously been cataloged, an additional 51 galaxies for which no redshifts were previously known, and 16 galaxies for which the cataloged optical velocities disagree. Of the 91 newly catalogued BGC sources, only four are definite \\HI\\ clouds: while three are likely Magellanic debris with velocities around 400\\kms, one is a tidal cloud associated with the NGC~2442 galaxy group. The remaining 87 {\\em new} BGC sources, the majority of which lie in the Zone of Avoidance, appear to be galaxies. We identified optical counterparts to all but one of the 30 {\\em new} galaxies at Galactic latitudes $|b| > 10\\degr$. Therefore, the BGC yields no evidence for a population of \"free-floating\" intergalactic \\HI\\ clouds without associated optical counterparts.  HIPASS provides a clear view of the local large-scale structure. The dominant features in the sky distribution of the BGC are the Supergalactic Plane and the Local Void. In addition, one can clearly see the Centaurus Wall which connects via the Hydra and Antlia clusters to the Puppis filament. Some previously hardly noticed galaxy groups stand out quite distinctively in the \\HI\\ sky distribution. Several new structures are seen for the first time, not only behind the Milky Way. ",
        "introduction": "% \\label{sec:intro} Neutral hydrogen (\\HI) is a major component of the interstellar medium (ISM) in galaxies. The ISM provides fuel for the initial formation of molecular clouds and stars, and conversely acts as a reservoir for recycled gas from stars and supernovae. The total \\HI\\ content of galaxies at a given epoch provides important constraints on the evolution of galaxies (Pei, Fall \\& Hauser 1999). At high redshifts, \\HI\\ content is commonly estimated through observations of the damped Lyman-$\\alpha$ (DLA) absorption-line systems (Rao \\& Turnshek 2000; Storrie-Lombardi \\& Wolfe 2000; Turnshek \\& Rao 2002). These observations appear to indicate that the \\HI\\ density of the Universe is substantially higher at redshift $z \\approx 1-3$ compared to the $z=0$ density measured from 21-cm observations of optically-selected galaxy samples (e.g. Rao \\& Briggs 1993). However, there are large uncertainties in the interpretation of DLA statistics at high redshift because of lensing and dust obscuration. There are also uncertainties at low redshifts because our knowledge of \\HI\\ in galaxies is heavily biased towards optically-selected samples. Are there, for example, many gas-rich galaxies or intergalactic clouds missing from the nearby galaxy census as some absorption-line observations of low column density systems ($10^{12.5} - 10^{15.5}$ cm$^{-2}$) appear to indicate (Penton, Shull \\& Stocke 2000; McLin et al. 2002; Ellison et al. 2002)\\,?  Within individual galaxies, \\HI\\ is frequently found to extend well outside the stellar radius (e.g. Broeils \\& van Woerden 1994; Meurer et al. 1996; Salpeter \\& Hoffman 1996; Broeils \\& Rhee 1997). Between galaxies neutral hydrogen is found mostly in the form of tidal tails, bridges and rings (Hibbard \\& van Gorkom 1996; Duc \\& Mirabel 1997; Hibbard et al. 2001; Ryder et al. 2001; Koribalski \\& Dickey 2004), which trace the interaction history of galaxies. Out of tidal \\HI\\ debris, which constitute part of the intergalactic medium, new galaxies can form. These are known as `tidal-dwarf galaxies' because of their similarity to classical dwarf irregulars and blue compact dwarf galaxies (Barnes \\& Hernquist 1992; Duc \\& Mirabel 1998; Duc et al. 2000; Hibbard et al. 2001; Braine et al. 2001). The closest and most prominent interacting galaxies are, of course, the Magellanic Clouds, which are associated with a range of extended tidal \\HI\\ features, most prominently the Magellanic Bridge between the LMC and SMC, the Magellanic Stream (Mathewson, Cleary \\& Murray 1974) and the recently discovered Leading Arm (Putman et al. 1998; Putman et al. 2003). Other \\HI\\ structures such as the Leo ring (Schneider 1989) and the tidal \\HI\\ remnant known as the Virgo cloud (Giovanelli \\& Haynes 1989; Chengalur et al. 1995) have been discovered by accident and indicate the potential for new discoveries in blind \\HI\\ surveys.  The \\HI\\ content of normal galaxies varies strongly, depending on morphological type, optical diameter and the environment (Roberts \\& Haynes 1994; Haynes, Giovanelli \\& Chincarini 1984; Giovanelli \\& Haynes 1990; Solanes, Giovanelli \\& Haynes 1996; Solanes et al. 2001; Verdes-Montenegro et al. 2001). Whereas early-type galaxies tend to be gas-poor (Duprie \\& Schneider 1996; Sadler 2001; Oosterloo et al. 2002), late-type galaxies are typically gas-rich (Matthews, Gallagher \\& Littleton 1995; Solanes et al. 2001). \\HI\\ can also be easily detected in many low surface-brightness (LSB) and late-type dwarf galaxies, often barely visible in the optical (e.g., Impey \\& Bothun 1997). Therefore, blind \\HI\\ surveys will reveal classes of galaxies that are under-represented in optical surveys. Since galaxy clustering is known to decrease from early- to late-type galaxies (Giuricin et al. 2001) this will affect our overall view of the local large-scale structure.  The galaxy distribution in the nearby Universe reveals many coherent structures (Geller \\& Huchra 1989; Fairall et al. 1990; Fairall 1998), such as the Supergalactic Plane (SGP), the Fornax and Centaurus Walls, as well as the Local Void (to name the most prominent in the southern sky). Since our knowledge about the local large-scale structure has, so far, mostly been determined by extensive optical and infrared surveys of the sky (e.g. Lauberts 1982, Jarrett et al. 2003), not much is known about the connectivity of structures across the artificial gap created by the extinction of dust in our own Galaxy. This can affect up to $\\sim$25\\% of the sky (e.g. Kraan-Korteweg \\& Lahav 2000) and obscures our view of important regions such as the putative Great Attractor (Lynden-Bell et al. 1988). However, 21-cm observations are unaffected by dust extinction, so are easily able to complement optical and infrared galaxy surveys and substantially improve the census of galaxies and the measurement of the baryon content of the local Universe.  \\subsection{\\HI\\ Surveys} \\label{sec:HIintro} Many \\HI\\ surveys have been targeted at cataloged galaxies, galaxy groups, clusters and selected regions --- for a summary see Salzer \\& Haynes (1996) or Zwaan et al. (1997). The biggest targeted \\HI\\ surveys to-date are those by (a) Fisher \\& Tully (1981a) who obtained \\HI\\ spectra for 1171 galaxies (out of 1787 targets) with the NRAO 91-m and 43-m telescopes and the Bonn 100-m telescope, (b) Schneider et al. (1990) who obtained \\HI\\ spectra for 762 dwarf and other LSB galaxies with the Arecibo telescope, (c) Mathewson, Ford, \\& Buchhorn (1992; hereafter MFB92) who obtained \\HI\\ spectra for 551 galaxies with the Parkes telescope, (d) Haynes et al. (1997) who obtained \\HI\\ spectra for about 500 galaxies in 27 Abell clusters visible from Arecibo, (e) Theureau et al. (1998) who present 2112 \\HI\\ spectra of cataloged spiral galaxies obtained with the Nan\\c{c}ay radiotelescope, and (f) Haynes et al. (1999) who obtained \\HI\\ spectra for 881 galaxies, over 500 of which were detected with the 91-m Green Bank telescope.  Overall, there are currently about three times more \\HI\\ measurements in the northern hemisphere ($\\ga$10300) than in the south ($\\ga$3500), as determined from the Lyon/Meudon Extragalactic Database (LEDA). At declinations south of --45\\degr\\ ($\\sim$30\\% of the southern sky) only 560 \\HI\\ measurements are available. The most prominent (targeted) southern \\HI\\ catalog is that by MFB92 who published \\HI\\ profiles for 551 galaxies of type Sb to Sd. Earlier, Reif et al. (1982) and Longmore et al. (1982) successfully obtained 196 and 100 \\HI\\ profiles, respectively, with the Parkes telescope, and Fouqu\\'e et al. (1990a) published \\HI\\ profiles of 242 southern late-type galaxies in the declination range --38\\degr\\ to --17\\degr\\ obtained with the Nan\\c{c}ay telescope.  One of the earliest {\\em blind} \\HI\\ surveys was that by Kerr \\& Henning (1987) who detected 16 previously uncataloged galaxies in the Zone of Avoidance (ZOA) and considered the possibility of future whole sky surveys with multibeam instruments (see also Henning 1995). Schneider (1996) gives a summary of other blind \\HI\\ survey programs. Until recently, the Arecibo Dual-Beam Survey (ADBS) and the deeper Arecibo \\HI\\ strip survey (AHISS, see Zwaan et al. 1997) were the largest blind 21-cm surveys (see Table~\\ref{tab:hisurveys}) carried out, covering areas of 430 and 65 square degrees, respectively. With the advent of a 21-cm multibeam system at the 64-m Parkes telescope (Staveley-Smith et al. 1996) as well as new observing and data reduction software (Barnes et al. 2001), much larger and deeper surveys are now possible. For example, Henning et al. (2000) did a shallow survey of 1800 square degrees of sky in the southern Zone of Avoidance, finding 110 \\HI\\ sources, 67 of which had no previously published optical counterpart. \\\\  This paper is organized as follows: the observations and data reduction are described in Section~2, followed by a detailed description of the `HIPASS Bright Galaxy Catalog' in Section~3. The \\HI\\ properties from the BGC are discussed in Section~4, followed by a study of the large-scale structure in Section~5. Our conclusions are given in Section~6.  In related papers: Ryan-Weber et al. (2002) study 138 galaxies from the BGC which had no redshift measurements prior to the commencement of the \\HI\\ multibeam surveys. Zwaan et al. (2003) analyze the BGC completeness and derive a new \\HI\\ mass function. ",
        "conclusions": "\\label{sec:conclusions} The HIPASS Bright Galaxy Catalog is a subset of HIPASS and contains the 1000 \\HI-brightest sources in the southern sky ($\\delta < 0\\degr$) based on their peak flux density (\\Speak\\ $\\ga$ 116 mJy) as measured from the spatially integrated \\HI\\ spectrum. We obtained previously catalogued optical (or infrared) counterparts for 909 of the 1000 BGC sources. The HIPASS BGC is the second largest sample of galaxies from a blind \\HI\\ survey to date, covering a volume of $\\sim2 \\times 10^6$ Mpc$^3$, only surpassed by the HIPASS Catalog itself. While the BGC is complete in peak flux density, only a subset of $\\sim$500 BGC members can be considered complete in integrated \\HI\\ flux density (\\FHI\\ $\\ga 25$ Jy\\kms).  The lowest detectable \\HI\\ flux density for relatively narrow line-width sources in the BGC (galaxies as well as \\HI\\ clouds) is given by \\FHI\\ = $0.7 \\times$ \\Speak\\ $\\times$ \\wfi\\ (see Fig.~\\ref{fig:w50fhi}) corresponding to an \\HI\\ mass of $2 \\times 10^4$ \\wfi\\ $D^2$ (see Fig.~\\ref{fig:logdmhi}). This means \\HI\\ sources with \\wfi\\ = 50\\kms\\ and \\MHI\\ $\\ga 10^8$ (10$^9$) \\Msun\\ can be detected out to 10 (30) Mpc.  Of the 91 newly catalogued BGC sources, 87 appear to be galaxies. A large fraction of these (37) lie at Galactic latitudes $|b| < 5\\degr$; all but five were previously uncovered by the \\HI\\ Zone of Avoidance shallow survey (HIZSS; Henning et al. 2000). For all but one of the 30 {\\em new} galaxies at $|b| > 10\\degr$ optical counterparts were discovered in the Digitized Sky Surveys (see Ryan-Weber et al. 2002). For a further 51 galaxies in the BGC no redshifts were previously known. In addition, we identify 16 galaxies (i.e., $\\sim$2\\% of the BGC) where the cataloged optical velocity is potentially incorrect, resulting in a total of 158 new redshifts.  There are four \\HI\\ clouds in the BGC, three of which are most likely Magellanic debris (HIPASS J1712--64, J1718--59, and J1616--55) with systemic velocities in the range 390--460\\kms. Further, detailed investigations are under way. We discovered one intergalactic \\HI\\ cloud, HIPASS J0731--69, which is associated with the galaxy NGC~2442 (Ryder et al. 2001). The cloud, which lies at a projected distance of 38\\arcmin\\ (176 kpc) from NGC~2442 ($D$ = 15.9 Mpc) has an \\HI\\ mass of nearly 10$^9$\\Msun. We could have detected similar \\HI\\ clouds up to a distance of 30 Mpc, with a separation of at least 20\\arcmin\\ from an associated \\HI-rich galaxy. This detection suggests that their space density is about 1/1000th that of galaxies with the same \\MHI\\ (see also Briggs 1990). Other known tidal \\HI\\ clouds such as those near NGC~3263 (HIPASS J1029--44b) and NGC~7582 (HIPASS J2318--42), see Appendix~\\ref{app:hiwidths}, are contained in the BGC as part of their associated galaxy or galaxy group.  Therefore, no definite free-floating \\HI\\ clouds without stars have been discovered in the search volume. Following Fisher \\& Tully (1981b) we derive an \\HI\\ density of $\\rho_{\\rm HIclouds} < 5.4 \\times 10^{5}$\\Msun\\,Mpc$^{-3}$ or $<3.7 \\times 10^{-35}$\\,g\\,cm$^{-3}$ for \\HI\\ clouds with \\wfi\\ = 100\\kms\\ and \\MHI\\ = $10^7 - 10^{10}$\\Msun. We compare the limit for $\\rho_{\\rm HIclouds}$ with the critical density for the closure of the Universe, $\\rho_{\\rm c}$, the baryonic mass density, $\\rho_{\\rm b}$, and the \\HI\\ mass density of galaxies, $\\rho_{\\rm HIG}$, (see Zwaan et al. 2003): $\\rho_{HIclouds}$/$\\rho_{\\rm c} < 3.5 \\times 10^{-6}$, $\\rho_{HIclouds}$/$\\rho_{\\rm b} < 9.9 \\times 10^{-5}$, $\\rho_{HIclouds}$/$\\rho_{\\rm HIG} < 8.9 \\times 10^{-3}$.  Since radio wavelengths suffer little absorption in the Galactic Plane, HIPASS provides a clear view of the local large-scale structure. The dominant features in the sky distribution of the BGC are the Supergalactic Plane and the Local Void. In addition, one can clearly see the Centaurus Wall which connects via the Hydra and Antlia clusters to the Puppis filament. Some previously hardly noticed galaxy groups stand out quite distinctively in the \\HI\\ sky distribution. Several new structures can be mapped for the first time, not only behind the Milky Way."
    },
    "0404/astro-ph0404600_arXiv.txt": {
        "abstract": "We use the LUQAS sample (Kim et al. 2004), a set of 27 high-resolution and high signal-to-noise QSO absorption spectra at a median redshift of $z=2.25$, and the data from Croft et al. (2002) at a median redshift of $z=2.72$, together with a large suite of high-resolution large box-size hydro-dynamical simulations, to estimate the linear dark matter power spectrum on scales $0.003\\,{\\rm s/km} < k < 0.03\\,{\\rm s/km}$.  Our re-analysis of the Croft et al.~data agrees well with their results if we assume the same mean optical depth and gas temperature-density relation.  The inferred linear dark matter power spectrum at $z=2.72$ also agrees with that inferred from LUQAS at lower redshift if we assume that the increase of the amplitude is due to gravitational growth between these redshifts.  We further argue that the smaller mean optical depth measured from high-resolution spectra is more accurate  than the larger value obtained from low-resolution spectra by Press et al.~(1993) which Croft et al.~used. For the smaller optical depth we obtain a $\\approx 20$\\% higher value for the rms fluctuation amplitude of the matter density. By combining the amplitude of the matter power spectrum inferred from the Ly$\\alpha$ forest with the amplitude on large scales inferred from measurements of the CMB we obtain constraints on the primordial spectral index $n$ and the normalisation $\\sigma_8$. For values of the mean optical depth favoured by high-resolution spectra, the inferred linear power spectrum is consistent with a $\\Lambda$CDM model with a scale-free ($n=1$) primordial power spectrum. ",
        "introduction": "The prominent absorption features blue-ward of the \\lya emission in the spectra of high-redshift quasars (QSOs) are now generally believed to arise from smooth density fluctuations of a photoionised warm intergalactic medium (see Rauch 1998 and Weinberg et al. 1999 for reviews).  This has opened up the possibility to probe the density fluctuation of matter with the flux power spectrum of QSO absorption lines (Croft et al. 1998, Hui 1999, Croft et al. 1999b, McDonald et al. 2000 [M00], Hui et al. 2001, Croft et al. 2002 [C02], McDonald 2003, Viel et al. 2003).  The flux power spectrum is mainly sensitive to the slope and amplitude of the linear dark matter power spectrum for wave-numbers in the range $0.002\\, \\vm <k< 0.05\\, \\vm$.  Croft et al. (1999) inferred an amplitude and slope which was consistent with a COBE normalised $\\Lambda$CDM model with a primordial scale invariant fluctuation spectrum (Phillips et al. 2001).  M00 and C02, using a larger sample of better quality data, found a somewhat shallower slope and smaller fluctuation amplitude.  The WMAP team used the later data in combination with their CMB data to claim that there is evidence for a tilted primordial CMB-normalised fluctuation spectrum ($n<1$) and/or a running spectral index (Bennet et al. 2003; Spergel et al. 2003; Verde et al. 2003).  A number of authors have argued that the errors in the inferred dark matter (DM) power spectrum have been underestimated (Zaldarriaga, Scoccimarro \\& Hui, 2003; Zaldarriaga, Hui \\& Tegmark, 2001; Gnedin \\& Hamilton 2002; Seljak, McDonald \\& Makarov 2003).  We here use a suite of high-resolution hydro-dynamical simulations and the flux power spectrum obtained from LUQAS (Large Sample of UVES QSO Absorption Spectra), together with the published flux power spectrum of C02 to further investigate these issues.   The plan of the paper is as follows. In Section \\ref{data} we describe the two data sets and give an overview of the determinations of the effective optical depth found in the literature.  The flux power spectra obtained from the hydro-dynamical simulations are discussed in Section \\ref{hydro}.  Section \\ref{inferring} describes the method and uncertainties of inferring the linear matter power spectrum.  In Section \\ref{linear} we present our results, compare with previous results and discuss implications for $\\sigma_8$ and $n$. Section \\ref{conclu} contains a summary and our conclusions. ",
        "conclusions": "\\label{conclu}  We have used the observed flux power spectrum of two large samples of high-resolution spectra, a sample drawn from LUQAS at a median redshift of $z=2.125$, and the sample compiled by Croft et al. (2002) at $z=2.72$, together with a suite of high-resolution numerical simulations, to infer the dark matter power spectrum on scales $0.003< k $(s/km)$<0.03$.  We have obtained the following results: \\begin{enumerate} \\item{With the same assumptions for effective optical depth, density-temperature relation, and cosmology, our inferred linear matter power spectrum agrees very well with that inferred by Croft et al.~(2002).}  \\item{We confirm previous results that the inferred rms amplitude of density fluctuations depends strongly on the assumed $\\tau_{\\rm eff}$. It increases by 20\\% if we assume an optical depth of $\\tau_{\\rm eff}= 0.305$ a value suggested by studies of high-resolution absorption spectra. We find, however, a dependence on $\\tau_{\\rm eff}$ which is weaker than that of C02 and stronger than that of GH02.}  \\item{For values $\\tau_{\\rm eff}$ suggested by high-resolution absorption spectra the linear power spectrum of the best fitting running spectral index model of Spergel et al. (2003) falls significantly below the linear power spectrum inferred from both the LUQAS and the C02 sample}.  \\item{The decrease of the amplitude of the flux power spectrum between $z=2.7$ and $z=2.1$ is consistent with that expected due to the decrease of $\\tau_{\\rm eff}$ and the increase of the amplitude of matter power spectrum due to gravitational growth.} \\item{Our estimate of the systematic uncertainty of the rms fluctuation amplitude of the density ($\\sim 14.5 \\%$) is a factor 3.5 larger than our estimate of the statistical error ($\\sim 4 \\%$).  The systematic uncertainty is dominated by the uncertainty in the mean effective optical depth and --- somewhat surprisingly --- by the uncertainties between the numerical simulations of different authors. Reducing  the overall errors will thus mainly rely on a better understanding of a range of systematic uncertainties.} \\item{By combining the CMB constraint (assuming that there is no contribution from tensor fluctuations) on the amplitude of the DM power spectrum on large scale with the high-resolution \\lya forest data we obtain $n= 1.01 \\; (\\Omega_{\\rm 0m}h^2/0.135)^{-0.35}\\; \\pm 0.02 {(\\rm statistical)}\\pm 0.06{( \\rm systematic)}$ for the spectral index. The corresponding rms fluctuation amplitude is, $\\sigma_8 = 0.93 \\pm 0.03 {(\\rm statistical)}\\pm 0.09{(\\rm systematic)}.$} \\end{enumerate}"
    },
    "0404/hep-ph0404103_arXiv.txt": {
        "abstract": "\\vskip 1cm We explore the potential of measuring the solar neutrino oscillation parameters in the proposed gadolinium loaded Super-Kamiokande (SK-Gd) detector. Gadolinium dissolved in water can detect neutrons much more efficiently than pure water. This imparts the detector the ability to observe electron type antineutrinos, transforming Super-Kamiokande into a huge reactor antineutrino detector with an event rate approximately 43 times higher than that observed in KamLAND. We simulate the reactor antineutrino data expected in this high statistics detector. We use these prospective data to study the precision with which the solar neutrino oscillation parameters, $\\Delta m^2_{\\odot}$ and $\\sin^2\\theta_{\\odot}$, can be determined i) with the SK-Gd detector, and ii) by combining the SK-Gd data with the global data on solar neutrino oscillations. For comparison and completeness the allowed regions of $\\Delta m^2_{\\odot}$ and $\\sin^2\\theta_{\\odot}$, expected to be obtained from the data of the solar neutrino and KamLAND experiments, are also presented. We find that the SK-Gd experiment could provide one of the most precise (if not the most precise) determination of the solar neutrino oscillation parameters $\\Delta m^2_{\\odot}$ and $\\sin^2\\theta_{\\odot}$. ",
        "introduction": "\\label{section:introduction}  At present we have compelling evidences for oscillations of both solar ($\\nu_e$) and atmospheric ($\\nu_{\\mu}$ and $\\bar{\\nu}_{\\mu}$) neutrinos, driven by nonzero neutrino masses and neutrino mixing \\cite{BPont57,Pont67}. They have been accumulated over a long period of time in the experiments with solar neutrinos \\cite{Cl98,Kamiokande,SAGE,GALLGNO,SKsol,SNO1,SNO2,SNO3}, Homestake, Kamiokande, SAGE, GALLEX/GNO, Super-Kamiokande and SNO, in the KamLAND experiment with reactor antineutrinos \\cite{KamLAND}, and in the studies of the fluxes of atmospheric neutrinos by the Super-Kamiokande collaboration \\cite{SKatmo98,SKatmo03,SKdip04}. Indications for oscillations of neutrinos have also been obtained in the K2K long base-line experiment \\cite{K2K}.  Evidences for oscillations of the solar $\\nu_e$, induced by nonzero neutrino masses and neutrino mixing \\cite{Pont67}, have been reported first by the pioneering Davis et al. (Homestake) experiment \\cite{Pont46,Davis68}. They have been confirmed and reinforced later by Kamiokande, SAGE, GALLEX/GNO and Super-Kamiokande experiments. The evidences for mixing and oscillations of the solar $\\nu_e$ have been made compelling during the last three years by the data from the SNO solar neutrino and KamLAND reactor antineutrino experiments \\cite{SNO1,SNO2,SNO3,KamLAND} (see also, e.g., \\cite{snonc}). Under the assumption of CPT-invariance, the observed disappearance of reactor $\\bar{\\nu}_e$ in the KamLAND experiment, in particular, confirmed the interpretation of the solar neutrino data in terms of $\\nu_e \\rightarrow \\nu_{\\mu,\\tau}$ oscillations, induced by nonzero neutrino masses and nontrivial neutrino mixing. The KamLAND results \\cite{KamLAND} practically established the large mixing angle (LMA) MSW solution (see, e.g., \\cite{SNO1}) as unique solution of the solar neutrino problem. The solar neutrino and KamLAND data, including the salt phase SNO results, strongly favor the low-LMA solution with \\cite{SNO3,salt,saltother} $\\Delta m^2_{21} \\sim 7\\times 10^{-5}$ eV$^2$ and $\\sin^2\\theta_{12} \\sim 0.3$, $\\ms \\equiv \\Delta m^2_{\\odot}$ and $\\theta_{12} \\equiv \\theta_{\\odot}$ being the neutrino mass squared difference and mixing angle driving the solar neutrino oscillations. The high-LMA solution (see, e.g., \\cite{KamLAND}), characterized by $\\Delta m^2_{21} \\sim 1.5\\times 10^{-4}$ eV$^2$ and similar value of $\\sin^2\\theta_{12}$, is only allowed at 99.14\\% C.L. by the data \\cite{salt}.   Strong evidences for oscillations of the atmospheric $\\nu_{\\mu}$ ($\\bar{\\nu}_{\\mu}$) have been obtained in the Super-Kamiokande experiment from the observed i) Zenith angle dependence of the multi-GeV and sub-GeV $\\mu$-like events \\cite{SKatmo98,SKatmo03}, and ii) the recently reported observation of an ``oscillation dip'' in the $L/E-$dependence of the (essentially multi-GeV) $\\mu-$like atmospheric neutrino events \\cite{SKdip04} \\footnote{The sample used in the analysis of the $L/E$ dependence consists of $\\mu-$like events for which the relative uncertainty in the experimental determination of the $L/E$ ratio does not exceed 70\\%.}, $L$ and $E$ being the distance traveled by neutrinos and the neutrino energy. As is well known, the SK atmospheric neutrino data is best described in terms of dominant two-neutrino $\\nu_{\\mu} \\rightarrow \\nu_{\\tau}$ ($\\bar{\\nu}_{\\mu} \\rightarrow \\bar{\\nu}_{\\tau}$) vacuum oscillations  with \\cite{SKatmo03,SKdip04} $\\ma \\equiv \\Delta m^2_{31} \\sim (2.0 - 3.0)\\times 10^{-3}$ eV$^2$ and maximal mixing, $\\sin^22\\theta_{\\rm atm} \\equiv \\sin^22\\theta_{23}\\sim 1.0$. The observed dip is predicted due to the oscillatory dependence on $L/E$ of the $\\nu_{\\mu} \\rightarrow \\nu_{\\tau}$ and $\\bar{\\nu}_{\\mu} \\rightarrow \\bar{\\nu}_{\\tau}$ oscillation probabilities. This beautiful result represents the first ever observation of a direct effect of the oscillatory dependence on $L/E$ of the probability of  neutrino oscillations in vacuum.  Information on the third mixing angle $\\theta_{13}$, present in the case of 3-neutrino mixing which is required to describe the solar and atmospheric neutrino (and KamLAND) data in terms of neutrino oscillations, is provided essentially by the short baseline reactor antineutrino experiments CHOOZ and Palo Verde \\cite{CHOOZPV}. Using the range of allowed values of $\\ma$ found in the latest SK data analysis \\cite{SKatmo03,Fogliatm0308055}, a combined 3-$\\nu$ oscillation analysis of the solar neutrino, CHOOZ and KamLAND data gives (at 99.73\\% C.L.) $\\sin^2 \\theta_{13} < 0.074$ \\cite{salt}. A somewhat more stringent limit was obtained in the global analysis of all the available solar, atmospheric and reactor neutrino data, performed in \\cite{ConchaNOON04}: $\\sin^2 \\theta_{13} < 0.050$ (99.73\\% C.L.).  Neutrino flavor oscillations have also been claimed to have been observed in the LSND experiment \\cite{lsnd}. This interpretation of the LSND results is being currently tested in the MiniBOONE experiment \\cite{mboone}.  After the remarkable progress made in the last few years in establishing the existence of neutrino oscillations, one of the main goals of the future experimental studies of neutrino mixing is to measure with high precision the parameters which drive the solar and atmospheric neutrino oscillations, $\\Delta m^2_{21}$, $\\Delta m^2_{31}$, $\\sin^2\\theta_{12}$ and $\\sin^2\\theta_{23}$. The potential of the current and future experiments for high precision determination of the solar neutrino oscillation parameters, $\\Delta m^2_{21}$ and $\\sin^2\\theta_{12}$, was studied recently in \\cite{th12,th12jnb,th12hlma,shika} (see also \\cite{talks}). The precision in the measurement of $\\Delta m^2_{21}$ and $\\sin^2\\theta_{12}$, which can be  reached in the Borexino and the so-called LowNu experiments, was explored in \\cite{th12,th12jnb}. The impact of the prospective increase in statistics of the \\kl data on the determination of $\\Delta m^2_{21}$ and $\\sin^2\\theta_{12}$ was investigated in detail in \\cite{salt,th12,shika}. If the $e^+-spectrum$ measured at KamLAND is simulated at a point in the low-LMA region, the allowed 3$\\sigma$ area in the high-LMA zone reduces in size in the case of 0.41 kTy of data, and disappears if the statistics is increased to 1.0 kTy \\cite{salt}. In this case $\\Delta m^2_{21}$ will be determined with high precision. If, however, the spectrum observed in \\kl conforms to a point in the high-LMA region, the conflicting trend of solar and \\kl data would make the high-LMA solution reappear at 90\\% C.L. and the determination of $\\Delta m^2_{21}$ would remain ambiguous \\cite{salt}. In both cases $\\sin^2\\theta_{12}$ cannot be determined with high precision. Actually, reaching high accuracy in the measurement of $\\sin^2\\theta_{12}$ (e.g., 10\\% error at 1$\\sigma$) is a rather challenging problem. It can be solved by performing a dedicated reactor $\\bar{\\nu}_e$ experiment \\cite{th12,th12hlma}, or by the future solar neutrino experiments LENS and XMASS, aiming to measure the pp neutrino flux \\cite{LENSNOON04,XMASSNOON04}.  Recently it was proposed in \\cite{gdpaper} to dope the SK detector with Gadolinium by dissolving 0.2\\% gadolinium trichloride in the water \\footnote{The authors of \\cite{gdpaper} called  ``GADZOOKS!'' the resulting detector. We will use in what follows the abbreviation SK-Gd for it.}. The added gadolinium would make it possible to detect the neutrons, released in the $\\anue$ capture on protons, with a relatively high efficiency. This would allow the SK experiment, in particular, to detect the reactor antineutrinos coming from the numerous powerful nuclear reactors located in Japan, thus transforming SK into a huge reactor antineutrino detector with an event rate approximately 43 times higher than that observed in KamLAND. In this paper we study the prospects of measuring the solar neutrino oscillation parameters by observing the reactor  antineutrino oscillations in the proposed gadolinium loaded SK (SK-Gd) experiment. Since $\\Delta m^2_{21} \\ll \\Delta m^2_{31}$ and the mixing angle $\\theta_{13}$ is restricted to be relatively small, the third (heaviest) neutrino with definite mass is expected to have  negligible impact on the determination of the values of the solar neutrino oscillation parameters \\footnote{The effect of  $\\sch$ can become important if $\\sin^22\\theta_{12}$ is measured with a $\\sim 10\\%$ precision, which, as we will see, is higher than the precision which can be reached in the SK-Gd experiment.} in the experiment of interest and we perform our study in the framework of the 2$\\nu-$ mixing scenario ($\\sch = 0$).  We begin in Section \\ref{solklsec} with an overview of the currently allowed ranges of values of the solar neutrino oscillation parameters by the existing global solar neutrino and KamLAND data. We present further the regions in the $\\Delta m^2_{21} - \\sin^2\\theta_{12}$ plane, which are expected to be allowed after taking into account future higher statistics and lower systematic data from SNO and from KamLAND experiments. The SK-Gd detector is considered  in Section \\ref{sens1}, where we discuss the assumptions made to simulate the data in this proposed modification of the SK experiment, as well as the procedure used to statistically analyze the simulated data. We investigate further the potential of the SK-Gd experiment in reducing the uncertainties in the values of $\\ms$ and $\\sss$. In Section \\ref{solgdsec} the expected results from a combined analysis of the solar neutrino data and the prospective SK-Gd simulated spectrum data are given. In Sections \\ref{sens1} and \\ref{solgdsec} we also discuss the bounds on $\\ms$ and $\\sss$ one would obtain if some of the reactors in Japan would be ``switched off''. Section \\ref{concl} contains the main conclusions of our study. ",
        "conclusions": "\\label{concl}  In this paper we have studied the prospects of high precision determination of the solar neutrino oscillation parameters, $\\ms$ and $\\sss$, by investigating the reactor  $\\bar{\\nu}_e$ oscillations with the recently proposed gadolinium loaded Super-Kamiokande (SK-Gd) experiment \\cite{gdpaper}. Doping the  SK detector with Gadolinium by dissolving 0.2\\% gadolinium trichloride in the water \\cite{gdpaper} would make it possible to detect the neutrons, released in the $\\anue$ capture on protons, with a relatively high efficiency. This would allow the SK experiment, in particular, to detect the reactor $\\bar{\\nu}_e$ coming from the numerous powerful nuclear reactors located in Japan, thus transforming SK into the largest long baseline reactor $\\bar{\\nu}_e$ detector with an event rate approximately 43 times higher than that observed in the KamLAND experiment.  Working in the framework of two-neutrino mixing, we have presented first an overview of the currently allowed ranges of values of the solar neutrino oscillation parameters $\\ms$ and $\\sss$ by the existing global solar neutrino and KamLAND data (Fig. 1 and Table 1), as well as the regions in the parameter space which are expected to be allowed after taking into account future higher statistics and lower systematic data from SNO and from KamLAND experiments (Fig. 2). The SK-Gd detector is considered  in Section \\ref{sens1}, where we discuss the assumptions made to simulate the data in this proposed modification of the SK detector, as well as the procedure used to statistically analyze the simulated data.  The results of our analysis show that the SK-Gd experiment has a remarkable potential in reducing the uncertainties in the values of $\\ms$ and $\\sss$ (Figs. 3 and 7 and Table 1). Combining the SK-Gd data, taken over a period of 3 years, with the data from the solar neutrino experiments, will allow to determine $\\ms$ with a 3\\% error at 99\\% C.L., the error being by a factor of 2 smaller if one uses only the SK-Gd data. With only $\\sim 0.5$ years of data, the SK-Gd experiment would rule out completely the spurious high-LMA solution if low-LMA is the true solution. In contrast, the high-LMA solution could be ruled out if one uses the data from the solar neutrino experiments and data from the \\kl detector, corresponding to 1kTy of statistics. If the future \\kl data would conform to a point in the high-LMA region, both the high-LMA and the low-LMA solutions would be allowed at 90\\% C.L. \\cite{salt}. In this case the high-LMA -- low-LMA solution ambiguity can be resolved completely by the SK-Gd experiment. In what regards the solar neutrino mixing angle, using the data from the solar neutrino experiments and the SK-Gd 3 year data will allow to determine $\\sss$ with a $\\sim 15\\%$ error at 99\\% C.L. One would have the same uncertainty in $\\sss$ using data only from the SK-Gd experiment, corresponding to 5 years of measurements. Thus, the SK-Gd experiment could provide one of the most precise (if not the most precise) determination of the solar neutrino oscillation parameters $\\ms$ and $\\sss$."
    },
    "0404/astro-ph0404422_arXiv.txt": {
        "abstract": "We analyse near-infrared integral field spectroscopy of the central starburst region of NGC 1140, obtained at the Gemini-South telescope equipped with CIRPASS. Our $\\sim 1.45 - 1.67 \\mu$m wavelength coverage includes the bright [Fe {\\sc ii}]$\\lambda 1.64 \\mu$m emission line, as well as high-order Brackett (hydrogen) lines. While strong [Fe {\\sc ii}] emission, thought to originate in the thermal shocks associated with supernova remnants, is found throughout the galaxy, both Br 12--4 and Br 14--4 emission, and weak CO(6,3) absorption, is predominantly associated with the northern starburst region. The Brackett lines originate from recombination processes occurring on smaller scales in (young) H{\\sc ii} regions. The time-scale associated with strong [Fe {\\sc ii}] emission implies that most of the recent star-formation activity in NGC 1140 was induced in the past $\\sim 35-55$ Myr. Based on the spatial distributions of the [Fe {\\sc ii}] versus Brackett line emission, we conclude that a galaxy-wide starburst was induced several tens of Myr ago, with more recent starburst activity concentrated around the northern starburst region.\\\\ This scenario is (provisionally) confirmed by our analysis of the spectral energy distributions of the compact, young massive star clusters (YMCs) detected in new and archival broad-band {\\sl Hubble Space Telescope} images. The YMC ages in NGC 1140 are all $\\lesssim 20$ Myr, consistent with independently determined estimates of the galaxy's starburst age, while there appears to be an age difference between the northern and southern YMC complexes in the sense expected from our CIRPASS analysis. Our photometric mass estimates of the NGC 1140 YMCs, likely upper limits, are comparable to those of the highest-mass Galactic globular clusters and to spectroscopically confirmed masses of (compact) YMCs in other starburst galaxies. Our detection of similarly massive YMCs in NGC 1140 supports the scenario that such objects form preferentially in the extreme environments of interacting and starburst galaxies. ",
        "introduction": "\\label{intro.sec}  The southern starburst dwarf galaxy NGC 1140 (Mrk 1063) is remarkable in the number of compact, young, massive star clusters (YMCs) that it has formed recently in the small volume of its bright emission-line core (roughly $600 \\times 200$ pc at the galaxy's distance\\footnote{based on the galaxy's heliocentric velocity, corrected for the Virgocentric flow, and assuming H$_0 = 70$ km s$^{-1}$ Mpc$^{-1}$; adopted from the {\\sl HyperLeda} database at {\\tt http://www.univ-lyon1.fr/}} of $(m-M)_{\\rm kin} = 31.49$ mag or $D = 20$ Mpc). Based on H{\\sc i} and optical observations, and comparison with models (Hunter, van Woerden \\& Gallagher 1994b, hereafter H94b), it exhibits the characteristics of a recent merger event, which has presumably induced the YMC formation in its centre; the galaxy is unusually bright ($M_V^0 = -20.76$; de Vaucouleurs et al. 1991, hereafter RC3) for its dwarf (or amorphous) morphological type, and shows low-level shells and filaments in its H{\\sc i} distribution (H94b), which are presumably remnants of a past interaction.  Within 3.3 Holmberg radii $(R_{\\rm H}$, the radius where the {\\it V}-band surface brightness reaches $\\mu_V = 26.4$ mag arcsec$^{-2}$; at $\\mu_V = 25$ mag arcsec$^{-2}$ the galaxy measures about $9.9 \\times 5.2$ kpc), it contains a total mass of $1 \\times 10^{11} {\\rm M}_\\odot$, based on its H{\\sc i} rotation curve; $7.5 \\times 10^9 {\\rm M}_\\odot$ of this mass is in the form of H{\\sc i}. The total mass contained within $1 R_{\\rm H}$ is $1.4 \\times 10^{10} {\\rm M}_\\odot$ (H94b).  The formation of compact YMCs seems to be a key feature of intense episodes of recent, active star formation. They have been identified, mostly with the {\\sl Hubble Space Telescope (HST)}, in several dozen galaxies, often interacting or undergoing active starbursts (see, e.g., Holtzman et al. 1992, O'Connell, Gallagher \\& Hunter 1994, de Grijs, O'Connell \\& Gallagher 2001, Whitmore 2003, and references therein). Their sizes, luminosities, and masses -- usually based on broad-band optical observations -- are generally consistent with theoretical expectations for young globular clusters (see also de Grijs, Bastian \\& Lamers 2003b). However, they cover a wide range of masses and luminosities, and also include those of the most extreme supermassive and superluminous clusters. YMCs are also important as probes of their host galaxy's star formation history, its chemical evolution, the stellar initial mass function (IMF), and other physical characteristics in starbursts.  This is so because each cluster approximates a coeval, single-metallicity, simple stellar population. Such systems are among the simplest to model, and their ages and metallicities and, in some cases, IMFs can be estimated from their integrated spectra (Ho \\& Filippenko 1996a,b, Gallagher \\& Smith 1999, Smith \\& Gallagher 2001, Mengel et al. 2002).  \\subsection{NGC 1140: A disturbed appearance across the board}  Observations in optical broad-band, H$\\alpha$ narrow-band, and the H{\\sc i} radio continuum are consistent with the interpretation that NGC 1140 is in the final stages of a galactic merger (e.g., Hunter et al. 1994a, hereafter H94a; H94b). At faint optical light levels it shows multiple, misaligned shell-like structures. These are reminiscent of the shells associated with some elliptical galaxies that are thought to be remnants of a past tidal encounter (e.g., Schweizer \\& Seitzer 1988, Hernquist \\& Spergel 1992). The galaxy is also unusual in its content and overall morphology: its optical appearance is dominated by a supergiant H{\\sc ii} region encompassing most of its centre, with an H$\\alpha$ luminosity equivalent to that produced by $\\sim 10^4$ OB stars (H94a), far exceeding that of the giant H{\\sc ii} region 30 Doradus in the Large Magellanic Cloud (LMC; H94b). Its gas fraction, gas-to-luminosity ratio, and H{\\sc i} velocity dispersion are unusually high, and its H$\\alpha$ velocity profiles remarkably broad, by roughly an order of magnitude, for the mid-spiral type galaxy one would expect, based on its total mass (H94b).  Although NGC 1140 has a small dwarf companion galaxy (which might suggest that gravitational interactions with this companion may have induced part of the recent active star cluster formation), these observations led H94b to conclude that the galaxy is most likely the product of the accretion of a low surface brightness gas-rich companion by a relatively normal mid-type spiral galaxy of roughly equal mass in the last $\\sim 1$ Gyr. This picture is consistent with the low metallicities (see Section \\ref{hst1.sec}) and complex H{\\sc i} morphology (including a warped disc) and velocity field.  \\subsection{The massive star cluster population}  The presence and characteristics of the central supergiant H{\\sc ii} region in NGC 1140 are consistent with simple conceptual models of enhanced central star formation associated with a tidal encounter (see H94b). Its core contains six blue, very luminous compact YMCs\\footnote{H94a show convincingly that these objects are much more likely to be YMCs than either luminous red supergiants or objects similar to the most luminous unstable stars} -- with ages $\\lesssim 15$ Myr (H94a) -- such as commonly found in regions of intense star formation.  These are comparable to or brighter than Galactic globular clusters (GCs) at similar ages, and thought to be entirely responsible for the ionisation of the supergiant H{\\sc ii} region in NGC 1140 (H94a). It appears that most of the recent star formation in NGC 1140 has occurred in, or is associated with, the YMCs: while its overall average recent star formation rate, $\\dot{M} \\sim (1-6) {\\rm M}_\\odot$ yr$^{-1}$ (based on a variety of tracers of the recent star formation rate and depending on -- among others -- the slope and high-mass cut-off assumed for the IMF; e.g., Storchi-Bergmann, Calzetti \\& Kinney 1994, H94b, Calzetti 1997, Takagi, Arimoto \\& Hanami 2003), is similar to the equivalent rates of the significantly more massive giant irregular and spiral galaxies, the average recent star formation rate {\\em in the YMCs alone} is implied to be near 1 M$_\\odot$ yr$^{-1}$ (H94a), if their IMFs are similar to the solar neighbourhood IMF.  In Section \\ref{observations.sec} we describe in detail both the broad-band {\\sl HST} imaging observations of the few dozen compact YMCs (i.e., the six very bright clusters discussed above, in addition to a larger population of less luminous compact star clusters) in NGC 1140's central starburst, and our near-infrared (NIR) integral-field spectroscopy. We then apply sophisticated age dating techniques to the {\\sl HST} observations in Section \\ref{hst.sec} and discuss the spectroscopic evidence for a possible age gradient (see Section \\ref{cirpass.sec}). These results are discussed in the context of NGC 1140's recent star formation history and its current supernova rate (SNR) in Sections \\ref{context.sec} and \\ref{snr.sec}, respectively, upon which we summarise our main results and conclusions in Section \\ref{summary.sec}. ",
        "conclusions": "\\label{summary.sec}  We have analysed NIR integral field spectroscopy of the nearby southern starburst galaxy NGC 1140, obtained at the Gemini-South telescope equipped with CIRPASS. The wavelength coverage used for this programme, $0.22 \\mu$m from $\\sim 1.45 - 1.67 \\mu$m, includes the bright [Fe {\\sc ii}]$\\lambda 1.64 \\mu$m emission line, as well as the high-order Brackett lines Br 12--4 and Br 14--4, which are all clearly resolved in the NGC 1140 central starburst region, and located well away from contaminating OH sky lines. Additional {\\sl HST} imaging observations covering the UV/optical range have been used to place these CIRPASS results in the context of interaction-induced star cluster formation and evolution.  NGC 1140 is remarkable in the number of compact YMCs formed in its bright, compact emission-line core, which were most likely triggered by a recent merger event with a gas-rich companion galaxy. The total energy output of these YMCs far exceeds that of the local ``super'' star cluster R136 in the core of the 30 Doradus starburst region in the LMC. Our CIRPASS spectroscopy, combined with the broad-band {\\sl HST} images, probes the complex transitions between the YMCs and the galaxy's ISM, which we use to test predictions of models for radiative and supernova-induced shocks.  The flux distribution of the [Fe {\\sc ii}] line is clearly more extended than those of the Brackett lines. While [Fe {\\sc ii}] emission is found throughout the galaxy, both Br 12--4 and Br 14--4 emission is predominantly associated with the northern starburst region. It is now well-established that the {\\it strong} [Fe {\\sc ii}]$\\lambda 1.64 \\mu$m line emission in starbursts originates predominantly in SNRs. The Brackett lines, on the other hand, are formed by recombination in ionised gas associated with sources having a strong Lyman continuum, such as H{\\sc ii} regions created by young, massive stars. The [Fe {\\sc ii}] emission and the young region as a whole may, in fact, represent the broad base of a galactic outflow, or wind, in its early stages.  The presence of strong [Fe {\\sc ii}] emission in NGC 1140's core, indicative of SNRs, can help to set limits on its star formation history. The last SNe in a quenched starburst region occur at a time comparable to the longest lifetime of an SN progenitor after the end of the starburst activity. Thus, for NGC 1140 we conclude that the presence of strong [Fe {\\sc ii}] emission throughout the central galaxy implies that most of the recent star-formation activity was induced in the past $\\sim 35-55$ Myr.  The Br 12--4 and Br 14--4 emission, concentrated around NGC 1140-N, is most likely associated with recombination in (younger) H{\\sc ii} regions. This implies, therefore, that the most recent starburst event has transpired in this more confined area of the NGC 1140 disc, around the northern starburst region and the massive, compact YMCs [H94a]1--3.  This scenario is provisionally confirmed by our fits based on broad-band {\\sl HST} imaging observations (despite the large associated uncertainties), and by the preliminary analysis of H94a. Using the method developed in Anders et al. (2004a) and de Grijs et al. (2003c), we derive the YMC ages, masses, and extinction values simultaneously from the {\\sl HST} broad-band passband coverage; we assume a generic metallicity for all NGC 1140 YMCs of $0.4 Z_\\odot$, which we justify by emphasizing technical limitations of our metallicity fits at the young ages expected for the compact NGC 1140 YMCs. Such a metallicity is also supported by previous, independent estimates. The ages derived for the NGC 1140 YMCs are all $\\lesssim 20$ Myr (or perhaps significantly younger, in view of the uncertainties inherent to our method of analysis of broad-band SEDs), consistent with independently determined estimates of the galaxy's starburst age, while there appears to be an age difference between YMC complexes [H94a]1--3 and [H94a]6--7 in the sense expected from our CIRPASS analysis. Our photometric mass estimates of the NGC 1140 YMCs, likely representing upper limits, are comparable to the Galactic GCs comprising the high-mass wing of the GC mass distribution and to the spectroscopically confirmed masses of the compact YMCs in M82 and the Antennae galaxies. Our detection of similarly massive YMCs in NGC 1140 supports the scenario that such objects form preferentially in the extreme environments of interacting and starburst galaxies.  Finally, we derive lower limits to the SN rates in the NGC 1140 starburst regions S and N of $\\sim 0.3$ SN yr$^{-1}$, within a factor of $\\sim 2$, which implies a similar starburst vigour as seen in the actively starbursting nucleus of the nearby prototype starburst galaxy M82. The two spatially-confined starburst sites S and N emit $\\sim 40$ per cent of the integrated [Fe {\\sc ii}] flux in the galaxy, with the remainder being distributed in a diffuse component. The high-order Brackett lines, on the other hand, exhibit a negligible diffuse component, but are confined to starburst region N."
    },
    "0404/astro-ph0404108_arXiv.txt": {
        "abstract": "{We study an example of instability in presence of a multiplicative noise, namely the spontaneous generation of a magnetic field in a turbulent medium. This so-called turbulent dynamo problem remains challenging, experimentally and theoretically. In this field, the prevailing theory is the Mean-Field Dynamo \\cite{Krause80} where the dynamo effect is monitored by the mean magnetic field (other possible choices would be the energy, flux,...). In recent years, it has been shown on stochastic oscillators that this type of approach could be misleading. In this paper, we develop a stochastic description of the turbulent dynamo effect which permits us to define unambiguously a threshold for the dynamo effect, namely by globally analyzing the probability density function of the magnetic field instead of a given moment. }   \\PACS{ {02.50.-r } {Probability theory, stochastic processes and statistics} \\and {47.27.Gs } {Isotropic turbulence; homogeneous turbulence} \\and {47.27.Jv} {High-Reynolds-number turbulence} } ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404278_arXiv.txt": {
        "abstract": "{We compare recent HST observations of Seyfert and quasar NLRs and find that type--2 AGNs follow a relation consistent with that expected for a distribution of gas ionized by a central source $R_{\\rm NLR,2} \\propto L^{0.32 \\pm 0.05}$, while type--1 objects are fit with a steeper slope of $0.55 \\pm 0.05$. The latter is comparable to the scaling found for the BLR size with continuum luminosity (slope: 0.5--0.7). Therefore, we investigate what we can learn about the BLR size if the NLR size is only determined by the AGN luminosity. We  find that NLR and BLR size are related linearly following $R_{\\rm BLR} \\propto R_{\\rm NLR,1}^{0.88 \\pm 0.1}$. This relation can be used to estimate BH masses.}  \\firstsection ",
        "introduction": " ",
        "conclusions": "Sizes and morphologies of broad-- and narrow--line regions (BLRs, NLRs) provide an ideal probe of the distribution of dust and gas in the central parts of AGNs. However, until today, even very basic issues remain open: What determines the size and structure of these emission--line regions? Do they grow with luminosity, and if so, how? For the NLR, the latter was investigated for the first time by \\cite[Bennert \\etal\\ (2002)]{ben02} for a sample of seven radio--quiet PG quasars observed in the [OIII]\\,$\\lambda$5007 line with HST. We found a NLR size--luminosity relation $R_{\\rm NLR} \\propto L_{\\rm [OIII]}^{0.5}$ when including a sample of seven Seyfert--2 galaxies (\\cite[Falcke \\etal\\ 1998]{fal98}). This result is remarkable if it implies that it is not the Str\\\"omgren radius that limits the NLR but an apparent threshold in ionizing flux, which can be expressed by a constant product of ionization parameter and density at the rim of the NLR. It is not clear whether this new relation will hold at all luminosities and redshifts. If, for example, the luminosity--scaled size of the NLR becomes larger than the size of the galaxy, the relation may flatten as the emission lines fade out and disappear (\\cite[Netzer 2004]{net04}). This may be evident already in some spectral line studies of Two Degree Field quasars (\\cite[Croom \\etal\\ 2002]{cro02}).  More recently, \\cite[Schmitt \\etal\\ (2003)]{sch03} studied the extended [OIII] emission in a sample of 60 Seyfert galaxies and find that the NLR size--luminosity relation follows a simple Str\\\"omgren law $R_{\\rm NLR} \\propto L_{\\rm [OIII]}^{0.33}$. What causes the apparent different slopes for our quasar--dominated sample (0.5) and their Seyfert sample (0.33): simple statistical uncertainty or selection effects? Are the NLRs of Seyferts and quasars intrinsically different? Or is there a difference in the NLR size--luminosity relation between type--1 and type--2 objects? Puzzled by these questions, we compared the two samples and calculated a fit to all type--1 and type--2 AGNs separately: While type--1 objects follow the relation $R_{\\rm NLR,1} \\propto L^{0.55 \\pm 0.05}$, type 2s can be fit by approximately a Str\\\"omgren--law $R_{\\rm NLR,2} \\propto L^{0.32 \\pm 0.05}$ (Fig. 1, {\\it left}). This result explains the different slopes found by the two groups: In the \\cite[Bennert \\etal\\ (2002)]{ben02} sample, basically type--1 quasars define the slope, whereas the slope of the \\cite[Schmitt \\etal\\ (2003)]{sch03} sample is dominated by Seyfert 2s which outnumber the Seyfert 1s by a factor of 1.4.  While the NLR size can be easily measured from images, the BLR is too compact to derive its size directly. Instead, observations of correlated variations in the line and continuum emission are used to determine the BLR size indirectly by the reverberation mapping techniques (e.g. \\cite[Wandel \\etal\\ 1999]{wan99}; \\cite[Kaspi \\etal\\ 2000]{kas00}). Investigations of the BLR size--luminosity relationship lead to controversy about the slope: While \\cite[Kaspi \\etal\\ (2000)]{kas00} report a relation of $R_{\\rm BLR} \\propto L^{0.7}$, \\cite[McLure \\& Jarvis (2002)]{mcl02} find $R_{\\rm BLR} \\propto L^{0.5}$. If the NLR size is only determined by the AGN luminosity, it is of interest what we can learn about the BLR size and BH mass. Therefore, we compared all measured BLR sizes with their corresponding NLR sizes. Unfortunately, the overlap is rather poor and leaves us with 11 objects. Applying a weighted linear least--squares fit, we find that the sizes of NLR and BLR are proportional: $R_{\\rm BLR} \\propto R_{\\rm NLR,1}^{0.88 \\pm 0.1}$ (Fig. 1, {\\it right}).  \\begin{figure} \\centering \\resizebox{6.5cm}{!}{\\includegraphics{bennert1.eps} } \\resizebox{6.5cm}{!}{\\includegraphics{bennert2.eps} } \\caption[]{{\\it Left:} Radius of the NLR versus the emission--line luminosity in [OIII] on logarithmic scales. The solid line represents the fit to type--1 AGNs ($R_{\\rm Maj}$ (NLR,1) $\\propto L_{\\rm [OIII]}^{0.55 \\pm 0.05}$), the dashed line the fit to type--2 AGNs ($R_{\\rm Maj}$ (NLR,2) $\\propto L_{\\rm [OIII]}^{0.32 \\pm 0.05}$). {\\it Right:} Distribution of NLR and BLR radii on logarithmic scales. The error bars indicate the uncertainty in defining the BLR radius. The solid line represents the fit corresponding to $R$ (BLR) $\\propto R_{\\rm Maj}$ (NLR,1)$^{0.88 \\pm 0.1}$.} \\end{figure}  Since deriving BLR sizes via reverberation mapping is very time consuming, relationships with luminosity or NLR size can be extremely useful for estimating BH masses. Given the NLR size (either directly measured or derived from the [OIII] luminosity), the BLR size can be estimated, which, assuming Keplerian motions, yields $M_{\\rm BH}$: \\begin{eqnarray*} M_{\\rm BH} = (10^{5.22} M_{\\odot}) v_{3000}^2 R_{\\rm NLR,1}^{0.88} \\hspace*{0.5cm} {\\rm or} \\hspace*{0.5cm} M_{\\rm BH} = (10^{8.83} M_{\\odot}) v_{3000}^2 L_{44, \\rm [OIII]}^{0.48} \\end{eqnarray*} (where $v_{3000}$ = FWHM (H$\\beta$)/3000 km\\,s$^{-1}$, $R_{\\rm NLR,1}$ in parsec, and $L_{44, \\rm [OIII]}$ =  [OIII] luminosity/$10^{44}$ erg\\,s$^{-1}$). The latter equation is very similar to the correlation of $M_{\\rm BH}$ and continuum luminosity at 5100 \\AA~(\\cite[Shields \\etal\\ 2003]{shi03}).  The observed slopes of the NLR size--luminosity and the BLR--NLR--size relation can be explained by a simple model based on the unified scheme and the receding torus model, taking into account the NLR size dependency on viewing angle and luminosity (\\cite[Bennert \\etal\\ 2004]{ben04})."
    },
    "0404/astro-ph0404564_arXiv.txt": {
        "abstract": "The unidentified infrared bands (UIBs) have been observed in virtually every dusty astrophysical environment investigated.  The UIB carrier must be abundant and ubiquitous. Strong evidence points to polycyclic aromatic hydrocarbons (PAHs) as likely candidates, but the identification is not complete.  Additional diagnostics are needed to further constrain the UIB carrier, such as probing excitation sources which range from UV-strong to UV-weak, in order to determine the ``band gap'' of the UIB carrier.  Observations and models suggest that the UIBs can be found in sources with weak UV fields.  To that end, we present new results of observing the 3.3 $\\mu$m spectral region in six stars embedded in reflection nebulae, and in six Vega-like stars.  These objects have effective temperatures ranging from  3500 to 12,000 K. Their environments include dust that should be relatively unprocessed (reflection nebulae) and dust that has most likely undergone significant processing (Vega-like) by the embedded illumination source. Together with data from the literature, we have a sample of 27 sightlines. Our analysis suggests that neither the strength of the UV field impinging on the dust nor the effective temperature of the star is the determining factor in whether the 3.3 $\\micron$ UIB emission is present in an object. We found three detections of the 3.3 \\micron~emission band, one in a Vega-type object, one in a Herbig Ae/Be object and one in a reflection nebula, and all with disks.  The role of disk geometry is likely to be important in revealing or obscuring the photo-dissociation regions from which the UIB emission arises. ",
        "introduction": "The unidentified infrared bands (UIBs), at 3.3, 6.2, 7.7, 8.6, 11.3, and 12.7 $\\mu$m, are the most prominent features in the near-infrared (NIR). They appear in spectra of the Galactic disk and high-latitude cirrus clouds (Matilla et al. 1996, Lemke et al. 1998) and other galaxies, including starburst galaxies (Roche 1989, Genzel and Cesarsky 2000), and in a wide range of astrophysical objects such as reflection nebulae (NGC 7023), planetary nebulae, dark clouds and photo-dissociation regions such as the Orion Nebula (Sellgren et al. 1990). Discovered in the 1970s (Gillett et al. 1973, Soifer et al. 1976 and Russell et al. 1977), today they have been observed to be present in most UV-rich astrophysical environments.  The proposed carriers of the UIBs include polycylic aromatic hydrocarbon (PAH) molecules (Leger \\& Puget 1984; Allamandola, Tielens, \\& Barker 1985) and grains containing PAHs (Duley \\& Williams 1981, Leger \\& Puget 1984).  Previous observations show UIBs occur in environments with moderate to high UV-excitation conditions, however, few UV-poor objects have been observed in order to determine the lower-limit for excitation of the UIB at 3.3 $\\mu$m.  UIB spectra can be split into two groups, those seen in interstellar dust (A-type) and those which are observed in circumstellar dust (B-type) (Geballe 1997; Tokunaga 1997). Class A-type UIBs, which include somewhat narrow, emission features at 3.3, 6.2, 7.7, 8.7, 11.3, and 12.7 $\\mu$m, are often observed in UV-rich sources.  Most UIB spectra appear to fall under Class A, having been observed in reflection nebulae, planetary nebulae, H II regions, starburst galaxies and even the diffuse interstellar medium (ISM).  Class B-type UIBs, which include narrow emission features at 3.3, 6.2, and 6.9 $\\mu$m and broad emission bumps at 3.4, 8, and 12 $\\mu$m (Geballe 1997; Tokunaga 1997), have been found in circumstellar dust illuminated by F- and G-type post-asymptotic giant branch (AGB) stars forming planetary nebulae (Uchida et al. 2000; Geballe et al. 1992).  In addition to spectra in these classes, another very small group of objects exhibits a unique pattern of features between 3.3 and 4 $\\mu$m. Van Kerckhoven et al. (2002), observed features at 3.3, 3.43 and 3.53 $\\mu$m in the star Elias 1 (A0-6), embedded in the reflection nebula IC 359 in the Taurus-Auriga complex (Elias 1978).  They attributed the 3.53 $\\micron$ feature to the presence of nanodiamonds in the circumstellar material.  This unusual grouping has only been observed in two other objects, HD 97048, an A0pshe Herbig Ae/Be, and HR 4049, a B9.5 Ib-II star evolving into a protoplanetary nebula.  The central question, which motivated our observations, was whether the 3.3 $\\mu$m UIB can be detected in weak UV field environments illuminated by stars with low T$_{eff}$.  This relates directly to the UV-absorption characteristics of the UIB carrier.  Models of PAH emission, building from extensive PAH laboratory work, suggest that a wide variety of PAHs may exist in dusty environments, with a typical size of 20-50 C atoms (Geballe et al 1989) (possibly larger), and absorption cross-sections dependent on size, molecular weight, and whether they are free-flying or incorporated into grains (Allamandola et al. 1989; Tielens 1993, Schutte et al. 1993, Cook and Saykally 1998; Bakes et al. 2001; Li \\& Draine 2001).  No firm UV-absorption cutoff has been established for astronomical PAHs, nor has observational evidence of PAH absorption in the UV been found (Clayton et al. 2003), but laboratory evidence suggests little absorption in the visible (Sellgren 2001).  One might expect, therefore, to observe a relationship between the presence of the 3.3 $\\mu$m UIB and environment (interstellar versus circumstellar), and specifically, with respect to the UV radiation field illuminating the dust.  It has been suggested that strong UV fields leave interstellar PAHs mostly neutral and positively charged, but weak UV fields are predicted to leave PAHs neutral or negatively charged (Bakes \\& Tielens 1994, 1998; Salama et al. 1996; Dartois \\& d'Hendecourt 1997).  Variations in the width of the 3.3 $\\mu$m feature have been observed among different objects such as in Orion (Sellgren, Tokunaga, \\& Nakada 1990), the reflection nebula NGC 7023, the planetary nebula (with a WC nucleus) IRAS 21282+5050 (Nagata et al. 1988) and toward the central star of the bipolar nebula HD 44179.  However, Sellgren et al. (1990) conclude that no dependence of width on UV field is evident.  Sellgren, Werner \\& Allamandola (1996) found the equivalent width of the 3.3 $\\micron$ UIB in reflection nebulae to be independent of $T_{eff}$.  Their coolest 3.3 $\\micron$ detections were for the nebula vdB 10 around HD 20041 (an A0 Ia-type star, $T_{eff}$ = 9400 K) and vdB 133 around HD 195593 (A+B) (an F5 Iab and a B7 II-binary, $T_{eff}$ = 6800 + 12,000 K).  Uchida et al. (2000) saw no spectroscopic differences among UIBs observed from 5-15 $\\mu$m for sources with T$_{eff}$ ranging from 3600-19,000 K.  The binary HD 195593 (A+B) illuminating vdB 133, is thought to be unique.  Uchida et al. (1998) note that vdB 133 is illuminated by the lowest ratio of UV to total flux of any nebulae observed to have the UIBs above 5 $\\mu$m.  This led Li \\& Draine (2002) to suggest that UV photons are not required to excite UIB emission. See Figure 1.  In this paper, we present new observations of vdB 133 and 11 other objects, chosen to cover both Class A and B-type sources with stars ranging from 3500-11,000 K in order to probe the dependence of the 3.3 $\\mu$m emission feature on UV radiation field.  The following section (2) describes the data reduction, and in section (3), we discuss the results and present our conclusions. ",
        "conclusions": "Li \\& Draine (2002) have suggested that UV photons are not required to excite UIB emission. The T$_{eff}$ of the exciting star will determine how much UV radiation is produced.  But the T$_{eff}$ does not take into account the dilution effect of the distance of the dust from the star.  A more useful parameter is the UV radiation field experienced by the dust at some distance from the star.  We estimated the UV radiation fields for the 27 objects in our sample by integrating Kurucz (1992) stellar atmospheric models from 91-240 nm, and then multiplying by a dilution factor of $(R_{*}/d)^{2}$, where $R_{*}$ is the stellar radius and d is the distance from the star at which the UV radiation field is being estimated.  A UV-cutoff of 240 nm was chosen in order to make a conservative estimate of photons available for the excitation of the 3.3 $\\mu$m UIB feature, as the UV-absorption spectrum for potential carriers such as PAHs is as yet unknown.  The UV radiation field $15^{''}$ NS of each star, corresponding to the slit length available on SpeX, is expressed as $G_{0}$. This is the local interstellar radiation field as determined by Habing (1968), in units of 1.6 x $10^{-3} erg s^{-1} cm^{-2}$ (Boulanger 1998, Dartois \\& d'Hendecourt 1997, Uchida et al. 2000).  The estimated values of $G_{0}$ for each object are listed in Table 2. The $G_{0}$ value of 2700 estimated for HD 195593 (A+B) at an offset of $15^{''}$, is consistent with the $G_{0}$ value from Uchida et al. (2000) of 75 for an offset of $100^{''}$ from the star.  Li \\& Draine (2002) estimate a $G_{0}$ value of $\\approx$ 1215 ``at the nebula'', using the linear combination of the Kurucz (1979) model atmospheres integrated for each star (A \\& B) between 91.2-1000 nm.  Figure 3 plots the estimated UV radiation field versus effective temperature of the exciting stars for our sample with detections of the 3.3 $\\micron$ feature marked. Clearly, no correlation exists between the strength of the UV radiation field ($G_{0}$) and the detection of the 3.3 $\\micron$ emission feature in the observed sample (ours and those of Meeus et al.).  The calculation of $G_{0}$ assumes that the dust is situated at a distance from the exciting star corresponding to 15 $\\arcsec$, projected on the plane of the sky.  Since matter around the stars is in many cases distributed out to 15 $\\arcsec$ (the available SpeX slit length), this represents the minimum value for $G_{0}$ impinging on the dust.  Table 2 includes the IRAS total surface brightness of the sky surrounding a target in a 0.5 degree beam at 100 $\\mu$m; a value above 30 MJy/Sr indicating dust with appreciable column density.  As can be seen from the table, all but five of the stars for which measurements were available have surface brightnesses larger than 30 MJy/Sr. The interesting exception being that for HD 34700, one of the three positive 3.3 $\\mu$m emission detections in our original sample, which has a measurement of 28 MJy/Sr.  Since all the stars with detections in our sample have disks, the emission is likely to be mostly from dust in a circumstellar disk much closer to the star than $15^{''}$ \\footnote{vdB 133 shows the 3.3 $\\micron$ emission in the nebulosity located beyond $15^{''}$ from the exciting star.}. This circumstellar disk dust will have similar dilution factors from star to star so T$_{eff}$ can be used to estimate the relative strength of the UV radiation field. But it can be seen from Figure 3 that T$_{eff}$ is no better correlated with the presence of the 3.3 $\\micron$ feature than with $G_{0}$. So neither T$_{eff}$ nor $G_{0}$ can be used to predict whether PAH emission will be seen along a particular sightline. One clue may come from the fact that 5 of the 6 stars that show PAH emission in the Meeus sample also have IR excesses indicating the presence of warm dust. So geometry may be important. Portions of the disk which are warm will probably correspond to photo-dissociation regions that are receiving photons from the star. So we may be viewing these disks more pole-on. Whereas, thick disks viewed edge-on will be less likely to show an IR excess or PAH emission because we are seeing only cold dust that is not being illuminated by any stellar photons. Flared disks tend to have UIBs whereas non-flared disks have weak to no UIB detections (Meeus et al. 2001). Perhaps the flaring of these disks increases the viewing angle where warm dust can be seen.  Since cool stars with small UV radiation fields such as HD 34700 show 3.3 $\\mu$m emission, the trend in Figure 3 supports the prediction of the astronomical PAH model of Li \\& Draine (2002), that the 3.3 $\\micron$ UIB should be observable in UV-weak objects.  But Figure 3 also shows that relatively hot stars with high UV radiation fields do not show any such emission.  So our results suggest that the strength of the UV radiation field impinging on the dust is not the only or even the most important factor in determining whether the 3.3 $\\micron$ UIB is present in a dusty astrophysical object. In addition to simple disk-viewing angle considerations, electron density variations from star to star may be an important factor in whether the 3.3 $\\mu$m emission feature is observable in an object.  However, a full spectrum of each star extending from 3.3 to 10 $\\mu$m is needed in order to solidly model the UIB features and estimate the electron density.  It is also likely that the ionization state of the PAHs in each object will affect the observability of the 3.3 $\\mu$m emission feature.  Laboratory studies (Allamandola et al. 1999) show the 3.3 $\\mu$m absorption feature to be very strong, relative to the features in the 6 to 10 $\\mu$m region, for a mixture of neutral PAHs, and to be much weaker, with respect to features in the 6 to 10 $\\mu$m, for the same mixture of PAHs in their positive state.  We conclude that the task of determining the role of the impinging radiation field on the observability and variations in the 3.3 $\\mu$m emission feature in our targets would be furthered by spectra from the same stars covering the 6-10 $\\mu$m region."
    },
    "0404/astro-ph0404202_arXiv.txt": {
        "abstract": "The recent observations of type Ia supernovae strongly support that the universe is accelerating now and decelerated in the recent past. This may be the evidence of the breakdown of the standard Friedmann equation. Instead of a linear function of the matter density, we consider a general function of the matter density to modify the Freidmann equation. We propose a new model which explains the recent acceleration and the past deceleration. Furthermore, the new model also gives a decelerated universe in the future. The new model gives $\\Omega_{\\rm m0}=0.46$ and $z_{\\rm T}=0.44$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0404/astro-ph0404034_arXiv.txt": {
        "abstract": "We present the stellar kinematics of $48$ representative elliptical and lenticular galaxies obtained with our custom-built integral-field spectrograph \\sauron\\ operating on the William Herschel Telescope. The data were homogeneously processed through a dedicated reduction and analysis pipeline. All resulting \\sauron\\ datacubes were spatially binned to a constant minimum signal-to-noise. We have measured the stellar kinematics with an optimized (penalized pixel-fitting) routine which fits the spectra in pixel space, via the use of optimal templates, and prevents the presence of emission lines to affect the measurements. We have thus generated maps of the mean stellar velocity $V$, the velocity dispersion $\\sigma$, and the Gauss--Hermite moments $h_3$ and $h_4$ of the line-of-sight velocity distributions. The maps extend to approximately one effective radius. Many objects display kinematic twists, kinematically decoupled components, central stellar disks, and other peculiarities, the nature of which will be discussed in future papers of this series. ",
        "introduction": "\\label{sec:intro} We are carrying out a study of the structure of 72 representative nearby early-type galaxies and spiral bulges based on measurements of the two-dimensional (2D) kinematics and line-strengths of stars and gas with \\sauron, a custom-built panoramic integral-field spectrograph for the William Herschel Telescope (WHT), La Palma. \\sauron\\ is based on the {\\texttt {TIGER}} microlens concept (\\citeauthor{betal95} \\citeyear{betal95}; see also \\citeauthor{becm00} \\citeyear{becm00}), and is described in detail in Paper I (\\citeauthor*{betal01b} \\citeyear{betal01b}). The objectives of the \\sauron\\ survey are summarized in Paper II (\\citeauthor*{zetal02} \\citeyear{zetal02}), which also contains the definition and global properties of the sample. Here we present maps of the stellar kinematics for the 48 elliptical (E) and lenticular (S0) galaxies in the survey.  The morphology and kinematics of the ionized gas, the line-strength distributions, similar measurements for the spiral galaxies in the sample, and higher spatial resolution integral-field observations obtained with \\oasis\\ on the Canada-France-Hawaii Telescope (CFHT) will be presented separately. The data and maps presented in this paper will be made available via the \\sauron\\ WEB page http://www.strw.leidenuniv.nl/sauron/.  In Section~\\ref{sec:observations} we summarize our observational campaign, as well as the principles of the data reduction and the measurement of the line-of-sight velocity distributions. We present the kinematic maps in Section~\\ref{sec:maps}, and our conclusions in Section~\\ref{sec:conclusions}. \\vfill ",
        "conclusions": "\\label{sec:conclusions}  The maps presented in this paper are the result of the first survey of the stellar kinematics of nearby early-type galaxies with an integral-field spectrograph. This unique dataset demonstrates that early-type galaxies display significant and varied structure in their kinematic properties, reinforcing the preliminary evidence shown in Paper II, and the hints provided by earlier long-slit spectroscopy. The present \\sauron\\ maps demonstrate that 2D coverage is a prerequisite to properly understand the dynamics of nearby galaxies.  The maps shown here constitute only part of the information contained in the \\sauron\\ spectra. The morphology and kinematics of the emission-line gas, as well as maps of the various line-strengths derived from the same observations will be presented in future papers."
    },
    "0404/astro-ph0404344_arXiv.txt": {
        "abstract": "{ We investigate the sample of 1175 new nonmagnetic DA white dwarfs with the effective temperatures $\\te \\ge 12000$ K, which were extracted from the Data Release 1 of the Sloan Digital Sky Survey. We determined masses, radii, and bolometric luminosities of stars in the sample. The above parameters were derived from the effective temperatures $\\te$ and surface gravities $\\log g$ published in the DR1, and the new theoretical $M-R$ relations for carbon -- core and oxygen -- core white dwarfs. Mass distribution of white dwarfs in this sample exhibits the peak at $M=0.562 M_\\odot$ (carbon core stars), and the tail towards higher masses. Both the shape of the mass distribution function and the empirical mass -- radius relation are practically identical for white dwarfs with either pure carbon or pure oxygen cores. ",
        "introduction": "The Sloan Digital Sky Survey is the very extensive and ambitious ground--based research project, aimed at the determination of positions, brightnesses, and acquiring of the optical spectra of millions of various celestial bodies. The SDSS was described in detail on their web page, {\\rm http://www.sdss.org/}, and in York et al. (2000). While this project was essentially designed to study galaxies, quasars, and large-scale structures in the Universe, its secondary targets are stars of various spectral types. We take advantage of the SDSS data on white dwarf stars of various classes, which recently were made publicly available as the subset of the Data Release 1.  Comprehensive analysis of the photometric and spectral data on white dwarfs, which are available in the DR1, was presented in Kleinman et al. (2004). They have found 2935 white dwarfs of various types and related stars, and only few of them were spectroscopically identified previously. The number of well identified white dwarfs in their sample approaches the number of white dwarf stars known previously, see McCook \\& Sion (1999), and the newest version of their catalogue available at {\\rm http://www.astronomy.villanova.edu/WDCatalog/index.html/.  Kleinman et al. (2004) performed fitting of the observed spectra of 1833 nonmagnetic DA white dwarfs to the widely used theoretical grid of pure hydrogen model atmospheres by Detlev Koester (see Finley, Koester \\& Basri 1997). They have obtained effective temperatures in the range $7220\\le \\te \\le 93855$ K. Surface gravities of stars in the sample range } from low values corresponding to subdwarfs to the most massive white dwarfs of the highest $\\log g = 9.0$.  The referee pointed out, that the SDSS DR1 values of $\\log g$ determined for stars with $\\te < 12000$ K are systematically larger than for hotter stars, and are larger than expected for white dwarfs, cf. Fig. 9 in Kleinman et al. (2004). The explanation of this effect was not found, however, they favor the possibility that cooler model atmospheres exhibit systematic errors in the parametrisation of convection, or in the fitting procedure. As the result, they do not claim that the high values of $\\log g$ are real. To avoid distortion of the high-mass tail presented in Fig.~1 we analyse a subset of the original SDSS DR1 sample which includes 1175 DA white dwarfs with $\\te \\ge 12000$ K.  The Koester's grid does not include spectra of models with high gravities, $\\log g > 9.0$. Therefore, available determinations of $\\log g$ also do not exceed 9.0, and this is an artificial bias which influences the determination of $\\log g$ for over 20 most massive DA white dwarfs in the whole SDSS DR1 sample.  One can note, that white dwarfs of the highest gravity exceeding 9.0 appear among previously known nonmagnetic DA stars. Compilation of the existing mass determinations by Nale\\.zyty \\& Madej (2004) lists 6 such white dwarfs, with the maximum $\\log g = 9.12$. Also the most recent paper by Dahn et al. (2004) reports the measurement of $\\log g = 9.46$ in the apparently most massive star LHS 4033. Inferred mass of this nonmagnetic DA star is $M \\approx 1.33 \\, M_\\odot$, and the corresponding small radius $R=0.00368 \\, R_\\odot$, place this star close to the limiting Chandrasekhar mass. ",
        "conclusions": "\\noindent {\\bf 1.} We have determined the distribution of DA white dwarf masses for the magnitude limited sample of 1175 stars in the SDSS DR1 table. Our mass distribution exhibits the peak at $M=0.562 \\, M_\\odot$, and slight asymmetry towards both higher and lower masses.  The peak mass of this paper is in excellent agreement with earlier studies, which used much less numerous samples and various methods of white dwarf mass determination. We present here a brief listing of previously determined peak masses of DA white dwarfs. Unfortunately, the early paper by Shipman (1979) does not provide data corresponding to data of other papers, which are collected in Table 1.  There exists a striking consistency of all peak mass determinations, in spite of relatively nonnumerous samples in the above listing. Our research adds the peak mass $M_{\\rm peak} = 0.562 \\, M_\\odot$ to the list (carbon core stars), which perfectly confirms the value of $M_{\\rm peak}$ by Bergeron, Saffer \\& Liebert (1992). Moreover, we claim that have reliably determined the shape of mass distribution of DA white dwarfs, except for the single bin in Fig. 1, corresponding to the most massive stars.  \\vskip2mm \\noindent {\\bf 2.} We note, that spectroscopic masses $M$ and radii $R$ in the present and earlier papers can still exhibit systematic errors due to uncertainties in the present theories of plasma physics. Theoretical profiles of the Balmer lines in spectra of DA white dwarfs are very strongly broadened, and the dominant contributor is the pressure broadening of line wings. The correct description of the pressure broadening in dense plasma of high gravity white dwarf atmospheres is a very difficult problem.  Recently Wujec et al. (2002) proposed the new method of Stark broadening calculations based on simulation techniques. Their theory is able to compute the broadening, shift, and asymmetry parameters of Stark broadened hydrogen line opacity profile. Previous Stark broadening calculations were not able to include all the three plasma effects simultaneously. However, no grids of theoretical Balmer line profiles in white dwarf spectra by Wujec et al. (2002) are currently available.  \\vskip2mm \\noindent {\\bf 3.} Catalogues of masses, radii, resulting bolometric luminosities for both C -- core and O -- core white dwarfs, and all the input parameters are available from authors on request. Bolometric luminosities and their errors are given by the following expressions \\begin{equation} L_{\\rm bol} = 4\\pi R^2 \\, \\sigma_{\\rm rad} \\te ^4 \\, , \\end{equation} \\begin{equation} \\delta L_{\\rm bol} = 8\\pi R \\,\\sigma_{\\rm rad}\\te^3 \\, \\left(\\te \\,\\delta R + 2 R \\, \\delta\\te \\right) \\, . \\end{equation}   \\begin{table} \\caption{Peak masses of the published DA white dwarf mass distributions, in  $M_{\\rm peak} / M_\\odot $.  } \\renewcommand{\\arraystretch}{1.1} \\begin{tabular}{rll} -- & $0.58 $ & Koester, Schulz \\& Weidemann (1979), average \\\\ &         & mass of 122 DA white dwarfs                  \\\\ -- & $0.603$ & Weidemann \\& Koester (1984), 70 DA WD stars     \\\\ -- & $0.571$ & McMahan (1989), 53 DA white dwarfs              \\\\ -- & $0.60 $ & Weidemann (1990)                                \\\\ -- & $0.562$ & Bergeron, Saffer \\& Liebert (1992), 129 DA white\\\\ &         & dwarfs                                          \\\\ -- & $0.56 $ & Liebert \\& Bergeron (1995), 200 white dwarfs from \\\\ &         & the Palomar Green survey (Green et al. 1986)   \\\\ -- & $0.570$ & Finley, Koester \\& Basri (1997), 174 DA white   \\\\ &         & dwarfs, some with cool companions               \\\\ -- & $0.55 $ & Marsh et al. (1997a,b), 89 stars of the ROSAT  \\\\ &         & All-Sky X-ray and EUV Surveys                  \\\\ -- & $0.56 $ & Vennes et al. (1997), 110 EUV selected DA WDs   \\\\ -- & $0.57 $ & Vennes (1999), 141 EUV/X-ray selected WDs       \\\\ \\end{tabular} \\end{table}"
    },
    "0404/astro-ph0404458_arXiv.txt": {
        "abstract": "We present initial results of an ongoing search for interferometric calibrators at submillimeter (sub-mm) wavelengths with the Submillimeter Array (SMA). Powerful radio galaxies are commonly used as calibrators at centimeter and millimeter wavelengths, but many are not strong enough to serve as calibrators at sub-mm wavelengths because of their rapidly declining flux densities toward shorter wavelengths. The inability to find a calibrator close to the target source may limit or even prevent us from imaging many interesting sources at sub-mm wavelengths. Here, we investigate whether high-mass protostellar objects and ultracompact \\htwo regions can serve as useful calibrators for the SMA. The dust emission associated with these objects makes them among the brightest sub-mm sources in the sky. Our observations at 0.85~mm (345 GHz) with an angular resolution of $\\sim$3$\\arcsec$ reveal that although a large fraction of the dust emission originates from an extended ``halo'' component, a compact unresolved component often remains that when sufficiently strong may serve as a useful calibrator. These observations also provide a first glimpse at the small-scale distribution of dust around ultracompact \\htwo regions and high-mass protostellar objects at sub-mm wavelengths. We discuss the origin of the core-halo structure seen in many sources, and conclude with suggestions for future searches for calibrators with the SMA. ",
        "introduction": "Calibrators used for determining the amplitude and phase (hereafter referred to as gain) of an interferometer are usually compact (i.e., spatially unresolved) and bright. The difference between the measured and expected gain of the calibrator, introduced by the Earth's atmosphere and telescope electronics, are used to correct the gain measured towards the target source. Because transmission (affecting the amplitude) and refraction (affecting the phase) through the Earth's atmosphere can differ significantly along different lines of sight, the calibrator should lie close in the sky to the target source. At centimeter and millimeter wavelengths, powerful radio galaxies are sufficiently compact, bright and numerous to serve as calibrators in virtually any region of the sky. At submillimeter (sub-mm) wavelengths, however, most of these radio galaxies are too weak to serve as calibrators for the Submillimeter Array\\footnote{% The Submillimeter Array is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy and Astrophysics, and is funded by the Smithsonian Institution and the Academia Sinica.} (SMA). This is because their synchrotron emission usually decreases rapidly in intensity towards shorter wavelengths, and furthermore the instrumental noise is higher because of higher system temperatures at shorter wavelengths. We estimate that only about eighty radio galaxies are bright enough to serve as calibrators for the SMA at 0.85 mm (345 GHz), so that the mean angular separation between calibrators in the sky is $\\sim$20$^\\circ$. This may limit our ability to image many interesting sources, especially at high angular resolutions where accurate gain corrections (and hence close proximity of a calibrator) become increasingly important. At 0.45 mm (690 GHz) the situation is drastically worse, with only about a dozen radio galaxies useable as calibrators for the SMA; none are located at R.A. 5$^h$$-$19$^h$ and Dec. $>$ 5$^\\circ$. Even with the addition of planets and their moons to the list of available calibrators, we still may not be able to image sources in large swaths of the sky.  By contrast with synchrotron emission, thermal emission from dust has a rapidly increasing intensity towards shorter wavelengths, so that dust sources are good candidates for calibrators at sub-mm wavelengths. A particularly promising class of objects are young high-mass stars, which can often be found embedded in their natal molecular cloud cores even when they have produced sufficient UV photons to excite an ultracompact (UC) \\htwo region. The emission from their dust cocoons makes them among the brightest sources in the Galaxy at sub-mm wavelengths; in single-dish observations, high-mass star-forming regions can exhibit continuum emission of up to a few hundred Jy at 0.35 mm \\citep{san94,hun00}. Their dust distribution on the relative small spatial scales that interferometers probe, however, is not known. Recently, high angular resolution ($\\sim$0\\farcs5) mid-IR observations of the UC \\htwo region G29.96$-$0.02 showed a good spatial coincidence between warm dust grains and ionized gas imaged at 1.3 cm \\citep{deb02}. If the cooler dust probed at sub-mm wavelengths has a similar distribution, UC \\htwo regions may make suitable calibrators at sub-mm wavelengths. Their location in the Galactic plane makes them an especially important addition to the available calibrators for studies of star-forming regions and other Galactic objects. In this Letter, we present the first results at 0.85 mm to determine whether high-mass star forming regions can serve as calibrators for the SMA. Because of the relative paucity of available calibrators at 0.45 mm, this also is a necessary first step towards determining whether young high-mass stars can potentially serve as calibrators at shorter wavelengths. ",
        "conclusions": "\\subsection{Structure of the 0.85 mm Continuum Emission} We detected continuum emission towards all target sources. The two sources G10.841$-$2.592 and Onsala~1 show distinct multiple components, as shown in the left panel of Figure 1 for Onsala 1. These sources are for the time being ruled out as suitable calibrators at 0.85 mm. For the remaining six sources, five show what appears to be a single compact component, whereas in the case of G15.04$-$0.68 extended emission is also detected. As diffuse/extended structures can be difficult to discern in the image plane, we also examined the amplitude of the measured visibilities as a function of projected baseline length (i.e., {\\it uv} distance). For an extended source, the visibility amplitude generally decreases with increasing {\\it uv} distance, whereas a point source shows a constant visibility amplitude over all {\\it uv} range. As an example, the right panel of Figure~1 shows a map of the 0.85 mm continuum towards G45.07+0.13. This image reveals a compact ``core'' component that is at best only marginally resolved. The visibility amplitude plot of Figure~\\ref{g45-uv}, however, clearly reveals an extended ``halo'' structure. Indeed, all the five sources that appeared to be compact in our maps were found to exhibit extended emission.  To estimate the flux densities and sizes of the different components in each target source, we first fitted the measured visibilities with a one-dimensional Gaussian source structure (i.e., halo component). When this fit proved unsatisfactory, we added a point source (i.e., core component). As shown in Figure \\ref{g45-uv}, the visibility amplitudes of G45.07+0.13 can be fit with a core-halo structure with flux densities of 0.90 Jy and 2.71 Jy for the two respective components, and a half-power width for the halo of $\\sim$3\\farcs8. The results of the fits for all six sources are listed in Table~2. Two are satisfactorily fit by a single Gaussian component, whereas the four sources G15.04$-$0.68, G45.07+0.13, IRAS 18089$-$1732, and IRAS 19217+1651 require a core component with flux density of at least 0.35 Jy in addition to a halo component.  \\subsection{Suitability as Calibrators} The first two sub-mm bands at which the SMA can operate are 0.85 and 0.45 mm. With a system temperature of 500/2000 K and aperture efficiency of 0.70/0.60 at 0.85/0.45 mm, and a bandwidth of 2 GHz, the rms noise level for a single baseline in a 5-min integration is about 0.07/0.35 Jy. We require the calibrators to be strong enough that the phase variations due to random noise fluctuations in a given calibration scan are smaller than those produced by the atmosphere between scans. For typical weather at Mauna Kea, over a 15 minute interval the atmospheric phase variations are 6 to 30 degrees for 10~m to 100~m baselines at 0.85 mm, and a factor of 2 higher at 0.45 mm \\citep{mas94}. Such phase variations are comparable with those seen in our observations of Uranus/Neptune at 0.85 mm. For the SMA, the (antenna-based) rms phase noise is 28/X degrees \\citep{bri89}, where X is the signal-to-noise ratio for a single baseline. To achieve an rms phase precision of $\\sim$6 degrees over a 5-min calibrator scan, we require X to be 5, corresponding to a flux density of 0.35/1.8 Jy at 0.85/0.45 mm. All four of the above-mentioned sources with point-like components therefore qualify as calibrators at 0.85 mm, provided that short baselines sensitive to the halo component are excluded and enough baselines remain to derive (antenna-based) gain solutions. Observations in superior weather conditions requiring higher phase precision will naturally require stronger calibrators and/or longer integration times.  For the above-mentioned sources to be suitable also as calibrators at 0.45 mm, we require them to have a frequency spectrum \\emph{S}$_\\nu$ $\\propto$ $\\nu$$^{>2.5}$. If the 0.85 mm emission is produced entirely (or predominately) by dust, then we require a dust opacity index $\\beta$ $\\gtrsim$ 0.5; such a dust opacity index is commonly found in star-forming regions \\citep[e.g.,][]{mol00}. We therefore anticipate that G15.04$-$0.68, IRAS~18089$-$1732, and IRAS 19217+1651, whose cm-wave continuum emission is relatively weak ($\\lesssim$30 mJy) and hence their 0.85 mm emission dominated by dust, are suitable as calibrators also at 0.45 mm. This, however, may not be the case for G45.07+0.13. As shown in Figure \\ref{g45-sed}, the flux densities of G45.07+0.13 at 2$-$1.3 cm (15$-$22 GHz) measured with the VLA range from 0.6 to 0.9 Jy, comparable to the measured flux density of the core component at 0.85 mm. In this case, a large fraction of the core emission may be produced by free-free emission.   \\subsection{Dust Distribution of UC \\htwo Regions and HMPOs}  Three of the sources in Table~2 have been observed in the continuum at 0.85~mm with the JCMT \\citep{san94,jen95,wal03}. Although the primary beam of the JCMT is a factor of 2.5 smaller than that of the SMA, nevertheless the single-dish flux density is a factor of 1$-$3 larger than that inferred for our core and halo components combined. Because our measurements do not provide stringent constraints on the radial intensity profile of the halo component (modelled here as a symmetric Gaussian), we may have underestimated its total flux density. For example, if the halo component has a power-law radial density and/or temperature profile instead of a Gaussian profile, its visibility amplitude will decrease more rapidly with increasing {\\it uv} distance and hence its total flux density higher. On the other hand, there may be a distinct, very extended ($\\gtrsim$10$\\arcsec$) component that our interferometric observations do not detect. The halo components that we observe have sizes ranging from 0.02 to 0.2 pc, which is comparable to or smaller than the molecular envelopes traced in CS (\\emph{J}=5$-$4) around young high-mass stars of $\\sim$0.3 pc \\citep{shi03}. Thus, the halo component may trace the inner regions of this dense molecular envelope.  The presence of a core component implies a compact region that is denser and/or warmer than its immediate surroundings. \\citet{van00} also find compact ($<$2\\arcsec) dust components in several HMPOs and UC \\htwo regions at 3~mm and 1.3~mm, and suggest that these components may originate from circumstellar disks. Circumstellar disks are commonly seen around low-mass protostars, but have been inferred only around a few massive young stellar objects. \\citet{yor96} postulate that the photoevaporation of such disks is required to inject material to sustain UC \\htwo regions. For a dust temperature of 50~K (and $\\beta$ = 1), the inferred masses (for an interstellar gas to dust ratio of $\\sim$100) of the compact components are 3$-$50 \\emph{M}$_\\sun$. By comparison, \\citet{she01} reported a (dust) disk around G192.16$-$3.82 with a mass of 3 \\emph{M}$_\\sun$, while \\citet{zha98} reported a (NH$_3$) disk around IRAS 20126+4104 with a minimum mass of $\\sim$5 \\emph{M}$_\\sun$. Of course, a part or all of the core emission could originate instead from a centrally collapsed region that surrounds the central star. For UC \\htwo regions, the compact component could in addition be produced by warm dust either cospatial with or in a relatively thin layer around the ionized gas, as seen in the mid-IR \\citep{kra01,kra03}.   \\subsection{Future Directions} In our initial search for calibrators among eight young high-mass stars, we identified four potential candidates. Although promising, we still need to demonstrate that calibration off the point-source component in these sources can be successful (this test was not possible here given the limited number of baselines that sampled the point-source component). If not, we plan to investigate whether detailed maps of these sources can be used as an input for gain calibration (analogous to the case in VLBI experiments). Bright and compact \\htwo regions whose free-free emission is by itself sufficiently strong to serve as calibrators at sub-mm wavelengths (e.g., G45.07+0.13) should also be explored.  In its sparsest configuration, the SMA can attain an angular resolution of $\\sim$0\\farcs3 at 0.85 mm and $\\sim$0\\farcs15 at 0.45 mm. Even if sufficiently compact dust sources associated with high-mass star-forming regions can be found, it is unlikely that they will be bright enough to serve as useful calibrators. To illustrate, an opaque dust source with a temperature of $\\sim$200 K (the highest so far measured in high-mass star-forming regions; \\citet{ces94}) and a size of 0\\farcs3/0\\farcs15 at 0.85/0.45 mm has a flux density of $\\sim$1.8 Jy. This is at the lower limit required for calibration at 0.45 mm, and for more reasonable dust temperatures and opacities, may not exceed the required flux density to allow calibration even at 0.85 mm. Thus, for observations at the highest angular resolutions with the SMA where no known calibrators close to the target source exist, we may have to resort to 2-D mapping of relatively compact ($\\lesssim$ 1$\\arcsec$) dust sources (if such sources exist). Finally, the SMA was designed to observe in two different wavelength bands simultaneously to allow observations of a calibrator at a longer wavelength to be used to correct the gain of the telescope at a shorter wavelength, but this capability remains to be implemented and tested."
    },
    "0404/astro-ph0404172_arXiv.txt": {
        "abstract": "{We report the detection of CO emission from the optically luminous, radio-quiet quasar J140955.5+562827 (hereafter \\sourcename), at a redshift $z_{\\rm CO} =2.583$.  We also present VLA continuum maps and VLBA high spatial resolution observations at 1.4~GHz. Both the CO(3\\to2) and CO(7\\to6) emission lines are detected using the IRAM Plateau de Bure interferometer. The 3\\to2/7\\to6 line luminosity ratio is about 1/3, indicating the presence of warm and dense molecular gas with an estimated mass of $6 \\times 10^{10}$~\\Msun.  The infrared-to-CO luminosity ratio $L_{\\rm FIR}/L^\\prime_{\\rm CO(1\\rightarrow0)} \\approx 500$~\\Lsun~(K~km~s$^{-1}$~pc$^2$)$^{-1}$, comparable to values found for other high-$z$ sources where CO line emission is seen. \\sourcename~ is detected using the VLA with a 1.4~GHz rest-frame luminosity density of $4.0\\times10^{25} \\, \\rm W \\, Hz^{-1}$.  The rest-frame radio to far-infrared ratio, $q$, has a value of 2.0 which is similar to the values found in star forming galaxies.  At the 30~mas resolution of the VLBA, \\sourcename~ is not detected with a 4$\\sigma$ upper limit to the surface brightness of 0.29~mJy~beam$^{-1}$. This implies a limit to the intrinsic brightness temperature of $\\rm 2 \\times 10^5 \\, K$ at 8~GHz, typical for nuclear starbursts and more than two orders of magnitude fainter than typical radio-loud active galactic nuclei. Both the properties of the CO line emission and the radio emission from \\sourcename~ are therefore consistent with those expected for a star forming galaxy. In \\sourcename~ young massive stars are the dominant source of dust heating, accounting for most of the infrared luminosity. The massive reservoir of molecular gas can sustain the star formation rate of a few $\\rm 1000 \\, M_\\odot \\, yr^{-1}$ implied by the far-infrared luminosity for about 10 million years. ",
        "introduction": "\\label{sec:Introduction}  In the recent decade, millimeter and submillimeter deep blank field surveys \\citep[e.g.,][]{Ivison2000,Blain2002} and pointed observations \\citep[e.g.,][]{Omont2003,Priddey2003,Bertoldi2003} have provided a view of the dust content of a few hundred galaxies and quasars in the redshift range $1 < z < 6.4$.  In these objects, a significant fraction of the energy generated by star formation is processed by dust and re-emitted at far-infrared wavelengths. The average space density of ultraluminous infrared galaxies and quasars at high-$z$ is found to be thousandfold greater than in the local universe. Observations at millimeter and submillimeter wavelengths provide a direct way to trace the bulk of the star formation in the early universe.  Pointed submillimeter and millimeter observations of optically selected quasars or radio galaxies have the advantage over deep field surveys in providing unique source identification and redshifts.  In particular, targeted observations of optically luminous quasars have been successful in revealing massive bursts of star formation in their host galaxies. Although optically luminous quasars, which trace the most massive collapsed structures to have formed in the early universe, are rare objects which bear little relation to the objects that make up most of the far-infrared and submillimeter background, they provide powerful probes to study the relation between star formation, massive gas reservoirs, and the growth of super-massive black holes in the dark ages of the Universe \\citep{Fan2003,Bertoldi2003b,Walter2003}.  The study of the redshift range $2 < z <3$ is important in studying the star formation history in the universe since it traces the peak of the space density of the quasar population \\citep{Shaver1996}, and corresponds to the median redshift of submillimeter galaxies \\citep{Chapman2003}.  Recent pointed observations of the 1.2 and 0.85~mm thermal dust continuum emission of $z\\approx2$ optically luminous, radio-quiet quasars reveal that 1/3 of these quasars are also luminous in the infrared with far-infrared luminosities of $L_{\\rm FIR} \\rm \\sim 10^{13}$~\\Lsun~ and estimated dust masses of typically $\\rm \\sim 10^8 \\, M_\\odot$ \\citep{Omont2003,Priddey2003}.  The optically very bright ($M_B=-28.4$), radio-quiet quasar \\sourcename~ is by far the strongest mm source in the \\citet{Omont2003} survey. With an estimated far-infrared luminosity of $4 \\times 10^{13}$~\\Lsun, \\sourcename~ ranges amongst the most luminous infrared high-$z$ sources found to date.  Under the assumption that the dust is predominantly heated by massive stars, the inferred star formation rate is several $\\rm 1000 \\, M_\\odot \\, yr^{-1}$ \\citep{Omont2003}.  However, the fraction of infrared luminosity due to dust heated by star formation and dust directly heated by the Active Galactic Nuclei (AGN) in \\sourcename~ is still an open issue, as is the case for almost all of the high-$z$ quasars studied to date.  Both the large masses of dust and the high star formation rate indicate that \\sourcename~ should have a copious reservoir of molecular gas to sustain the formation of a few $1000 \\, M_\\odot$ per year over its dynamical time.   In high-$z$ sources, the search for molecular gas is best done by using the redshifted rotational lines of CO. The molecular gas in star-forming galaxies is warm and dense enough to excite the higher CO rotational levels ($\\rm J>3$) which are shifted into the millimeter windows for redshifts greater than 2.  However, a good determination of the redshift of the target object, either from optical or near-infrared spectroscopy, is a prerequisite for a detection because the bandwidths of the current heterodyne receivers are still relatively narrow. For example, at the Plateau de Bure Interferometer, the available instantaneous bandwitdh is $\\Delta \\nu = 580$~MHz which corresponds at 96~GHz to a velocity range of 1810~\\kms or $\\Delta z \\sim 0.02$ at $z=2.6$.  Despite this current limitation, over the last few years CO line emission has been detected in 24 far-infrared luminous high-$z$ quasars, radio and (sub)mm galaxies \\citep[][ and references therein]{Carilli2004}.  These observations reveal large reservoirs of molecular gas ($\\rm 10^{11} \\, M_\\odot$), a prerequisite for efficient star formation, and prompted also the first studies on the properties of the dense interstellar gas in starburst galaxies in the early universe \\citep[e.g.][]{Barvainis1997,Solomon2003,Bertoldi2003b}.  In addition, these observations allow us to estimate, unhindered by extinction, the dynamical masses in these systems, which are the key constraining hierarchical models \\citep[see, e.g.][]{Neri2003}.  In this paper, we present a study of the CO emission in \\sourcename~ made with the Plateau de Bure interferometer together with observations of the radio emission carried out with the Very Large Array and the Very Large Baseline Array. The data provide strong evidence that the origin of the infrared luminosity is dominated by the starburst activity of the host galaxy of the QSO \\sourcename. Throughout this paper, we adopt the concordance $\\Lambda$-cosmology with $H_0=71\\rm~km~s^{-1}~Mpc^{-1}$, $\\Omega_\\Lambda=0.73$ and $\\Omega_m=0.27$ -- \\citep{Spergel2003}. At $z=2.58$, the luminosity distance $D_L$ is 21.5~Gpc and an angular scale of $1^{\\prime\\prime}$ corresponds to 8.1~kpc. ",
        "conclusions": "\\label{sec:conclusions}  The millimetre and radio observations described in this paper provide strong evidence that the far-infrared luminosity of \\sourcename~ is related to massive star formation.  First, the huge reservoir of dense and warm molecular gas ($ 6 \\times 10^{10} \\, M_\\odot$) can sustain the star formation rate of a few thousand solar masses per year implied by the far-infrared luminosity for about 10 million years. Second, the radio continuum emission is found to be consistent with the radio/far-IR correlation of star forming galaxies. The limits on the size of the CO and radio continuum emitting region suggest that in \\sourcename~ star formation is taking place in an extended region, which could be a torus or disk-like structure not larger than 5~kpc. Observations of the CO emission at higher spatial resolution would be useful to further constrain the size of the star-forming region in \\sourcename~ and thereby its dynamical mass.  The submillimeter and radio properties of the $z=2.58$ quasar \\sourcename, which are similar to those seen in other high-$z$ optically luminous quasars, are consistent with a coeval starburst and AGN.  However, the relative importance of the warm dust heated by the star formation and the dust directly heated by the AGN is still poorly constrained.  Observations at higher frequencies are needed to analyze the rest-frame near- to mid-infrared spectral energy distribution of distant quasars like \\sourcename.  Such observations will allow us to derive the total infrared luminosity of high-$z$ quasars, to estimate the contribution of the hot dust heated by the AGN, to better constrain the star formation history in their associated host galaxies and, thereby, to investigate further the coeval growth of stars and massive black holes and its evolution over cosmic time."
    },
    "0404/astro-ph0404581_arXiv.txt": {
        "abstract": "The main aim of microlensing experiments is to evaluate the mean mass of massive compact halo objects (MACHOs) and the mass fraction of the Galactic halo made by this type of dark matter. Statistical analysis shows that by considering a Dirac-Delta mass function (MF) for the MACHOs, their mean mass is about that of a white dwarf star. This result is, however, in discrepancy with other observations such as those of non-observed expected white dwarfs in the Galactic halo which give rise to metal abundance, polluting the interstellar medium by their evolution. Here we use the hypothesis of the spatially varying MF of MACHOs, proposed by Kerins and Evans to interpret microlensing events. In this model, massive lenses with a lower population contribute to the microlensing events more frequently than do dominant brown dwarfs. This effect causes the mean mass of the observed lenses to be larger than the mean mass of all the lenses. A likelihood analysis is performed to find the best parameters of the spatially varying MF that are compatible with the duration distribution of Large Magellanic Cloud microlensing candidates of the MACHO experiment. ",
        "introduction": "The rotation curves of spiral galaxies (including Milky Way) show that this type of galaxies have dark halo structure \\cite{bro01}. One of the candidates for the dark matter in the Galactic halo may be massive compact halo objects (MACHOs). Paczy\\'nski (1986) proposed a gravitational microlensing technique as an indirect way of detecting MACHOs. Since his proposal, many groups began monitoring millions of stars of the Milky Way in the directions of the spiral arms, the Galactic bulge and the Large and Small Magellanic Clouds (LMC \\& SMC) and detected hundreds of microlensing candidates \\cite{ans04,der01,sum03,afo03}. Looking in the direction of the LMC and SMC (which are the most important for estimating Galactic halo MAHCOs), Exp\\'erience de Recherche d'Objets Sombres (EROS)\\footnote{http://eros.in2p3.fr/} and MACHO\\footnote{http://wwwmacho.mcmaster.ca/} observers found only a dozen of microlensing candidates \\cite{las00,alc00}. The interpretation of LMC and SMC events is based on the statistical analysis of the distribution of the duration of the events. The result of this analysis attribute a mean mass to MACHOs and their mass contribution in the Galactic halo. With the standard halo model, the mean mass of lenses is evaluated to be about half of that of the solar mass with a 20 per cent contribution in the Galactic halo mass.\\\\ The results obtained by the analysis of LMC microlensing events, however, do not agree with other observations \\cite{gat01}. Studying the kinematics of white dwarfs that have been discovered \\cite{opp01} has shown that halo white dwarfs corresponds to $1-2$ per cent of the halo mass. Recent re-analysis of the same data \\cite{spa04,tor02} shows that this fraction is an order of magnitude smaller than the value derived in Oppenheimer et al. (2001). On the other hand, if there were as many white dwarfs in the halo as suggested by the microlensing experiments, they would increase the abundance of heavy metals via the evolution of white dwarfs and Type I Supernova explosions \\cite{can97}. The other problem is that for the mass of the MACHOs to be in the range proposed by microlensing observations, the initial mass function (IMF) of MACHO progenitors of the Galactic halo should be different from those of the disc \\cite{ada96,cha96}, otherwise we should observe at the tail of mass function (MF) a large number of luminous stars and heavy star explosions in the Galactic halos.\\\\ In this study we use the hypothesis of a spatially varying MF (instead of uniform Dirac-Delta MF for the halo MACHOs) to interpret the LMC microlensing candidates (Kerins \\& Evans 1998, hereafter KE). The physical motivation for the hypothesis of spatially varying MFs of MACHOs comes from baryonic cluster formation theories \\cite{ash90,car94,de95}. These models predict the spatial variation of MF in the galactic halo in such a way that the the inner halo comprises partly visible stars, in association with the globular cluster population, while the outer halo comprises mostly low-mass stars and brown dwarfs.\\\\ We extend the work of KE by (i) using spatially varying MF model in the power-law halo model \\cite{alc96}, including the contribution disc, spheroid and LMC for comparison with the latest (5 yr ) LMC microlensing data (Alcock et al. 2000); (ii) using a statistical approach applied by Green \\& Jedamzik (2002) and Rahvar (2004) to compare the distribution of the duration of the observed events with the galactic models; and (iii) performing a likelihood analysis to find the best parameters of the inhomogeneous MF model.\\\\ The advantage of using spatially varying MF models is that the active mean mass of lenses as the mean mass of observed lenses is always larger than mean mass of overall lenses. This effect is shown by a Monte-Carlo simulation, and taking it into account may resolve the problems with interpreting microlensing data.\\\\ This paper is organized as follows. In Section 2 we give a brief account of the hypothesis of spatially varying MFs and the galactic models used in our analysis. In Section 3 we perform a numerical simulation to generate the expected distribution of events, taking into account the observational efficiency of the MACHO experiment. In Section 4 we compare the theoretical distribution of the duration of the events with the observation. We also perform a likelihood analysis to find the best parameters of the MF for compatibility with the observed data. The results are discussed in Section 5. ",
        "conclusions": "In this work we extended the hypothesis of a spatially varying MF proposed by KE as a possible solution resolving discrepancies between microlensing results and other observations. The main point of this is to investigate the contradiction where microlensing experiments predict large numbers of white dwarfs which have not been observed.\\\\ The advantage of using a spatially varying MF is that we can modify our interpretation of microlensing data. We showed that in this model, in contrast to the Dirac-Delta MF, massive MACHOs contribute in the microlensing events more frequently than the low-mass ones do. To quantify our argument we defined two mass scales, the active mean mass of MACHOs as the mean mass of lenses that can be observed by the gravitational microlensing experiment and the passive mean mass of MACHOs as the overall mean mass of them. We showed that the active mean mass of MACHOs is always larger than the passive mean mass, except in the case of a uniform Dirac Delta MF where they are equal. \\\\ To test the compatibility of this model with the observed microlensing events, we compared the duration distribution of the events in this model with the LMC candidates of MACHO experiment. We used two statistical parameters - the mean and the width of the duration of events - to compare the observed data with the theoretical models. We showed that amongst power-low halo models some of them with a Dirac-Delta MF are compatible with the data, while in the case KE model, almost none of them are compatible with the data. The best parameters for the KE model were obtained with likelihood analysis. In the spatially varying MF using the new parameters, some Galactic models (such as standard model and models A, B, C and D) were compatible with the data. The hypothesis of a spatially varying MF of MACHOs may be tested by measuring the proper motions of white dwarfs in the Galactic halo \\cite{tor02}.\\\\  ACKNOWLEDGMENT\\\\ The author thanks David Bennett for his useful comments on blending correction of the duration of events and Sepehr Arbabi and Mohammad Nouri-Zonoz for reading the manuscript and giving useful comments."
    },
    "0404/astro-ph0404062_arXiv.txt": {
        "abstract": "Type Ia supernova data have recently become strong enough to enable, for the first time, constraints on the time variation of the dark energy density and its equation of state. Most analyses, however, are using simple two or three-parameter descriptions of the dark energy evolution, since it is well known that allowing more degrees of freedom introduces serious degeneracies. Here we present a method to produce uncorrelated and nearly model-independent band power estimates of the equation of state of dark energy and its density as a function of redshift. We apply the method to recently compiled supernova data.  Our results are consistent with the cosmological constant scenario, in agreement with other analyses that use traditional parameterizations, though we find marginal (2-$\\sigma$) evidence for $w(z)<-1$ at $z<0.2$. In addition to easy interpretation, uncorrelated, localized band powers allow intuitive and powerful testing of the constancy of either the energy density or equation of state.  While we have used relatively coarse redshift binning suitable for the current set of $\\sim$ 150 supernovae, this approach should reach its full potential in the future, when applied to thousands of supernovae found from ground and space, combined with complementary information from other cosmological probes. ",
        "introduction": "Recent measurements of the distance-redshift relation using type Ia supernovae (SNe Ia) \\cite{Rieetal04,Baretal04} obtained using the Hubble Space Telescope further strengthened the evidence that the rate of expansion of the universe is increasing in time \\cite{perlmutter-1999}. This accelerated expansion is ascribed to a mysterious component called dark energy that comprises about 70\\% of the energy density of the universe. In addition to supernova data, additional pieces of evidence come from the combined study of the large scale structure and the cosmic microwave background anisotropy measurements \\cite{Tegetal03}.  While the presence of dark energy is by now well established, we are at an early stage of studying and understanding this component. It is hoped that more accurate cosmological measurements will further constrain parameters describing dark energy and eventually shed light on the underlying physical mechanism.  Dark energy is most simply described by its present day energy density relative to the critical value, $\\Omega_{DE}$, and its equation of state defined as the ratio of pressure to density, $w\\equiv p_{DE}/\\rho_{\\rm DE}$ \\cite{turner_white}. In general, $w$ is allowed to freely vary with time (or redshift), as is $\\rho_{DE}$. In practice, it is difficult to constrain $w(z)$ or, say, the scaled energy density $f(z)\\equiv \\rho_{DE}(z)/\\rho_{DE}(0)$, when they are described by more than a few parameters due to severe parameter degeneracies entering the observable quantity (luminosity distance, in the case of SNe Ia). Even though it is in principle possible to recover the function $f(z)$ or $w(z)$ directly from supernova measurements \\cite{Star98}, in practice one has to fit the noisy data with a smooth functional form \\cite{HutTur01} which introduces error and bias (for a valiant attempt to do this with current data, see \\cite{DalDjo04}).  Another general approach is to model $w$ or $f$ using a cubic spline in redshift (e.g. \\cite{WanTeg04}), but again the paucity of data limits the spline to a few points in redshift, while having more points would correlate the measurements making the interpretation somewhat difficult.  Constraints from the new SN Ia data~\\cite{Rieetal04} suggest that dark energy is consistent with the cosmological constant scenario \\cite{Rieetal04,WanTeg04}, agreeing with previous work \\cite{Alaetal03,WanFre04}.  However, these (and other) analyses are typically based on particular models --- either a linear variation with redshift \\cite{CooHut99} or the evolution that asymptotes to a constant $w$ at high redshift \\cite{Lin03}, or perhaps a more complicated parameterization \\cite{CorCop02} --- that are used to describe redshift variation of the dark energy equation of state.  While these forms do a very good job in fitting $w(z)$ due to a variety of proposed mechanisms that could be responsible for dark energy \\cite{Lin04}, one should keep in mind that we are far from having any solid leads as to what to expect for the dark energy evolution. Given the constant increase in the quality and quantity of SN Ia data, it is timely to consider whether one can use current data to derive model independent conclusions on the evolution of dark energy.  In this paper, we introduce a variant of the principal component analysis advocated in Ref.~\\cite{HutSta03}. We make use of the most recent type Ia supernova data from Ref.~\\cite{Rieetal04} and present a view of dark energy complementary to other approaches. At the same time, we are seeking to answer one of the most important questions at present: is dark energy consistent with the cosmological constant scenario or not? Our analysis is facilitated by the fact that our measurements are completely uncorrelated. Finally, we briefly comment on the applicability of this approach to future datasets. Throughout we assume a flat universe. ",
        "conclusions": "\\label{sec:Discussion}  We have used a variant of the principal component technique to produce uncorrelated, nearly model independent estimates of the equation of state of dark energy $w(z)$ and its scaled energy density $f(z)$.  We used four redshift bins in each case, and found results that are in good agreement with previous analyses.  We further argued that the present approach nicely complements other methods that use conventional parameterizations of $w(z)$ and $f(z)$. Given that our band powers are uncorrelated, the interpretation of the cumulative evidence is particularly easy.  If dark energy is due to the cosmological constant, then $w=-1$ and $f(z)=1$, and all of our band powers should be consistent with those values, independently of their window functions. Conversely, if we ever find strong statistical evidence that {\\it even just one} band power is different from $-1$ (for $w$) or $1$ (for $f$), we will have ruled out the cosmological constant scenario. While we do find a hint of such evidence in the first band power of $w(z)$, a definitive analysis will have to await more data and a careful assesssment of the systematics.  While we have presented an analysis with $\\sim$ 150 supernovae and restricted ourselves to four bins in redshift, the generality and power of this method should make it perfectly suitable for the analysis of future supernova datasets, when the error bars are expected to improve by up to an order of magnitude and enable a much more quantitative analysis and comparison with models. Furthermore, the same techniques can be applied to a variety of other cosmological probes, as one can expect that their complementarity will considerably strengthen the SN Ia results. Finally, one can customize the proposed technique specifically to maximize the return of any given test (say, whether $w(z)=-1$ or not).  With an increase in the number of type Ia supernovae at high redshift, it is likely that these interesting possibilities will be considered in the future.   {\\it Acknowledgments:} We thank Eric Linder for useful comments on the manuscript. This work has been supported by the DOE at Case Western Reserve University (DH) and the Sherman Fairchild foundation and DOE DE-FG 03-92-ER40701 at Caltech (AC)."
    },
    "0404/astro-ph0404312_arXiv.txt": {
        "abstract": "A \\chandra\\ image of the central X-ray source \\rxj\\ in the supernova remnant CTA~1 reveals a point source, a compact nebula, and a bent jet, all of which are characteristic of energetic, rotation-powered pulsars. Using the MDM 2.4m telescope we obtain upper limits in the optical at the position of the point source, \\pos, determined to an accuracy of $0.\\!^{\\prime\\prime}1$, of $B > 25.4$, $V > 24.8$, $R > 25.1$, and $I > 23.7$; these correspond to an X-ray-to-optical flux ratio $\\ga 100$. Neither a VLA image at 1425~MHz, nor a deep pulsar search at 820~MHz using the NRAO Green Bank Telescope, reveal a radio pulsar counterpart to an equivalent luminosity limit at 1400~MHz of 0.02~mJy~kpc$^2$, which is equal to the lowest luminosity known for any radio pulsar. The \\chandra\\ point source accounts for $\\approx 30\\%$ of the flux of \\rxj, while its compact nebula plus jet comprise $\\approx 70\\%$.  The X-ray spectrum of the point source is fitted with a power-law plus blackbody model with $\\Gamma = 1.6 \\pm 0.6$, $kT_{\\infty} = 0.13 \\pm 0.05$~keV, and $R_{\\infty} = 0.37$~km, values typical of a young pulsar.  An upper limit of $T^{\\infty}_{\\rm e} < 6.6 \\times 10^5$~K to the effective temperature of the entire neutron star surface is derived, which is one of the most constraining data points on cooling models. The 0.5--10~keV luminosity of \\rxj\\ is $\\approx 4 \\times 10^{31}\\ (d/1.4\\ {\\rm kpc})^2$ ergs~s$^{-1}$, but the larger ($\\sim 18^{\\prime}$ diameter) synchrotron nebula in which it is embedded is two orders of magnitude more luminous. These properties allow us to estimate, albeit crudely, that the spin-down luminosity of the underlying pulsar is in the range $10^{36} - 10^{37}$ ergs~s$^{-1}$, and support the identification of the high-energy $\\gamma$-ray source \\eg\\ as a pulsar even though its spin parameters have not yet been determined. ",
        "introduction": "The supernova remnant CTA~1 (G119.5+10.2) is a radio shell with a diameter of $\\approx 1\\fdg8$ \\citep*{sew95} and a center-filled X-ray morphology.  It has a kinematic distance of $d = 1.4 \\pm 0.3$~kpc derived from an associated \\ion{H}{1} shell \\citep{pin93}, and the remnant age from a Sedov analysis is $t_S \\approx 1.3 \\times 10^4$~yr \\citep{sla04}.  A detailed X-ray spectral study of the SNR using \\asca\\ data showed that the large-scale extended emission is likely of a synchrotron origin; its total luminosity is $5.6 \\times 10^{33}\\ (d/1.4\\ {\\rm kpc})^2$ ergs~s$^{-1}$ in the 0.5--10 keV band \\citep{sla97}.  The compact {\\it ROSAT\\/} PSPC source \\rxj\\ is located in the brightest part of the synchrotron nebula \\citep{sew95,sla97,bra98}.  \\rxj\\ has the requisite properties for a rotation-powered pulsar.  Its X-ray luminosity in the 0.1--2.4~keV band is $4.3 \\times 10^{31}\\ (d/1.4\\ {\\rm kpc})^2$ ergs~s$^{-1}$ \\citep{sla97}, typical of young pulsars, but atypically small compared to the total luminosity of the large synchrotron nebula in CTA~1. Appealing to an empirical relation between X-ray luminosity and spin-down power $\\dot E$, \\citet{sla97} argued that the current spin-down power of the pulsar, required to account for the present nonthermal emission from the entire synchrotron nebula, is $\\dot E = 1.7 \\times 10^{36}$ ergs~s$^{-1}$. Combining this with the Sedov age predicts a rotation period $P \\la 0.17$~s and a dipolar surface magnetic field strength $B_{\\rm p} \\la 6 \\times 10^{12}$~G.  Within the boundary of CTA~1 lies \\eg, one of the brighter unidentified EGRET sources.  With a 95\\% confidence localization diameter of only $28^{\\prime}$ and an intermediate Galactic latitude of $+10\\fdg2$, this EGRET/SNR coincidence is one of the most convincing among the unproven identifications \\citep{bra98}, especially because of the likely pulsar in \\rxj.  The $\\gamma$-ray source itself has all of the characteristics of the known pulsars. It is not variable; other $\\gamma$-ray pulsars show little if any change in flux while most blazars, the other major class of EGRET source, are often dramatically variable. Its photon spectral index of $1.85 \\pm 0.10$ \\citep{har99} ($1.58 \\pm 0.18$ between 70~MeV and 2~GeV, \\citealt{bra98}), is similar to other EGRET pulsars, and flatter than that of most blazars.  The estimated values of $d, t_S, \\dot E$, and $P$ are all typical of what one would expect for a young $\\gamma$-ray pulsar.  In particular, the distance to CTA~1 is typical for the predicted EGRET pulsar population \\citep*[e.g.,][]{hal93,kaa96,rom95,yad95,yad97}, and the inferred energetics are comparable to those of the Vela pulsar and other sources with similar values of $\\dot E$.  Recent observations with \\xmm\\ resolved \\rxj\\ into a point source and a diffuse nebula, and found a two-component (blackbody plus power-law) spectral fit that is consistent with a young pulsar \\citep{sla04}. Here, we present a higher resolution \\chandra\\ observation of these features that further resolves out a prominent jet-like structure and possible evidence of a torus, features typical of pulsar wind nebulae (PWNe) found around the most energetic of pulsars. In \\S 2 we present our \\chandra\\ imaging and spectroscopic results, in \\S 3 the optical observations, and in \\S 4 new radio imaging and deep radio search for pulsations from \\rxj.  In \\S 5 we discuss the constraints on $\\dot E$ for the putative pulsar in CTA~1 based on its X-ray and (likely) $\\gamma$-ray luminosities. Throughout this paper, for clarity, we refer to the small-scale nebula resolved by \\chandra\\ within \\rxj\\ as the PWN, as distinct from the large synchrotron nebula, which is also likely wind fed. ",
        "conclusions": "The most important unknown property of \\rxj\\ is its spin-down luminosity.  Predictions of $\\dot E$ typically make use of correlations between $L_X$ and $\\dot E$ among known pulsars \\citep*[e.g.,][]{sew88,pos02}, but this method is fraught with uncertainty, as well as special difficulties in the case of \\rxj.  The 0.5--10~keV luminosity of \\rxj\\ that we measure with \\chandra, including the point and diffuse components, is $4.0 \\times 10^{31}$ ergs~s$^{-1}$.  This is slightly smaller than the value of $5.4 \\times 10^{31}$ ergs~s$^{-1}$ measured by \\citet{sla04} using \\xmm, the difference probably attributable to an additional contribution of diffuse emission contained in the larger \\xmm\\ extraction region. The spin-down luminosity predicted by \\citet{sla04} using the \\citet{sew88} relation is then $\\dot E \\approx 5 \\times 10^{34}$ ergs~s$^{-1}$, comparable to the values found for older $\\gamma$-ray pulsars such as Geminga, which has an $\\dot E = 3.3 \\times 10^{34}$ ergs~s$^{-1}$. This interpretation is troubling in its neglect of the large synchrotron nebula, for which there is no proposed power source other than the pulsar.  It is impossible to measure the entire synchrotron nebula of CTA~1 with \\chandra\\ because of its large size, low surface brightness, and lack of a clear boundary.  But since the total synchrotron emission seen by \\asca\\ exceeds the compact source luminosity by more than two orders of magnitude, consideration of the total energy budget would favor a  minimum value of $\\dot E \\approx 1.7 \\times 10^{36}$ ergs~s$^{-1}$, as originally proposed by \\citet{sla97}.  We may also look to more detailed X-ray studies of pulsars with PWNe to estimate the spin-down power of CTA~1. From an empirical perspective, bright PWNe are found only around pulsars with $\\dot E > 3 \\times 10^{36}$ ergs~s$^{-1}$ \\citep{got04}.\\footnote{Extended emission recently found around the nearby ($d \\sim 160$~pc) Geminga pulsar \\citep{car03} has $L_X$ of only $\\sim 6.5 \\times 10^{28}$ ergs~s$^{-1}$, a mere $2 \\times 10^{-6}$ of its $\\dot E$ or 2\\% of its $L_X$. Such weak emission would not have been detected at the distance of CTA~1.} Above this threshold, PWN emission typically exceeds the pulsar point-source luminosity by an order of magnitude. Furthermore, the spectral index of the pulsar in the 2--10~keV band correlates with $\\dot E$ according to the \\cite{got03} relation $\\Gamma_{\\rm PSR} = 2.08 - 0.29\\,\\dot E_{38}^{-1/2}$, where $\\dot E_{38}$ is the spin-down power in units of $10^{38}\\ {\\rm ergs\\ s}^{-1}$ Applying this to CTA~1 for which $\\Gamma_{\\rm PSR} = (1.1-2.2)$ suggests that $\\dot E \\ga 6 \\times 10^{36}$ ergs~s$^{-1}$, allowing for scatter in the relation.  Further indication that $\\dot E$ is in this higher range is suggested by the $\\gamma$-ray luminosity of \\eg, which provides a lower-limit on the spin-down power. At $d = 1.4$~kpc, the $\\gamma$-ray flux integrated above 100~MeV corresponds to $L_{\\gamma} = 6 \\times 10^{34}$ ergs~s$^{-1}$ if isotropic, or $5 \\times 10^{33}$ ergs~s$^{-1}$ if beamed into 1~steradian as is often (but arbitrarily) assumed.  Thus, a prediction of $\\dot E \\approx 5 \\times 10^{34}$ ergs~s$^{-1}$ is incompatible with isotropic $\\gamma$-ray emission, and requires a substantial degree of beaming. On the other hand, if $\\dot E \\approx 6 \\times 10^{36}$ ergs~s$^{-1}$, then the ratio $L_{\\gamma}/\\dot E$ is only 0.01 in the isotropic case, which is similar to Vela, whose $\\dot E = 6.9 \\times 10^{36}$ ergs~s$^{-1}$.  Note that this argument is independent of whether the X-ray spectrum of \\rxj\\ can be extrapolated over four orders of magnitude in energy to join the $\\gamma$-ray spectrum.  Such a connection was suggested by \\citet{sla04}, but it is impossible to verify, as well as unnecessary for associating the sources.  The $0.5-10$~keV power law of the Geminga pulsar clearly {\\it does not\\/} connect with its $\\gamma$-ray spectrum \\citep{jac02}; rather, they are separate spectral components arising from different emission mechanisms."
    },
    "0404/astro-ph0404124_arXiv.txt": {
        "abstract": "{New Effelsberg 1.4~GHz observations of the central $10\\degr \\times 10\\degr$ of the Virgo cluster are presented. NVSS data are used to subtract point sources from our map. During the data reduction process special care is taken (i) to disentangle emission from the North Polar Spur from emission from the Virgo cluster, (ii) to disentangle emission from the strong M87 sidelobes from emission from the Virgo cluster, and (iii) to correct for non-linear ground emission due to the long scans. We detect a low surface brightness radio halo with a flux density of $5\\pm1.5$~Jy centered close to the elliptical galaxy M86. This halo is much weaker than that observed in the Coma cluster. It is reminiscent of a past interaction between the intracluster medium of M86 and a low density gas, belonging most probably to the Virgo cluster. \\keywords {Galaxies: clusters: individual: Virgo -- intergalactic medium -- Radio continuum: general}}  \\titlerunning{The Virgo cluster radio halo} ",
        "introduction": "}  The location of the largest diffuse radio emission sources in the local Universe are galaxy clusters. A very hot (several $10^{7}$~K) and tenuous ($n \\sim 10^{-4}$~cm$^{-3}$) gas pervades the whole cluster. This gas can be directly observed in X-rays (see e.g. Ebeling et al. 1996). The existence of a cluster-wide diffuse radio emission indicates that there are additional, important components of the intracluster medium (ICM): large-scale magnetic fields and relativistic electrons. In general these extended, diffuse radio sources are divided into two subgroups (see e.g. Giovannini \\& Feretti 2002a): (i) cluster halos have low surface brightness, are centered on the cluster center, and show a regular shape. Their prototype is Coma~C (Deiss et al. 1997). They have a steep spectrum and show a negligible degree of polarization. Their typical size is $\\sim 1$~Mpc. (ii) Relic sources also show low surface brightness and steep spectra, but in contrast to radio halos they are elongated structures that are found at the periphery of the cluster. In addition, their degree of polarization is high. A3667 (R\\\"{o}ttgering et al. 1997) represents an extreme member of this group.  Diffuse, extended radio sources are rare. Giovannini et al. (1999), using a complete cluster sample, found that 5\\% of clusters have radio halos and 6\\% show relic sources. The X-ray luminosity of clusters hosting a diffuse radio source is significantly higher than that of clusters without such a source (Giovannini \\& Feretti 2002b). The number of known diffuse radio sources has significantly increased to about 40 in the last years. Typical radio powers of diffuse sources are $10^{24} - 10^{25}$~W\\,Hz$^{-1}$ at 1.4~GHz. The strength of the magnetic fields calculated using equipartition between the relativistic electrons and the magnetic field is $0.1 - 1$~$\\mu$G.  There is strong evidence that diffuse radio sources are related to cluster mergers. Significant substructure in the X-ray emission is detected in clusters hosting a radio halo/relic (B\\\"{o}hringer \\& Sch\\\"{u}cker 2002). At present a merger scenario could not be clearly excluded for a cluster showing a radio halo (Giovannini \\& Feretti 2002a). The formation of halos and relics is tightly linked to the dynamical behaviour of the cluster. Whereas the magnetic fields and relativistic electron populations are built during the cluster formation and subsequent evolution, the energy supply for the maintenance of the diffuse radio sources is due to recent gravitational interactions between the cluster and an infalling galaxy group or cluster.  One of the best studied cluster radio sources is Coma (see e.g. Deiss et al. 1997), which is located at $\\sim$100~Mpc distance. Its extent on the sky is about $1\\degr$. In order to detect a large-scale, diffuse emission in a galaxy cluster with a single-dish telescope, additional high angular resolution observations at the same wavelength must be available to identify point sources, which have to be subtracted. It must also be taken into account that the region of interest should not be too large, because the low surface brightness of diffuse radio sources implies huge amount of observation time at high frequencies. In addition, observations of large areas may have problems in the determination of a constant baselevel. Because of these reasons, it is only now that we attempt to detect a diffuse radio source in the Virgo cluster with the Effelsberg 100-m telescope. The Virgo cluster, with a distance of 17~Mpc\\footnote{We use a Hubble constant of $H_{0}=70$~km\\,s$^{-1}$\\,Mpc$^{-1}$.}, is the nearest galaxy cluster on the northern hemisphere. We cover the central $10\\degr \\times 10\\degr$.  Virgo is a dynamically young cluster, i.e. the galaxy distribution is not symmetric (see e.g. Schindler et al. 1999). Several subgroups of galaxies are located closely in projected distance and velocity to the Virgo cluster. The galaxy distribution shows mainly a north-south and an east-west elongation. The ICM distribution observed with ROSAT in X-rays (B\\\"{o}hringer et al. 1994) shows a symmetric component around M87 and an asymmetric, low surface brightness emission elongated in north-south direction. The main part of this asymmetric component comes from the galaxy group around the elliptical galaxy M49 located in the south of M87. A small X-ray halo is also associated with M86, a big elliptical galaxy located to the west close to M87.  These asymmetries imply a high dynamical activity. At the same time the Virgo cluster hosts a central cooling flow (White \\& Sarazin 1988). Since it is thought that a major cluster merger disrupts a central cooling flow (see White et al. (1993) for the Coma cluster), the existence of a symmetric X-ray halo and a cooling flow around M87 thus implies that the Virgo cluster did not experience a major merger in the recent past. Nevertheless, this might happen in the near future (Tully \\& Shaya 1984).  From the central cD galaxy M87 a distorted radio jet emanates, surrounded by a radio minihalo (Andernach et al. 1979, Rottman et al. 1996, Owen et al. 2000). Its linear extent is about 80~kpc. This minihalo is directly related to the radio jet of M87 and can thus not be qualified as a radio halo as defined above.  We have attempted to detect diffuse radio emission on a Mpc scale in the Virgo cluster at 1.4~GHz with the Effelsberg 100-m telescope. When conducting a search for a presumably weak radio halo from the Virgo cluster, the presence of the quite bright radio source Virgo~A (3C274, M87) must be taken into account. At 1.4~GHz the peak flux density is about 150~Jy when measured with the Effelsberg telescope. Such a strong signal produces antenna sidelobes within the area of the cluster what must be adequately taken into account. Another complication arises from the North Polar Spur (NPS), an intense Galactic radio shell, which is likely the old remnant from a nearby supernova explosion. Its minimum distance was previously estimated to be about 100~pc. Its diameter is about $116\\degr \\pm 4\\degr$ (Berkhuijsen et al. \\cite{elly}). The NPS is seen in projection  towards the south-eastern area of the Virgo cluster. Virgo~A is close to the outer periphery of the NPS as shown in Fig.~\\ref{fig:NPS}, where a section of the 1.42~GHz northern--sky Stockert survey (Reich \\& Reich \\cite{pr}) is displayed to illustrate the unfavourable situation. However, this absolutely calibrated Stockert survey data may be helpful in a separation of the different contributions.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{0783fig1.ps}} \\caption{Section of the 1.4~GHz Stockert continuum survey showing a part of the North Polar Spur close to the Virgo cluster with Virgo~A or M87 at its center. The greyscale range is between 3.3~K (white) and 3.7~K (black). Intensities above 3.7~K are clipped (shown white). The map is in the equatorial coordinate system and centered at $\\rm 12^{h}57^{m}18.^{s}4, 12\\degr24\\arcmin24\\arcsec$ (J2000) or $6\\degr$ to the east of Virgo~A. The $10\\degr \\times 10\\degr$ field observed with the Effelsberg telescope is indicated.} \\label{fig:NPS} \\end{figure}  The observations, data reduction, and the results are described in Sect.~\\ref{sec:data} followed by the discussion in Sect.~\\ref{sec:discussion} and the our conclusions (Sect.~\\ref{sec:conclusions}). ",
        "conclusions": "}  We present Effelsberg 1.4~GHz observations of the central $10\\degr \\times 10\\degr$ of the Virgo cluster, in order to investigate if the cluster hosts a diffuse radio halo. Three major obstacles have to be overcome to detect a large-scale, low surface brightness structure: (i) the North Polar Spur (NPS) extends into the Virgo cluster region, (ii) the high flux density of the central cD galaxy M87 leads to strong sidelobes, and (iii) the ground radiation due to the large scans. The NVSS survey is used to subtract the point sources from our total intensity map. In this way we decrease the rms noise level from 7~mJy/beam at a resolution of $9\\farcm4$ to $3$~mJy/beam at a resolution of $12\\arcmin$. In the total intensity and polarized intensity map the NPS appears as elongated structures running along the north-east towards south-west direction in the south-eastern part of the maps. Parts of the NPS extend up to the location of M87. We corrected for the non--linear ground radiation in three ways: (i) we subtracted the observed mean radial profile, (ii) we subtracted an analytical profile, and (iii) we corrected for the large-scale structure with the help of absolutely calibrated 1.4~GHz observations of lower resolution. The remaining sidelobe confusion is analyzed using a holographic, deep antenna pattern at 11.7~GHz using the EUTELSAT satellite.  \\noindent It is concluded that \\begin{enumerate} \\item we do not detect a bright, large-scale radio halo, as it is observed in the Coma cluster. \\item We detect a radio halo around the elliptical galaxy M86 with an estimated radial extent of $\\sim 2\\degr$ and an estimated total flux density of $5\\pm1.5$~Jy. \\item The detection of the M86 radio halo is robust. \\item The radio halo is a witness of a past interaction of the M86 ICM with low density gas at the outskirts of the Virgo cluster. \\item This leads to a scenario where the ICM of M86 had an interaction $\\sim 1$~Gyr ago with a low density gas, whereas M86 itself undergoes an interaction with a relatively high density gas today (Finoguenov et al. 2004). \\end{enumerate}"
    },
    "0404/astro-ph0404197_arXiv.txt": {
        "abstract": "We are conducting a high angular resolution imaging survey of Galactic Wolf-Rayet stars using the Wide Field and Planetary Camera 2 aboard the {\\it Hubble Space Telescope}. We have found a small stellar cluster associated with the faint, close pair WR 38 and WR 38a. We present astrometric measurements and photometry in the wide-band F336W ($U$), F439W ($B$), and F555W ($V$) filter system for these cluster and nearby stars.   We combine our photometry with Johnson and IR magnitudes, and compare the observations with calibrated model results for reddened stars to adjust the HST zeropoints and to identify five probable main sequence members of the cluster. A least-squares fit of the colors and magnitudes of this set yields a cluster reddening of $E(B-V) = 1.63 \\pm 0.05$ mag and a distance of $10^{+12}_{-4}$ kpc for an assumed ratio of total-to-selective extinction of $R=3.1$.  We discuss the relationship of this cluster to other objects along the line of sight.  If situated at a distance of $\\approx 8$~kpc, then the cluster would reside in a dense region of the Carina spiral arm, close to a giant molecular cloud and the starburst cluster NGC~3603. ",
        "introduction": "%  Wolf-Rayet (WR) stars are the evolved, He-burning cores and hot stellar envelopes of once-massive O-type stars. Their strong winds (terminal velocities: $v_{\\infty}$~$\\approx$~1000~-~2500~km~s$^{-1}$) and heavy mass loss rates ($\\dot{M}$~=~$\\sim[1-6]~\\times$~10$^{-5}$~$M_{\\odot}$~yr$^{-1}$; \\citealt{wil91}; \\citealt{nug02}) power their characteristic strong, broad emission lines of He, N, C, and O, which correspond to increasingly evolved evolutionary states.  While the environment and multiplicity of O-type stars have been explored in depth \\citep{gie87,mas98}, these characteristics for the evolved WR phase are not as well determined. In 1996 we initiated a survey with the {\\it Hubble Space Telescope} Wide Field and Planetary Camera 2 ({\\it HST} WFPC2) to discover and quantify the multiplicity and environments of Galactic WR stars.  Here we discuss two of our targets, WR~38 and WR~38a, that were revealed to be both multiple and contained within a small stellar cluster \\citep{wal99}.  WR~38 is a WC4-type star \\citep{mac70} and WR~38a is a WN5-type \\citep{sha91}.  The two are separated by only $20\\arcsec$ in the sky, and they are neighbors to another even tighter $5\\arcsec$ pair of WR stars, WR~38b and WR~39, at a separation of $3\\farcm8$.  These four plus WR~37 are surrounded by a ring nebula \\citep{mar97}, although it is not clearly associated with any one of these stars.  Our view in this direction in the Galactic plane cuts through the Sagittarius-Carina spiral arm at distances of approximately 2.4 and 8~kpc \\citep{gra88,geo00,rus03}.  The photometric distances of these WR stars are generally close to 5~kpc \\citep{van01}, which would place them among the young population of the Carina spiral arm.  However, in a recent paper \\citet*{sho04} presented photometry for the stars in the immediate vicinity of WR~38 and WR~38a that indicated a much greater distance of $14.5 \\pm 1.6$~kpc.  This distance would place the pair in an extension of the outer Perseus arm some 14~kpc from the Galactic center.  There are only 4 of 227 known Galactic WR stars that have such a great distance from the Galactic center \\citep{van01}. \\citet{sho04} designate the cluster as C1104--610a, and images of the field surrounding the cluster are presented in \\citet{wra76} (Plate~4), \\citet{van81}, and \\citet{sho04} (Fig.~1 and 2).  Here we present new astrometric and photometric results on the stars found close to WR~38 and WR~38a based upon {\\it HST} PC images with superb angular resolution.  We identify which stars are associated with the WR pair based upon their derived reddening and location in a color-magnitude diagram.  Our results are mainly consistent with those of \\citet{sho04} but we argue that the acceptable distance range is such that the WR pair and its surrounding cluster may reside in the distant portion of the Carina spiral arm. ",
        "conclusions": ""
    },
    "0404/gr-qc0404119_arXiv.txt": {
        "abstract": "We provide a brief chronological guide to the literature on non-Riemannian cosmological models. Developments in this field are traced back to the early seventies and are given in table form. ",
        "introduction": "\\label{sec_introduction} Nowadays, cosmology is a prospering and rapidly changing field of physics. One could say that cosmology has entered a {\\it golden age} with a wealth of new experimental data available, cf. the other contributions in this proceedings. In addition theoretical efforts have led to what we call the {\\it standard cosmological model} or {\\it cosmological concordance model}. Within this model we are able to explain most of the observations. But, maybe as a trade-off for simplicity, we have to introduce concepts like dark matter and dark energy within this picture. Unfortunately, so far there has been {\\it no} direct detection of a dark matter particle and there seems to be {\\it no} deeper theoretical justification for the large amount of dark energy. This is clearly an embarrassing situation because we cannot tell what over 90\\% of the universe is really made of.  Since cosmology combines concepts of many different fields of physics there exist several distinct approaches to find a remedy for this unsatisfactory situation. One of them is to change the spacetime geometry and therewith the underlying gravity theory. In this report we want to focus on the history of non-Riemannian cosmological models, which we track back to the early seventies. For reviews covering non-Riemannian gravity theories the reader is referred to \\cite{PhysRep}, \\cite{Blagojevic}, \\cite{Hammond1} and references therein. ",
        "conclusions": "Challenging experimental results, like the (non) detection of dark energy and dark matter, are often the precursors of a shifting of scientific paradigms. Non-Riemannian gravity theories seem to offer an ideal playground for cosmologists and their search for an explanation of dark matter and dark energy.  Maybe cosmology might yield the definitive clue for the presence of non-Riemannian structures in our universe thereby utilizing a nearly 85 year old theoretical idea.  \\paragraph*"
    },
    "0404/astro-ph0404532_arXiv.txt": {
        "abstract": "We determine the expected surface brightness and photometric signature of a white dwarf remnant population, issued from primordial low-mass stars formed at high redshifts, in today galactic halos. We examine the radial dependence of such a contribution as well as its redshift dependence. Such a halo diffuse radiation is below the detection limit of present large field ground-based surveys, but should be observable with the HST and with the future JWST project. Since the surface brightness does not depend on the distance, the integration of several galactic dark halos along the line of sight  will raise appreciably the chances of detection. Both the detection or the non-detection of such a remnant diffuse radiation within relevant detection limits offer valuable information on the minimum mass for star formation in the early universe and on the evolution of the stellar initial mass function. ",
        "introduction": "Modern theories of structure formation in the early universe suggest that galaxies have formed out of a primordial generation of stars, with zero-like metal abundances, the so-called Population III stars. Our understanding of the general properties of Pop III stars has improved significantly over the past recent years (see Barkana \\& Loeb 2001, Bromm \\& Larson 2004 for recent reviews). The mass distribution of this first generation(s) of stars, however, remains ill-determined. According to recent hydrodynamical models (Abel, Bryan \\& Norman 2000, Bromm, Copi \\& Larson 1999, 2002) the first stars are very massive ($\\simgr 100\\,\\msol$), primarily because of the poor efficiency of the dominant cooling mechanism, due to molecular hydrogen (H$_2$) rotational transitions. Other calculations, however, find that the formation of Pop III stars following the very first-generation stars may extend to the $\\sim 1$-10 $\\msol$ range, depending on the protogalactic gas cloud conditions (Nakamura \\& Umemura 2001, 2002),  the effect of the intense radiation field (Omukai \\& Yoshii 2003) or on the efficiency of metal-enriched (Pop II.5) star formation due to supernovae (Mackey, Bromm \\& Hernquist 2003, Salvaterra, Ferrara \\& Schneider 2003). In that case, primordial star formation may have led to an early generation of extremely or totally (in the absence of mixing of  the ejected material in the SN-driven shell in the ISM) metal-depleted, long-live low-mass  stars. The minimum mass for star formation in the early universe thus remains uncertain. This minimum mass, however, is likely to decrease with time, as (i) the virial temperature of massive collapsing halos will exceed $T_{vir}>10^4$, inducing radiative line-cooling from atomic hydrogen, (ii) the minimum ambient temperature, set up by the cosmic background radiation, will decrease and so will the thermal Jeans mass, and (iii) a small fraction of metals will be processed in the primordial short-lived stars, yielding a much more efficient cooling mechanism (Clarke \\& Bromm 2003, Bromm \\& Loeb 2003). When the transition between Pop III+II.5, hereafter defined generically as our \"primordial star generation\", to well identified Pop II-like stars is occurring remains unknown, depending on ill-determined factors such as the effect of the ambient UV radiation on the main coolants or the efficiency of the medium heavy element enrichment due to the first SNII generations.  The evolution of the initial mass function (IMF) from the (nearly-zero $Z$) Pop III to (low-$Z$) Pop II may in fact correspond to transition from a high-mass to a low-mass fragmentation mode of star formation, illustrating the evolution of the IMF with structure formation in the universe. The recent discovery of a very metal-depleted ($[Fe/H]=-5.3\\pm 0.2]$) red giant ($m\\approx 0.8\\,\\msol$) in our Galactic halo (Christlieb et al. 2002) indeed points to such a low-mass mode of primordial star generation. Moreover, the metal enrichment of clusters of galaxies suggests an IMF at high redshift with an average stellar mass larger than for present-day conditions, in the range $\\sim 1$-$10\\,\\msol$ (Portinari et al. 2004 and references therein). Such an evolution bears important consequences on the early history and the subsequent evolution of the universe. While the UV radiation and $\\alpha$-element nucleosynthesis will be dominated by the more massive objects, a primordial, high-redshift stellar population extending to masses below about 10 $\\msol$ will (i) extend the nucleosynthetic products to C, N and Fe enriched elements as well as s-processed elements and (ii) produce a population of white dwarfs (WDs) sequestering part of the baryons. The quest for the direct detection of these remnants remains for now inconclusive, for present surveys do not reach deep enough magnitude or have a too small field of view to reach robust conclusions. An alternative observable signature of the presence of a primordial remnant population in today galactic halos would be their contribution to these halo surface brightness. N-body simulations (White \\& Springel 2000) suggest that the primordial star clusters might survive disruption during halo merging preceding galaxy formation and, through loss of kinetic energy by dynamical friction, end up in the central bulge, with low velocity. If, however, the low-mass stars are scattered by violent relaxation or tidal disruption of the cluster during halo merging, the remnants of these early stars would be now incorporated into galactic external halos. Which of these two processes dominates remains to be determined unambiguously and depends on various factors such as the mass ratio and the compactness of the mergers or the efficiency of the SN-driven winds in the clusters. In the present paper we will assume that the primordial remnant population follows the isothermal density distribution of the dark halo. If these remnants are located in the central bulge or spheroid, their surface brightness will be largely dominated by the main sequence (M-dwarfs) contribution. In that case only spectroscopic and high-proper motion detections of individual objects will allow their identification as PopIII remnants. The nature of these primordial remnants (black holes, neutron stars or white dwarfs) carries the signature of the minimum mass of the IMF at the early epoch of star formation in the universe. Since the surface brightness does not depend on the distance (see below), the integration along the line of sight of several galactic dark halos adds up each contribution, yielding eventually a detectable signature. Both the observation or the {\\it non} observation of such a stellar relic radiation would bring important information  on the low-mass end of the stellar mass spectrum and thus on the dominant process of star formation in the early universe. This is the aim of the present paper to examine this issue. ",
        "conclusions": "In this paper, we have examined the possibility for a fraction of the baryons in the universe to be trapped in the WD remnant population of a primordial generation of intermediate or low-mass stars, formed before the end of the reionization period, i.e. in the redshift range $5\\simle z \\simle 20$. For such redshift values, a non-negligible fraction of the nucleosynthetic baryons can be in the form of such WD progenitors without producing unacceptable light today. As a possible signature of such relics, we calculate their expected surface brightness both in our own halo and in external galactic halos. The expected radiation lies below the detection limits of ground-based present surveys, but the detection of this glow should be in reach of the HST or the JWST, in particular if several external galactic halos can be observed along a given line-of-sight. Both a successful or a null detection of such a diffuse light, within relevant detection limits, will provide precious information about the mode of star formation and the evolution of the IMF at early stages of galaxy formation. The background light contribution of these WDs is found to increase appreciably with redshift, suggesting possible detection in dedicated surveys, with appropriate color cuts. Contributions from red giants and main sequence stars at these redshifts, however, must be calculated consistently with chemico-evolution models, in order to examine the favored detection bands."
    },
    "0404/astro-ph0404018_arXiv.txt": {
        "abstract": "We report on the preliminary timing and spectral analysis of the High Mass X-ray Binary Pulsar SAXJ2103.5+4545 as observed with INTEGRAL during the Galactic Plan Scan of the Core Program. The source shows a hard spectrum, being detected up to 100 keV. The timing analysis performed on IBIS/ISGRI data shows evidence for a spin-up with respect to previous observations, performed in 1997 with BeppoSAX. ",
        "introduction": "SAX J2103.5+4545 is a transient HMXRB pulsar with a $\\sim$358~s pulse period discovered with the WFC on-board BeppoSAX during  an  outburst in 1997 (Hulleman et al., 1998). Its orbital period of 12.68~days has been found with the RXTE during the 1999 outburst (Baykal et al., 2000). The likely optical counterpart, a Be star with a magnitude V=14.2, has been recently discovered (Reig \\& Mavromatakis, 2003; Reig et al. 2004).  During the outburst in 1999, Baykal et al. (2002) observed for the first time, with RXTE, the transition from the spin-up phase to the spin-down regime, while the X--ray flux was declining. Indeed, the source underwent a spin-up phase during the initial part of the outburst (during which the pulse period decreased by $\\sim$0.9~s in 150~days), then the flux dropped (and the pulse frequency saturated), and, as the flux continued to decline, a weak spin-down phase started. Moreover, a correlation between spin-up rate and X--ray flux was observed (Baykal et al., 2002), suggestive of the formation of an accretion disk during the periastron passage.  A very preliminary spectral analysis of INTEGRAL public observations of the source region, performed in Dec.2002 during the performance verification phase, has been reported by Lutovinov et al. (2003).  Inam et al. (2004) observed a soft spectral component (blackbody with a temperature of 1.9~keV) and a transient 22.7~s QPO during a XMM-Newton observation performed in Jan, 2003.  SAX J2103.5+4545 has been observed several times during the Galactic Plane Scan (GPS) which INTEGRAL performs every 12 days as part of the Core Program. We report here the timing and spectral analysis of these observations.   \\begin{small} \\begin{figure*} \\centering \\includegraphics[width=13.0cm]{isgri.soft.eps} \\vskip -3cm \\caption{IBIS/ISGRI mosaic of the source field in the 20--40~keV energy range in galactic coordinates} \\end{figure*} \\end{small}   An IBIS/ISGRI mosaic of the region of the sky containing SAX J2103.5+4545 is shown in Fig.~1. The source long-term lightcurve, as measured with RXTE All Sky Monitor (ASM) is shown in Fig.~2. The times of the two previous outbursts observed with BeppoSAX (in 1997) and RXTE (in 1999) are indicated, as well as the epoch of the INTEGRAL observations reported here.   \\begin{small} \\begin{figure} \\centering \\includegraphics[width=8.0cm]{sax2103_asm_lc_lab.ps} \\caption{RXTE ASM lightcurve of SAX~J2103.5+4545. Times of the INTEGRAL observations have been marked} \\end{figure} \\end{small} ",
        "conclusions": ""
    },
    "0404/hep-ph0404231_arXiv.txt": {
        "abstract": "\\PRE{\\vspace*{.3in}} We investigate supergravity models in which the lightest supersymmetric particle (LSP) is a stable gravitino.  We assume that the next-lightest supersymmetric particle (NLSP) freezes out with its thermal relic density before decaying to the gravitino at time $t \\sim 10^4 - 10^8~\\s$.  In contrast to studies that assume a fixed gravitino relic density, the thermal relic density assumption implies upper, not lower, bounds on superpartner masses, with important implications for particle colliders.  We consider slepton, sneutrino, and neutralino NLSPs, and determine what superpartner masses are viable in all of these cases, applying CMB and electromagnetic and hadronic BBN constraints to the leading two- and three-body NLSP decays.  Hadronic constraints have been neglected previously, but we find that they provide the most stringent constraints in much of the natural parameter space.  We then discuss the collider phenomenology of supergravity with a gravitino LSP.  We find that colliders may provide important insights to clarify BBN and the thermal history of the Universe below temperatures around 10 GeV and may even provide precise measurements of the gravitino's mass and couplings. ",
        "introduction": "\\label{sec:introduction}  Supersymmetric theories predict the existence of a spin 3/2 particle, the gravitino, the partner of the spin 2 graviton.  The gravitino mass is \\begin{equation} \\mgravitino \\sim \\frac{F}{\\mstar} \\ , \\end{equation} where $F$ is the scale of supersymmetry breaking, and $\\mstar = (8 \\pi G_N)^{-1/2} \\simeq 2.4 \\times 10^{18}~\\gev$ is the reduced Planck mass.  The masses of scalar superpartners are derived from terms such as \\begin{equation} \\lambda_{ij} \\int d^4\\theta \\frac{Z^\\dagger Z \\Phi_i^\\dagger \\Phi_j}{\\mmed^2} \\ , \\end{equation} where $\\lambda_{ij}$ are unknown constants, $Z$ is a superfield whose auxiliary component develops the vacuum expectation value $F$, $\\Phi_i$ are standard model superfields, and $\\mmed$ is the mass scale of the interactions that mediate supersymmetry breaking.  Similar terms give the spin $1/2$ superpartners mass.  In supergravity, the mediating interactions are gravitational, and so $\\mmed \\sim \\mstar$, $F \\sim (10^{10}~\\gev)^2$, and the gravitino and all standard model superpartners have mass $\\sim F/\\mstar \\sim \\mweak$, with the precise ordering determined by unknown constants, such as $\\lambda_{ij}$.  Most studies of supergravity have assumed, either explicitly or implicitly, that the lightest supersymmetric particle (LSP) is a standard model superpartner.  This avoids potential complications resulting from the decay of standard model superpartners to a gravitino LSP, which naturally happens at time $t \\sim 10^4 - 10^8~\\s$, well after Big Bang nucleosynthesis (BBN).  However, the phenomenology and cosmology of gravitinos have also been considered in a number of studies~\\cite{Pagels:ke,Weinberg:zq,Krauss:1983ik,% Nanopoulos:1983up,Khlopov:pf,Ellis:1984eq,Ellis:1984er,% Juszkiewicz:gg,Ellis:1990nb,Moroi:1993mb,Bolz:2000fu,Khlopov:rs}. (See also related studies of axino and quintessino dark matter~\\cite{Covi:1999ty,Covi:2001nw,Bi:2003qa,Hooper:2004qf}.)  More recently, it has been shown that the gravitino LSP possibility does not destroy the beautiful predictions of BBN even when the gravitino LSPs produced in late decays have relic density $\\Omega_{\\gravitino} = 0.23$ and so are present in sufficient numbers to account for all of dark matter~\\cite{Feng:2003xh,Feng:2003uy,threebody}.  In fact, bounds from the cosmic microwave background (CMB) and, in some corners of parameter space, entropy production and the diffuse photon spectrum may be even more severe than bounds from BBN~\\cite{Feng:2003xh,Feng:2003uy}.  Nevertheless, all of these bounds were shown to be respected for some regions of parameter space with weak scale superpartners.  The possibility of superweakly-interacting massive particle (superWIMP) gravitino dark matter from NLSP decays thus appears to be viable.  The analogous scenario in extra dimensional theories~\\cite{Feng:2003xh,Feng:2003uy,Feng:2003nr}, as well as interesting astrophysical implications in this and related scenarios~\\cite{Bi:2003qa,Chen:2003gz,Sigurdson:2003vy} have also been discussed.  In this work, we take an approach that differs from the exploration of superWIMP gravitino dark matter.  Instead of assuming that gravitinos are the dark matter with $\\Omega_{\\gravitino} = 0.23$, we assume that the NLSP reaches its thermal relic density $\\Omega_{\\NLSP}^{\\text{th}}$ before decaying, and so $\\Omega_{\\gravitino} = (m_{\\gravitino}/m_{\\NLSP} ) \\Omega_{\\NLSP}^{\\text{th}}$.  That is, we relax the constraint that gravitinos from NLSP decays account for all of dark matter. Rather we assume the simplest thermal history for the Universe and ask what regions of $(\\mgravitino, \\mNLSP)$ parameter space are allowed.  The thermal relic density assumption has consequences that differ markedly from the fixed gravitino relic density assumption. To see this, assume that the gravitino and NLSP masses are both parametrized by a general superpartner mass scale $m_{\\text{SUSY}}$.  As noted in Ref.~\\cite{threebody}, if one assumes a fixed gravitino relic density, the NLSP number density scales as $1/m_{\\text{SUSY}}$.  Low superpartner masses are therefore disfavored.  In contrast, if one assumes a thermal relic density for the NLSP, $\\Omega_{\\text{NLSP}}^{\\text{th}} \\propto \\langle \\sigma v \\rangle ^{-1} \\propto m_{\\text{SUSY}}^2$, where $\\langle \\sigma v \\rangle$ is the thermally-averaged NLSP annihilation cross section. The NLSP number density then scales as $m_{\\text{SUSY}}$, and so high superpartner masses are disfavored.  This difference has obviously important implications for collider searches for new physics, and we discuss collider implications below.  Even given the NLSP thermal relic density assumption, gravitinos may still be all of dark matter --- for example, if the gravitino relic density from NLSP decays is too low, the remainder may be made up by gravitinos produced during reheating.  However, the existence of such alternative gravitino sources is either untestable, or testable only with strong assumptions about the early Universe.  In contrast, the existence of a gravitino component from NLSP decays makes several robust predictions that are testable at cosmological observatories and collider experiments, and we concentrate on this gravitino source here.  Before leaving the topic of reheating altogether, however, we note that the gravitino LSP scenario has an important virtue with respect to reheating.  For stable weak-scale gravitinos, the overclosure constraint is well-known to require a bound on reheat temperature of $T_R \\alt 10^{10}~\\gev$~\\cite{Bolz:2000fu}.  Recently, however, it has been shown that if the gravitino is not the LSP, hadronic BBN constraints greatly strengthen this bound~\\cite{Kawasaki:2004yh}.  For example, if the gravitino decays to the LSP $+$ hadrons with branching fraction $10^{-3}$, the reheat temperature must satisfy $T_R \\alt 10^6~\\gev$.  This is uncomfortably low.  The gravitino LSP scenario is therefore preferred if one requires a high reheat temperature, as might be desirable, for example, for scenarios of leptogenesis~\\cite{Fujii:2003nr}.  In the present analysis, in addition to the constraints on electromagnetic energy release considered previously, we include the recent results on hadronic BBN constraints~\\cite{Kawasaki:2004yh}. The work of Ref.~\\cite{Kawasaki:2004yh} represents a significant update to previous hadronic analyses~\\cite{Reno:1987qw,% Dimopoulos:1987fz,Dimopoulos:1988ue,Kohri:2001jx}.  To include these results correctly, we must, of course, determine the leading contributions to hadronic energy.  For slepton\\footnote{Throughout this work, ``slepton'' refers to a charged slepton.} and sneutrino NLSPs, the leading contribution is from three-body decays \\begin{eqnarray} \\tilde{l} &\\to& l Z \\tilde{G} \\ , \\ \\nu W \\tilde{G} \\nonumber \\\\ \\tilde{\\nu} &\\to& \\nu Z \\tilde{G} \\ , \\ l W \\tilde{G} \\ . \\label{threebody} \\end{eqnarray} The three-body decays have been studied in Ref.~\\cite{threebody}.  For a neutralino NLSP, the leading contribution to hadronic energy is from the two-body decays, such as \\begin{equation} \\chi \\to Z \\gravitino , \\, h \\gravitino\\ , \\label{twobody} \\end{equation} followed by $Z, h \\to q \\bar{q}$.  These decays, and the hadronic constraints on them, were neglected in previous works.  As we will see, however, they are the leading constraints in much of parameter space and they are especially important when the superpartner masses and their splittings are all of the order of the weak scale, the most natural possibility.  The decays of \\eqsref{threebody}{twobody} may be suppressed kinematically if $m_{\\text{NLSP}} - \\mgravitino < m_Z, m_W$ or dynamically, as when the neutralino is photino-like.  However, even in these cases, decays such as $\\tilde{l} \\to l q \\bar{q} \\gravitino$ and $\\chi \\to q \\bar{q} \\gravitino$ are still possible at higher order. We have included estimates of these in our analysis.  These decays are in some sense ``model-independent''; even in the extreme case where the dominant decay is to invisible particles, at higher order there will be contributions to hadronic cascades from such decays.  The hadronic bounds are so constraining that these should be considered for any late decaying particle, whether a superpartner, an axino, a modulus, or other particle.  After determining the regions of parameter space allowed by cosmology, we discuss the collider signals.  The upper bounds on superpartner masses resulting from the thermal relic density assumption imply promising prospects for superpartners to be within reach of future collider experiments.  In addition, we will see that the signals of supersymmetry in gravitino LSP scenarios may be completely different {}from the conventional supersymmetry signals.  In particular, if sufficient NLSPs can be collected and monitored for decays, the NLSP lifetime may be measured, which may considerably sharpen our understanding of BBN and the thermal history of the Universe at temperatures of 10 GeV and below.  Such studies may also provide the first direct measurements of the gravitino mass and the Planck scale from particle physics~\\cite{Buchmuller:2004rq,Feng:2004gn}.  The gravitino LSP possibility, assuming a thermal NLSP relic density, has been discussed recently in the context of minimal supergravity~\\cite{Ellis:2003dn}.  Our work is complementary in that we do not work in a specific model framework, but rather consider several NLSP candidates, as well as a wide range of gravitino and NLSP masses.  Our work also differs in that we consider the hadronic constraints and the leading two- and three-body decays that contribute to hadronic energy.  As noted above, we find that these are the leading constraints in the most natural regions of parameter space. ",
        "conclusions": "In this paper, we have determined the viability of supergravity scenarios in which the gravitino is the LSP.  We have considered the three possibilities in which the NLSP is a slepton, a sneutrino, and a neutralino.  In each case, we determined the branching fractions of the leading two- and three-body decays, and applied constraints from the dark matter density, CMB, and both EM and hadronic BBN bounds.  We found that the hadronic BBN constraints, previously neglected, are extremely important, providing the most stringent limits in natural regions of parameter space.  The gravitino LSP scenario opens up many connections between particle physics and cosmology.  Consider, for example, the slepton NLSP scenario.  At colliders, it may be possible to collect the less energetic metastable sleptons in a detector and monitor this detector for slepton decays.  By measuring the decay time distribution and the energy of each produced lepton, one could independently determine both the gravitino mass and the reduced Planck mass $\\mstar$~\\cite{Buchmuller:2004rq,Feng:2004gn}.   A measurement of the gravitino mass determines the scale of supersymmetry breaking, with implications for model building and dark energy.  At the same time, such a measurement would determine a particular place in the $(\\mgravitino, m_{\\text{NLSP}})$ plane. Colliders would therefore shed light on the possible role of such new physics in BBN, and more generally in the thermal history of the Universe after NLSP freeze out at temperatures of around 10 GeV.  For example, the $\\mu$ parameter sensitivity of \\eqref{mu} may be improved by two orders of magnitude in the future by the DIMES mission~\\cite{DIMES}.  Such improvement would extend the CMB sensitivity contour significantly.  The measurement of a $\\mu$ distortion consistent with the determination of $(\\mgravitino, m_{\\text{NLSP}})$ would provide a striking confirmation of the underlying gravitino scenario.  The measurement of $\\mstar$ would provide a precision test of the supersymmetry predictions relating the properties of gravitinos to those of gravitons, and also provide the first direct measurement of the Planck scale on microscopic scales~\\cite{Buchmuller:2004rq,Feng:2004gn}.  The crucial question is the feasibility of collecting a sizable sample of metastable NLSPs. There has been an earlier study on the collection of very long lived heavy charged leptons~\\cite{Goity:1993ih}.  A more detailed analysis of the trapping of sleptons at future colliders is also now under study~\\cite{FMS}."
    },
    "0404/hep-ph0404141_arXiv.txt": {
        "abstract": "We describe an algorithm and a C++ implementation that we have written and made available for calculating the fully nonlinear evolution of $5$D braneworld models with scalar fields. Bulk fields allow for the stabilization of the extra dimension.  However, they complicate the dynamics of the system, so that analytic calculations (performed within an effective $4$D theory) are usually only reliable for static bulk configurations or when the evolution of the extra dimension is negligible. In the general case, the nonlinear $5$D dynamics can be studied numerically, and the algorithm and code we describe are the first ones of that type designed for this task. The program and its full documentation are available on the Web at \\texttt{http://www.cita.utoronto.ca/{\\textasciitilde}jmartin/BRANECODE/}~\\footnote{We also maintain a mirror of the BRANECODE website at \\texttt{http://www.cita.utoronto.ca/{\\textasciitilde}kofman/BRANECODE/}}. In this paper we provide a brief overview of what the program does and how to use it. ",
        "introduction": "Many extensions of the Standard Model have in common the presence of extra dimensions. This has to be contrasted with the fact that our world looks four dimensional, so one has to explain why the presence of the extra space has not yet been detected. The traditional answer has been that the extra space is compact and very small, so that the fields associated with its excitations are too heavy to be observable in accelerators or cosmology. More recently, it has been realized that ordinary matter and gauge interactions may be confined on lower dimensional submanifolds, known as branes. In this case, they could be four dimensional objects, even if the geometry of the theory is higher dimensional. The situation is different for gravity, which propagates in the whole bulk space. Several questions naturally arise, such as why a compact space would remain small while the three non-compact dimensions are undergoing cosmological expansion, or why the expansion of the universe we see is described by $3+1$ dimensional general relativity so well. The presence of extra dimensions may cause  deviations from the standard FRW cosmology that is supported by observations.   In most cases, these two questions turn out to be intimately related. Only if the extra space is static can the evolution of the non-compact coordinates behave as in the standard four dimensional case. Hence, the dynamics of the hidden dimensions becomes a crucial ingredient in understanding the evolution of the ones we observe. In some particular cases, static bulk configurations can be achieved under the combined action of the bulk/brane gravity. In most realistic examples that could account for our observed four  dimensional cosmology the stability is due to the presence of additional fields that acquire nontrivial configurations in the bulk. While the stabilization has to be effective at relatively ``late'' times, the first stages of our universe (before primordial nucleosynthesis, for instance) are much less constrained.  The evolution of the bulk may have been significant at this phase, and this offers many new possibilities for phenomenology. This is particularly true with the addition of the fields responsible for the ``late'' time stabilization, since they constitute new dynamical degrees of freedom for the system.  While the above considerations are valid for all models with extra dimensions, significant computations have been performed in the framework of brane models. These models can be thought of as simplified, phenomenological (bottom-up) versions of branes in string theory. The string dynamics is ignored and the primary focus is on the classical dynamics from the point of view of General Relativity. Branes act as a source for the Einstein equations of the system, with their tension and possibly with the energy density of fields confined on them.  Additional sources are a bulk cosmological constant, or the energy density of possible bulk fields. This set-up is sufficiently rich to describe very interesting situations. For example, inflation in braneworlds can acquire a nice geometrical interpretation, with the inflaton associated with the distance between different branes, the Hubble parameter scale associated with the induced curvature on the branes, and with reheating through radion oscillations.  Despite this great simplification, the whole dynamics is still very complicated, in particular when a bulk scalar field is present. In this case, analytical computations are typically performed within an effective $4$ dimensional theory, obtained after the extra dimension is integrated out, or perturbatively, using linearized analysis around simple backgrounds. While these studies  are very useful when the extra space is static or quasi-static, they are not sufficient  to describe the system when the evolution of the bulk is important. For example, systems that are stable at low energy (low curvature of the branes) can become unstable when the energy/curvature is increased. The stability/instability can be studied analytically. However, one cannot determine analytically where the system will evolve towards when the initial configuration turns out to be unstable.  For example, we numerically examined the dynamics of brane collisions and found that, as the branes approach each other, the spacetime of the bulk asymptotically  approaches the Kasner-type solution.  Motivated by these limitations with the analytical treatment, we undertook a numerical study of these models. We developed a numerical algorithm, implemented in a C++ code, specifically designed for codimension one braneworlds, with a scalar field included in order to provide stabilization of the bulk at low energies. With the assumption of homogeneity and isotropy along the brane spatial coordinates ${\\bf x}\\,$ (corresponding to the standard assumption of homogeneity and isotropy of the non-compact coordinates), the problem is reduced to an effectively $2$ dimensional one. The independent variables are the bulk time $t$ and the bulk dimension $y$.  The program integrates numerically the full set of Einstein equations in the bulk, together with the Israel junction conditions at two orbifold branes. We discretize the two dimensional spacetime (time and the bulk coordinate) and solve the bulk differential equations by finite-differencing them using the so-called {\\it leapfrog} scheme. This algorithm is sufficient for our problem, and it provides a good compromise between accuracy and computational time (see Section \\ref{overview} for more details). The two branes act as (one dimensional) boundaries of this space, and the junction conditions provide the boundary conditions for the system at each time-step. The solution of a boundary value problem is required to provide generic initial conditions that fulfill the constraint equations and the boundary conditions at the beginning of the evolution. In the static configurations we have mentioned above, this boundary value problem is significantly simpler than in the general case. The setup of initial conditions is explained in more detail in Section \\ref{sec:init}.  The first results obtained with the code have been presented in~\\cite{Martin:2003yh}.  We are now making the code public on the World Wide Web under the name BRANECODE. Its website is at \\texttt{http://www.cita.utoronto.ca/{\\textasciitilde}kofman/BRANECODE/}. The website for the program has documentation, including derivations of all the equations used in the program. Here we present a short summary of what the program does and what it can be used for. For more details see the website. Section \\ref{overview} of this paper gives an overview of what the program is and how it works.  Sections \\ref{equations} and \\ref{sec:init} describe the evolution equations and the setting of initial conditions respectively. Section \\ref{output} describes some of the output generated by the program. The references section is limited to papers from our group related to the BRANECODE design and its first results~\\cite{Martin:2003yh, Frolov:2003yi, Kofman:2004tk, Contaldi:2004hr, Felder:2001da}. See these papers for a more complete set of references. ",
        "conclusions": "The main motivation for our work was to extend the knowledge of braneworld dynamics beyond the few situations where it was known analytically. Apart from these situations (characterized by a static or slowly evolving bulk), approximate methods based on effective $4$ dimensional computations are unreliable. In this short note we have presented an algorithm, together with its C++ implementation, designed for numerical computations in this framework. First results obtained from this code were presented in~\\cite{Martin:2003yh,Kofman:2004tk}. We could show that some bulk configurations which are stable at low energy (low value of the expansion rate $H$ of the two branes) become unstable as $H$ increases, in agreement with the analytical calculations of~\\cite{Frolov:2003yi}.  This can be interpreted as a part of a more general phenomenon of gravitational instability of compactification to four dimensional de Sitter geometry~\\cite{Contaldi:2004hr}.  The numerical integration allowed us to follow the evolution of the system starting from the unstable configuration. For certain bulk/brane potentials the system may evolve towards another static, but stable configuration, characterized by a lower value of $H$. The transition is typically a process of quick bulk reconfiguration.  In many cases, however, the second configuration does not exist or cannot be reached. The two branes then either move apart to infinite distance, or they collide.  Brane collisions is a very interesting subject by itself. The study of the numerical evolution led us to the conclusion that the asymptotic geometry is given by the universal Kasner-type that is typical of homogeneous but anisotropic strong gravity regimes. We did not investigated the case of branes departing from each other when one of them in general approaches a naked singularity in the bulk configuration.  We also did not studied the possibility of the formation of an apparent horizon between the branes.  Our algorithm is focused on the simplest possible set-up which allows for brane stabilization based on a generalized Goldberger-Wise mechanism. Only one bulk scalar field has been considered, although with arbitrary potentials in the bulk and at the two branes. The inclusion of more scalar fields would be straightforward. In particular, one could consider other fields, which are confined to the branes, and which are coupled to the bulk fields e.g. through the brane potentials. (The interplay between bulk/brane fields may lead to novel interesting features not considered in~\\cite{Martin:2003yh}). Another easy generalization would be the inclusion of perfect fluids on the branes (for example, describing standard matter and radiation). Less trivial but more interesting extensions could be the inclusion of other types of fields (form fields in the bulk, for instance), or evolution with more dimensions included. For example, relaxing the hypothesis of homogeneity and isotropy along the ordinary dimensions would allow the study of inhomogeneous perturbations of the system."
    },
    "0404/astro-ph0404526_arXiv.txt": {
        "abstract": "Light rays passing very close to a black hole may experience very strong deviations. Two geometries were separately considered in the recent literature: a source behind the black hole (standard gravitational lensing); a source in front of the black hole (retro-lensing). In this paper we start from the Strong Field Limit approach to recover both situations under the same formalism, describing not only the two geometries just mentioned but also any intermediate possible configurations of the system source-lens-observer, without any small-angle limitations. This is done for any spherically symmetric black holes and for the equatorial plane of Kerr black holes. In the light of this formalism we revisit the previous literature on retro-lensing, sensibly improving the observational estimates. In particular, for the case of the star S2, we give sharp predictions for the magnitude of the relativistic images and the time of their highest brightness, which should occur at the beginning of year 2018. The observation of such images would open fascinating perspectives on the measure of the physical parameters of the central black hole, including mass, spin and distance. ",
        "introduction": "The fact that light rays can wind an arbitrary number of times around a black hole before emerging back towards spatial infinity is well known since early times of General Relativity \\citep{Dar,Atk,Lum,Oha,Nem} (see e.g. the treatment in \\citet{Cha}). In practice, a source behind a black hole not only produces the two classical weak field gravitational lensing images, but also an infinite number of strong field images corresponding to photons with winding numbers running from 1 to infinity. This phenomenon has been revived in a work by \\citet{VirEll}, who showed that the supermassive black hole at the center of our Galaxy may be a suitable lens candidate. The resolution needed for such observation is very high but should be reached by next future Very Long Baseline Interferometry experiments, such as ARISE\\footnote{ARISE web page: arise.jpl.nasa.gov} and MAXIM\\footnote{MAXIM web page: maxim.gsfc.nasa.gov}. In order to describe gravitational lensing in such extreme cases, we cannot use any weak field approximation. However, it is possible to take advantage of the opposite limit and finally get a very simple and efficient analytical approximation, called the Strong Field Limit (SFL) \\citep{Boz1}. This method, firstly emerged in Schwarzschild black hole studies \\citep{Dar,Oha,BCIS}, was then applied to Reissner-Nordstrom black holes \\citep{ERT}, to charged black holes of heterotic string theory \\citep{Bha} and finally generalized to an arbitrary spherically symmetric metric \\citep{Boz1}. The microlensing situation was considered by \\citet{Pet}, while the Kerr black hole was explored analytically for quasi-equatorial motion \\citep{Boz2} and numerically for arbitrary motion \\citep{VazEst}. It is also interesting that a time delay measurement would give a precise estimate of the distance of the lens \\citep{BozMan}. For former and/or alternative formulations of strong field gravitational lensing see \\citet{FriNew,FKN,DabSch,Per}.  Already in the first studies of gravitational lensing in strong fields by \\citet{Dar}, it was noticed that the whole sky is mapped in the vicinity of what it was then called a \"compact sun\". In particular, it was known that a source in front of the lens could yield relativistic images and Einstein rings as bright as a source behind the lens (see e.g. \\citet{Lum}). In more recent years retro-lensing was re-discovered quite independently from standard gravitational lensing. \\citet{HolWhe} proposed that a black hole of a few solar masses passing within 1 pc from the solar system would redirect photons from the Sun backwards to the Earth. An observer would thus see a \"star\" lighting up and then switching off in the sky as the Earth in its motion enters and leaves the best alignment position. \\citet{DepS2} suggested the black hole in Sgr A* at the center of our Galaxy as a suitable retro-lens, proposing the star S2 (the star with the smallest average distance from Sgr A* discovered so far) as a candidate source. At the moment, this seems to be the best known candidate for gravitational lensing in strong fields, deserving a closest investigation. At the same time, \\citet{EirTor} considered retro-lensing by Sgr A*, using the analytical framework of the SFL method to give correct estimates for such a phenomenon; their investigation, however, is limited to small angles and cannot cover the case of S2. Finally, \\citet{DepRM} extended Holz \\& Wheeler proposal to Kerr retro-lensing, using the SFL method to calculate the position of the images but using the Holz \\& Wheeler formula for the amplification (which is inadequate for spinning black holes), finding no significant deviation from the light curves of the Schwarzschild case.  In this paper, we give up any limitation due to small-angle approximation and treat the standard and retro-gravitational lensing on a unified ground, where they just come up as particular cases. Besides recovering these two situations, we also address the gravitational lensing problem for any intermediate geometries, where the strong field images are still present as always, but they are usually dimmer (see Sect. 2). This allows us to treat in a very accurate way the case of the star S2, giving sharp predictions for the light curves of its images in the next years. Moreover, we clarify several aspects which were not clearly stated in the former literature, constructing a unique analytical framework for the whole phenomenology.  In Sect. 2 we treat the spherically symmetric lens; we give the formulae for the position and the magnification of the images, and discuss the differences with the magnification formula by \\citet{HolWhe}; we also comment on the importance of time delay measurements. Sect. 3 is devoted to the study of S2, the best candidate source for gravitational lensing in the SFL; we use our formalism to draw analytical curves of the images, accurately predicting the epoch of their luminosity peak. In Sect. 4, in the light of the results by \\citet{Boz2}, we guess about the possible changes in S2 relativistic images if the black hole at the center of our galaxy has non-vanishing spin; we also compare our results in Sun retro-lensing by Kerr black holes with those by \\citet{DepRM}. Sect. 5 contains a summary of the work. ",
        "conclusions": "Gravitational lensing in strong fields is a very interesting subject from the theoretical point of view. Potentially it is a very appealing phenomenon which is completely embedded within the full general relativity. It would thus be an exceptional probe for physics in the regions close to the event horizons of black holes and would give very important feedback on the correct theory of gravitation. This justifies the very strong efforts that have been done by several groups to find possible candidate sources and lenses which could make this fascinating phenomenon manifest. The astrophysical cases investigated up to now are such that the relativistic images should be \"almost\" observable or should become observable within not many years. The most important problem is whether these theoretical configurations are likely or not. The case of S2, on the contrary, is a concrete case where definite predictions for the image luminosities can be done.  In this work we have extended the analytical framework of the Strong Field Limit analysis to cover all the possible geometric configurations, giving up any small angle approximations which are not automatically encoded in the tracks of the light rays. In this way, we have been capable to revisit the so called retro-lensing, in particular the Sun retro-lensing proposed by \\citet{HolWhe}, with more valid mathematical instruments. Moreover, we can adequately cover any intermediate case, such that of the star S2 \\citep{DepS2}. For the first relativistic image of this star, we predict an epoch of maximal luminosity (better than 32 magnitudes) between the end of year 2017 and the first part of 2018.  Using previous work on equatorial Kerr lensing \\citep{Boz2}, we have guessed how the presence of extended caustics and their drift from the optical axis can affect the brightness of the relativistic images of S2 at the epoch of its crossing through the galactic plane. We have also put upper limits on the black hole spin in Kerr retro-lensing of the Sun in order to have still significant images.  As a final remark, we can say that strong field gravitational lensing opens really fascinating perspectives, like testing the General Relativity, measuring the distance, the spin or other physical parameters of the black hole in the Galactic center and in the centers of other galaxies. The nice luminosity estimates for S2 encourage to look for more stars orbiting very close to the central black hole as new potential candidate sources for such an amazing phenomenon."
    },
    "0404/astro-ph0404183_arXiv.txt": {
        "abstract": "Analysing mid--infrared ISOCAM  images of the cluster of galaxies  J1888.16CL, we  identified among its members several particularly active galaxies  with total infrared luminosities well above  $10^{11}~\\Lo$. If powered by dust enshrouded starbursts, as suggested by their optical spectra, these Luminous Infrared Galaxies would exhibit  star  formation rates surprisingly high in the cluster environment. The triggering mechanism  is unclear but could be tidal collisions within sub-structures or infalling groups. ",
        "introduction": "In the Universe, the galaxies  forming stars with the highest rates and/or exhibiting the most extreme nuclear radiation turn out to be the most obscured ones. Their  internal activity  being partly  dust-enshrouded, it remains underestimated in visible observing bands and  shows up in the infrared regime.  Such so-called Luminous Infrared Galaxies (LIRGs) seem  to be more prevalent  in relatively dense environments like groups where tidal galaxy interactions  play a role in enhancing their activity. On the other hand, the densest regions like clusters are usually considered as being particularly unfavorable for the global activity of their host galaxies. The combined action of  multiple high-speed collisions, of the global cluster tidal field, and of the ram-pressure exerted by  the rich intra-cluster medium  contribute  to strip galaxies  of their disk and halo gas reservoirs. Without this fuel, star formation is  quenched and eventually stops while  the galaxy center is no longer fed and  any  nuclear activity ends.  \\vspace{0.3cm}  We present here   data obtained with the ISOCAM camera on board the Infrared Space Observatory (ISO) that reveals  the presence of confirmed cluster members with  high infrared luminosities and inferred star formation rates of  several tens of solar masses per year.  These LIRGs were discovered in the X-ray luminous cluster J1888.16CL, situated at a redshift of 0.56. The observations performed in the LW3 band centered at 15 $\\mu$m covered a strip of $3 \\times 15$ square arcminutes. Follow-up ground based observations were carried out using ESO facilities in Chile. They consist of optical B,V,R WFI/2.2m  images, near-infrared J,H,Ks SOFI/NTT  images and  optical FORS1/VLT multi-slit spectra. This collection of data  allowed us to assess the cluster membership of the ISO sources and study their spectral properties.  \\begin{figure} \\centerline{\\resizebox{\\textwidth}{!}{\\includegraphics{duc.fig1.ps}}} \\caption{Mid-infrared emitters towards the rich cluster of galaxies  J1888.16CL. The 15 $\\mu m$  emission is shown with   contours corresponding to different signal to noise ratio levels starting with a SNR of 3. In this figure are  presented the six ISOCAM emitters firmly identified with  galaxies having a measured spectroscopic redshift consistent with  cluster membership. The two bright sources at the bottom may lie  in groups falling towards the cluster. } \\end{figure} ",
        "conclusions": "Although the effective triggering mechanism remains unclear  --perhaps tidal interactions in sub-structures --, our observations clearly indicate that  a significant activity may occur in  clusters, even at moderate redshifts and hence a long time after their collapse and the formation of the bulk of their stellar populations. This result is corroborated with other studies based on ISOCAM mid--infrared data  (Ricardo-Perez et al., in these proceedings), optical data (see the review by A. Dressler), radio continuum maps and perhaps millimetric SCUBA surveys   (Webb et al., these proceedings; \\cite[Best, 2002]{Best02}). The MIPS camera on board  the infrared satellite Sptizer will also very soon provide ample information on the dust  enshrouded activity in clusters. This secondary peak of activity may correspond to a phase in the evolution of clusters when accretion of galaxies, especially through the merging of groups, was conspicuous."
    },
    "0404/astro-ph0404306_arXiv.txt": {
        "abstract": "{ Since a large fraction of active galactic nuclei (AGN) is missed in common UV-excess surveys and is even hard to find in radio, near-IR and X-ray surveys, we have used a new AGN selection technique which is expected to be not affected by extinction. Within the scientific verification of the {\\it ISOCAM Parallel Survey} at $6.7\\,\\mu$m we have discovered objects with exceptional mid-infrared (MIR) emission. They are essentially not detected on IRAS-ADDSCANs and only very few of them show up in the NVSS and FIRST radio surveys. Various colour criteria of the $6.7\\,\\mu$m data with 2MASS and optical wavebands show that the sources reach more extreme IR colours than the sources in the Hubble Deep Field-South and the ELAIS survey. The comparison with known object types suggests that we have found AGN with a pronounced MIR emission, probably due to circum-nuclear dust. First results from optical spectroscopy of ten candidates corroborate this interpretation showing four AGN, two reddened LINER and four extremely reddened emission-line galaxies with MIR/FIR flux ratios higher than for known pure starburst galaxies. The results will make a significant contribution to the debate on the entire AGN population. ",
        "introduction": "\\label{section_introduction}  The realisation that obscuration plays a critical role in the classification of AGN fundamentally inspired the current research. Attempts to overcome the limitations of dust extinction and to identify the entire AGN population -- including type 2 and dust-enshrouded AGN -- encompass surveys in the near-IR, radio, and X-ray regimes. However, searching for very red AGN the colour selection via $J - K_{\\rm s} > 2$  (Cutri et al. 2001) excludes most of the known AGN (Barkhouse and Hall 2001), only about 30\\% of AGN are radio-loud (Urry \\& Padovani 1995), and there seems to exist many X-ray faint AGN (Wilkes et al. 2002). Thus, a considerable fraction of the AGN population must have escaped detection due to observational bias. Webster et al. (1995) found that their radio-selected quasar sample is significantly redder than an optical comparison sample and concluded that up to 80\\% of the quasars have been missed in conventional optical surveys, provided that the redder colours of the radio-loud quasars are due to dust reddening and that the radio-quiet quasars contain as much dust as the radio-loud ones. By new strategies in the optical, assuming that the narrow-line regions are sufficiently extended and that only the continuum emission is hidden, type-2 quasar candidates have been selected as objects with narrow permitted emission lines and high [\\ion{O}{III}]$\\lambda$5007 equivalent widths (Djorgovski et al. 2001, Zakamska et al. 2003). Applying a moderate colour cut $J-K_{\\rm s} > 1.2$, from about 1500 sources in 2MASS Francis et al. (2004) find only tentative evidence that Seyfert\\,2 nuclei are more common in the NIR selected survey than in blue selected galaxy surveys, and they can place only very weak constraints on any population of dusty AGN.  The disadvantage of heavy extinction in optical and NIR surveys can turn into a valuable detection tool, when observing dust-surrounded AGN at MIR wavelengths. There, the reemission of the hiding dust heated by the strong radiation field of the AGN should be seen easily as MIR excess. In fact, for known (powerful) AGN of both type 1 and type 2 a steep near- to mid-IR slope has been revealed by sensitive MIR observations (e.g. Haas et al. 2003,  Haas et al. 2004, Siebenmorgen et al. 2004).  We therefore started a new approach, searching for AGN by means of their MIR emission of the nuclear dust torus. However, one complication with this method has to be solved: Since luminous IR starburst galaxies may also show a pronounced MIR emission due to the PAH emission bands around $7.7\\,\\mu$m, it is of special importance to distinguish them from AGN. In this Letter we describe the new technique for the mid-IR selection of AGN candidates using IR colour diagrams and report about first results from optical spectroscopy.  \\vspace{-0.2cm} ",
        "conclusions": "We have discovered a sample of unique MIR excess sources. Various arguments suggest that they likely contain an AGN. The ongoing optical spectroscopy indicates 40\\% of the sample to be classical AGN and the remaining part to be dust-enshrouded sources. They appear enigmatic in terms of having a high MIR/NIR flux ratio, but relatively low FIR upper limits.  Since a moderate FIR/MIR flux ratio is more typical for AGN, it  argues against pure starbursts. This conjecture has to be confirmed. Because of the high extinction, optical observations might not be able to see the true AGN or starburst or composite nature of those sources. Therefore, future observations with XMM-Newton, the VLA and IR spectroscopy with the Spitzer Space Telescope are planned.  If all 80 ISOCP MIR excess sources turn out to be AGN and if 32 of them (= 40\\%) are classical AGN which could have been identified also in optical surveys, then the number counts of AGN may increase by a factor of up to 80/32 = 2.5 due to the inclusion of dusty AGN. Clearly, a reliable comparison will have to wait until spectra of more ISOCP sources are obtained.  The emphasis of this letter is on presenting a promising strategy how to find the presumed missing AGN. The selection of MIR excess sources via the $H-K_{\\rm s}$ and $K_{\\rm s}-LW2$ colour-colour diagrams in combination with the low 60$\\mu$m upper flux limits turns out to be efficient in separating starbursts from AGN. It will establish a long sought-after technique for selecting the full population of AGN largely free  of extinction."
    },
    "0404/astro-ph0404076_arXiv.txt": {
        "abstract": "We have calculated the opacities and secondary production mechanisms of high energy photons arising in gamma-ray burst internal shocks, using exact cross-sections for the relevant processes. We find that for reasonable choices of parameters, photons in the range of 10's to 100's of GeV may be emitted in the prompt phase. Photons above this range are subject to electron-positron pair production with fireball photons and would be absent from the spectrum escaping the gamma-ray burst. We find that, in such cases, the fireball becomes optically thin again at ultra-high energies ($\\gtrsim$ PeV). On the other hand, for sufficiently large fireball bulk Lorentz factors, the fireball is optically thin at all energies. Both for $\\gamma\\gamma$ self-absorbed and optically thin cases, the escaping high energy photons can interact with infra-red and microwave background photons to produce delayed secondary photons in the GeV-TeV range.  These may be observable with GLAST, or at low redshifts with ground-based air Cherenkov telescopes. Detection of the primary prompt spectrum constrains the bulk Lorentz factor, while detection of delayed secondary gamma-rays would provide a consistency check for the primary spectrum and the bulk Lorentz factor as well as constraints on the intergalactic magnetic field strength. ",
        "introduction": "GeV photons have been detected from a number of gamma-ray bursts (GRB) \\citep{dingus03} with the EGRET detector on the Compton Gamma Ray Observatory. Photons at energies in this range and higher can make pairs via $\\gamma\\gamma \\to {\\rm e}^+ {\\rm e}^-$ interactions in the GRB fireball, the optical depth being dependent on the bulk Lorentz factor $\\Gamma_b$, so spectral observations at this energy can provide useful constraints \\citep{baring94, lithwick01} on this key ingredient of GRB models.  In addition, high energy photons which escape $\\gamma\\gamma$ absorption in the source can still suffer such interactions far from the source but before reaching the observer, against cosmic infrared background (CIB) and/or cosmic microwave background (CMB) photons. The electron-positron pairs thus produced can inverse-Compton scatter again on the most numerous CMB photons, giving rise to a delayed secondary spectrum \\citep{plaga95, dai02a}, which ranges typically up to 100's of GeV.  In this paper we investigate, using the exact cross sections, the high energy interactions of photons produced in a simple GRB internal shock model, as a function of the luminosity, the typical variability timescale and the bulk Lorentz factor. We approximate the original GRB spectrum through an observationally motivated Band function \\citep{band93}, and calculate the emerging spectrum as a function of the burst parameters. This spectrum is then subjected to interactions with the CIB and CMB, and the shape of the secondary spectrum and its time delay is calculated assuming typical cosmological distances $z\\sim 1$ and various values for the intergalactic magnetic field strength. The pair-production cut-off of the primary spectrum and the delayed secondary spectrum fall, for typical GRB parameters, in the GLAST \\citep{gehrels99} energy range, and for a range of intergalactic magnetic field strengths the latter should be detectable by GLAST as well as by the newer generation of air-Cherenkov telescopes \\citep{weekes01}. ",
        "conclusions": "We have calculated the spectrum and time dependence of the high energy ($\\gtrsim$ GeV) spectrum emerging from GRB internal shocks as a function of the luminosity, time variability and the bulk Lorentz factor, using exact cross sections.  For typical bursts parameters the internal shock spectrum cuts off above a threshold energy $E_{\\gamma,{\\rm th}}\\sim 10-100$ GeV due to $\\gamma\\gamma$ iteractions within the shock region.  There is a qualitative difference with the results of \\citet{baring94}, who discussed the $\\gamma\\gamma$ absorption in external shocks, and who did not consider the secondary spectra from interactions with background radiation. More recently \\citet{wang04} also discussed the $\\gamma\\gamma$ absorption in external shocks, including interactions against background photons. The main difference with our work is due to the much lower $\\gamma\\gamma$ optical depth from the lower photon density at the radii of external shocks, which are much larger than the radii of internal shocks considered here.  Internal shock $\\gamma\\gamma$ absorption using approximate cross sections was discussed by \\citet{pilla98} and \\citet{lithwick01} without interactions against external background photons, and by \\citet{dai02a} including such interactions.  These authors did not consider the effects of synchrotron self-absorption.  The main qualitative difference between the emerging internal shock spectra of these authors and our work arises from our use of accurate cross sections, and more importantly, our inclusion of synchrotron self-absorption.  The cut-offs in the emerging spectra that we obtain are compatible with more detailed numerical calculations of the self-consistent pair creation and annihilation spectrum in internal shocks of \\citet{peer03}, without inclusion of interactions against external background photons. The main difference between their final spectra and ours is due to their using an ab-initio spectrum instead of a phenomenological Band spectrum as here, and to the additional effects of our inclusion here of interactions againts external photons, leading to a delayed secondary spectrum.  A new feature discussed in this paper is that the spectra that emerge from the GRB internal shock region shows, besides the expected high energy cutoff due to pair-production $E_{\\gamma,{\\rm th}}$, also a re-emergence of the spectrum at a higher energy $E_{\\gamma,{\\rm thin}} > E_{\\gamma,{\\rm th}}$ where the shock becomes optically thin again to $\\gamma\\gamma$ interactions.  This is due to our inclusion of synchrotron self-absorption, which reduces the internal photon target spectral density causing absorption of the highest energy emerging photons. For high enough bulk Lorentz factors (typically $\\Gamma_b \\gtrsim 800$), the fireball is optically thin to internal $\\gamma\\gamma$ interactions at all energies, and there is no internal $\\gamma\\gamma$ cut-off band in the emerging spectrum.  The emerging GRB spectrum is modified, on its way to the observer, by interactions with the diffuse infrared background (CIB) and the diffuse microwave background (CMB). Here, we have used the GRB internal shock emerging spectra including synchrotron self-absorption as the source of input photons, and calculate their interaction with a CIB spectrum redshifted to $z\\sim 1$, based on that of \\citet{malkan01}.  We then calculated the secondary GRB spectrum from the resulting secondary pairs which are up-scattered on CMB photons. We used two different escaping GRB primary spectra, one with an internal $\\gamma\\gamma$ cutoff band (for a typical burst bulk Lorentz factor $\\Gamma_b\\sim 300$) and one without an internal $\\gamma\\gamma$ cutoff band (for a burst with a larger $\\Gamma_b \\sim 800$). We used two different isotropic-equivalent luminosities, and different intergalactic magnetic field strengths ranging between $B_{\\rm IG}= 10^{-20}$ G and $10^{-17}$ G. This leads to a secondary spectrum, mainly in the 1-100 GeV range, in addition to and delayed respect to the prompt unabsorbed internal shock spectrum.  We also calculated the further interaction of the secondary gamma-ray spectrum with the CIB, which gives a similarly attenuated spectral shape as that of the primary spectrum. We did not consider the delayed spectrum of the tertiary and higher order pairs, since the delay time scales are longer and the fluxes of these components is much lower.  A different delayed GeV component can arise in the afterglow phase, when the inverse Compton spectrum peak energy in the external shock region sweeps across the GeV band \\citep{mr94, dermer00, se01, zm01}. This delayed GeV emission, from IC upscattering on external shock electrons, is predicted to be detectable by GLAST at $z \\sim 1$ \\citep{zm01}, as long as the shock parameters are such that the IC component is prominent (which is the case for the typical parameters inferred from broadband modeling of some well-studied bursts \\citep{wijers99, panaitescu02, yost03}).  The IC duration can be up to a few hours, which could overlap with the earlier part of the delayed spectra discussed here.  The temporal evolution of the IC spectral component is such that it hardens with time initially, but softens with time after the IC peak crosses the band. The delayed GeV component discussed in this paper has a rather different temporal evolution behavior. Since $\\Delta t (\\gamma_e)$ is essentially anti-correlated with $\\gamma_e$ (see Figure 3), the hard spectrum (due to the more energetic electrons) lasts a shorter timescale than the soft spectrum. Also the harder photons arrive earlier than the softer photons. Throughout the whole delayed phase, the hardest photons emerge first, followed by progressively softer photons. Later the hardest photons drop out first from the spectrum.  The delayed spectrum always progressively softens with time.  This distinct spectral evolution behavior can be used to distinguish the delayed emission from IC scattering in external shocks from the delayed emission of secondary spectra due to pair-production and IC scattering on the diffuse background discussed here. The secondary external spectrum from such an IC spectrum, however, should be similar to what we have discussed here, on timescales longer than a few hours, since the absorbed energy undergoes similar reprocessing.  The time delays and the flux level of the secondary spectra are sensitive to the intergalactic magnetic field strength. The best prospects for detection of such secondary delayed GeV spectra occurs for bursts with bulk Lorentz factors $\\gtrsim 800$ and intergalactic magnetic fields $\\lesssim 10^{-17}$ G, which might be typical of intergalactic void regions. Since the mean-free-path of the first generation high energy photons against $\\gamma\\gamma$ absorption typically exceeds several to tens of Mpc \\citep{coppi97, malkan01}, it is plausible that a substantial fraction of the pair formation will occur in void regions near the burst, even if the latter occured in a cluster galaxy.  The corresponding time delays are in the range 10-100 s and the fluxes are within the sensitivity range of GLAST, and at the high energy end, also of VERITAS, HEGRA, and other ACTs.  A detection of this secondary delayed GeV spectrum in GRB would provide a valuable diagnostic for the IGM magnetic field strength, as well as the bulk Lorentz factor, and would provide valuable clues for the typical size of the shocks responsible for the radiation."
    },
    "0404/astro-ph0404595_arXiv.txt": {
        "abstract": " ",
        "introduction": "Gamma-rays of TeV energies from distant BL Lacertae objects, or `blazars', provide us a problem so-called `TeV gamma crisis' \\cite{P&M}. For example, HEGRA collaboration observed the gamma-ray from Markarian 501 (Mrk501) in 139 Mpc distance, and the maximum measured energy of the gamma-ray reaches to about 20 TeV \\cite{M501}. An energetic gamma-ray can be attenuated by scatterings with the extra-galactic infrared photons in $\\gamma\\gamma \\rightarrow e^+e^-$ process. Since the evaluated mean free-path of a 20 TeV photon is only 13 Mpc, the gamma-rays should have been suppressed by a factor of $\\exp(-139/13)$ = $1/(4 \\times 10^4)$. It means that the original brightness of Mrk501 in 20 TeV is $4 \\times 10^4$ times brighter than what HEGRA saw. Comparing the observed energy spectrum of the gamma-rays of Mrk501 with that of much closer TeV gamma-ray sources like the Crab nebula, such a strong suppression in high energy side in the energy spectrum seems to be unrealistic. The same problem also holds for Mrk421, 1ES1959+650 and H1426+428 \\cite{four}. The observed highest energies of the photons from these blazars are listed in \\ref{T1}.  Some trials to solve the `crisis' have been proposed by many authors. The most conservative way is to assume a less density of the extra-galactic infrared photons. One can also evaluate the appropriate IR photon density so that the traveling gamma-rays do not absorbed so much \\cite{four}. There is another possibility to assume the violation of the Lorentz invariance in the space-time \\cite{many,P&M}. Such kind of idea is motivated by a recent development of the theories of the large extra-dimension. In the deformed space-time, it is expected that the energy threshold of the $\\gamma\\gamma \\rightarrow e^+e^-$ process can be shifted to reduce the absorption of the TeV gamma-rays.  In this paper, we examine the possibility to answer the `TeV gamma crisis' by a model with reduced degree of space freedom at a high energy scale \\cite{Suga}. In a model whose one spatial direction is discretized at a high energy scale, we evaluate the extra-galactic absorption of an energetic photon which acquires an effective mass in the model.  The paper is organized as follows. In the section \\ref{S2}, the adopted model is presented. The threshold shift of the scattering process and the mean free-path of the high energy photon in this model are evaluated in the section \\ref{S3}. The effects of the universe expansion on the photon propagation is discussed in the section \\ref{S4}. We give the conclusions in the last section. ",
        "conclusions": "\\label{S5}  To solve the `TeV gamma crisis', we have examined a model whose one spatial direction is discretized at a high energy scale. Assuming the standard extra-galactic IR photon distribution, we have evaluated the mean free-path of an energetic photon which acquires an effective mass in the model. For a wide range of the value of the lattice energy scale $a^{-1}$ between a few TeV and around $10^{10}$ GeV, the mean free-path of a TeV energy photon can be enlarged enough to solve the `crisis'. Taking into account the effects of the universe expansion, we have found bounds of $a^{-1}$: 2.5 TeV $< a^{-1}$ $\\simlt 10^{10.5}$ GeV for a 20 TeV photon from Mrk501, and 1.4 TeV $< a^{-1}$ $\\simlt 10^{9.5}$ GeV for a 5.37 TeV photon from H1426+428, under an assumption that the blazar gamma-ray travels to the discretized direction.  Further test of the present model can be performed by watching PeV photons. In the ordinary space-time, a 1 PeV photon has the mean free-path of 9 kpc. However, it can travel over 300 Mpc in our model with $a^{-1}$ $\\simlt 10^{10}$ GeV, as can be seen in \\ref{F1}. Observation of PeV photons from the near-by galaxies will support our model, which is a consequence of the item \\ref{c3} described in the end at the section \\ref{S3}. For $a^{-1}$ $\\sim 10^8$--$10^9$ GeV, a 1 TeV photon has shorter mean free-path than a few to several TeV photons in our model, as was also mentioned in the item \\ref{c1}. It results in a dent of the photon energy spectrum around 1 TeV for the objects which are distant as hundred Mpc.  We also mention some other possible effect of the space discretization. In the case of smaller values of $a^{-1}$ as a few TeV to several ten TeV, collider experiments may find some direct evidence that the energetic light particles fly slowly due to the effective mass. Modification to the QED vertex structure may affect to the magnetic and the electric dipole moments of the elementary charged particles, even in the interactions with low energy photons. These observations will bring us additional bounds of $a^{-1}$."
    },
    "0404/astro-ph0404289_arXiv.txt": {
        "abstract": "We report the results of a 45 ks \\cxo\\ observation of the cataclysmic variable V426 Ophiuchus. The high resolution spectrum from the high-energy transmission grating spectrometer is most consistent with a cooling flow model, placing \\targ\\ among the group of CVs including U Gem and EX Hya. An uninterrupted lightcurve was also constructed, in which we detect a significant 4.2~hr modulation together with its first harmonic at 2.1~hrs.  Reanalysis of archival \\ginga\\, and \\rosat\\ X-ray lightcurves also reveals modulations at periods consistent with 4.2 and/or 2.1~hrs. Furthermore, optical photometry in V, simultaneous with the \\cxo\\ observation, indicates a modulation anti-correlated with the X-ray, and later more extensive R band photometry finds a signal at \\til2.1~hrs.  The earlier reported X-ray periods at \\til0.5 and 1~hrs appear to be only transient and quasi-periodic in nature. In contrast, the 4.2~hr period or its harmonic are stable and persistent in X-ray/optical data from 1988 to 2003.  This periodicity is clearly distinct from the 6.85~hr orbit, and could be due to the spin of the white dwarf. If this is the case, V426 Oph would be the first long period intermediate polar with a ratio  $P_{spin}/P_{orb}$ of 0.6. However, this interpretation requires unreasonable values of magnetic field strength and mass accretion rate. ",
        "introduction": "The cataclysmic variable V426 Oph has proven to be somewhat enigmatic, even though it is a bright (V=11.5--13.4) system that has been studied for a number of years since its discovery as a nova-like with an emission line spectrum \\citep{herb60}. A detailed spectroscopic study by Hessman (1988) revealed a K3 dwarf secondary in a 6.85~hr orbit with a white dwarf, with an inclination of $53^\\circ$ and at a distance of 200 pc. He concluded this was a Z Cam type of dwarf nova with outbursts approximately every 22 days and some standstills in the light curve occurring about one magnitude fainter than the outburst brightness. This interpretation would indicate V426 Oph is a system with such a high rate of mass transfer that it is very close to the upper limit for dwarf novae outbursts \\citep{osak96}.  However, there are several observational clues that V426 Oph might also harbor a magnetic white dwarf. IUE observations show an unusually flat flux distribution in the UV, possibly indicating the truncation of a hot inner disk by a magnetic white dwarf (i.e. an intermediate polar), while the X-ray flux at quiescence is larger than for typical dwarf novae \\citep{szko86}. The emission lines show flaring activity and a phase shift from a location near the white dwarf \\citep{hess88}. Most importantly, quasi-periodic variability has been seen in the optical and X-ray at a period near 30 min at a brightness comparable to standstill \\citep{szko90} and near 1, 2.5 or 4.5~hrs in the X-ray at quiescent brightness (Szkody 1986, Rosen et al. 1994; but see Hellier et al. 1990). While none of these periods could be positively identified as that of the rotation of the white dwarf to confirm the identification as an intermediate polar, the different brightness states combined with the lack of a long stretch of continuous observation have hampered attempts to resolve the origins of the observed periodicities.  With the advent of {\\it Chandra}, high resolution X-ray spectroscopy can be combined with uninterrupted long light curves to obtain much better information on the accreting regions of cataclysmic variables. Results are now available for the dwarf novae U Gem \\citep{szko02}, and  WX Hyi \\citep{pern03}, the intermediate polar EX Hya \\citep{mauc01}, along with 5 other cataclysmic variables \\citep{muka03}. One of the surprising results is how similar the spectrum of the low mass accretion disk system U Gem is to that of the intermediate polar EX Hya. \\citet{muka03} have found that the seven systems can be divided into two groups. The first group, which includes U Gem and EX Hya, shows a bremsstrahlung continuum with strong H and He-like ion emission as well as ions from FeXVII-FeXXIV. The other group shows a harder X-ray continuum and little \\linebreak Fe L-shell emission. They suggest the differences in these two groups may be due to differences in the accretion rate per unit area. In the first group, the energy is released as optically thin radiation in a steady state cooling flow. In the second, the line emission arises from a plasma photoionized by the hard continuum. To further explore these ideas, and to try to resolve the nature of the accretion in V426 Oph, we obtained a {\\it Chandra} observation at quiescence, together with simultaneous optical photometry and a single optical spectrum. We also obtained additional optical photometry (in quiescence) at a later date. ",
        "conclusions": "Our analysis of a long uninterrupted \\cxo\\ X-ray lightcurve of \\targ\\ together with optical photometry and reanalysis of archival X-ray data, has revealed a persistent period at 4.2~hrs, unrelated to the known orbit.  In contrast, we find that previous candidate spin periods seen in past X-ray studies at \\til0.5 and 1~hrs are most likely non-persistent QPOs. Our simultaneous \\linebreak[4] X-ray/optical photometry shows that the modulations are anti-correlated and therefore the variations must arise from different parts of the system; possibly due to the reprocessing of beamed X-radiation.  Indeed, the period of 4.2~hrs is easiest to interpret as the rotation of the white dwarf, implying that \\targ\\ is, after all, an IP, and making it the first long period IP with a similar spin-orbit equilibrium to EX Hya, $P_{spin}/P_{orb}\\sim0.6$.  However, this interpretation leads to problems with the magnetic field strengths and mass accretion rates that are typical for intermediate polars.   The high-resolution spectrum  places \\targ\\ as a member of the group of CVs, including U Gem and EX Hya, where the X-rays are emitted from a cooling flow. Indeed, more detailed modeling of the line fluxes for such a shock model lends further support to the magnetic CV interpretation."
    },
    "0404/astro-ph0404240_arXiv.txt": {
        "abstract": "We address three questions regarding solar system planets. What determined their number?  Why are their orbits nearly circular and coplanar? How long did they take to form?  Runaway accretion in a disk of small bodies resulted in a tiny fraction of the bodies growing much larger than all the others. These big bodies dominated the viscous stirring of all bodies. Dynamical friction by small bodies cooled the random velocities of the big ones. Random velocities of small bodies were cooled by mutual collisions and/or gas drag.  Runaway accretion terminated when the orbital separations of the big bodies became as wide as their feeding zones. This was followed by oligarchic growth during which the big bodies maintained similar masses and uniformly spaced semi-major axes. As the oligarchs grew, their number density decreased, but their surface mass density increased. We depart from standard treatments of planet formation by assuming that as the big bodies got bigger, the small ones got smaller as the result of undergoing a collisional fragmentation cascade. It follows that oligarchy was a brief stage in solar system evolution.  When the oligarchs' surface mass density matched that of the small bodies, dynamical friction was no longer able to balance viscous stirring, so their velocity dispersion increased to the extent that their orbits crossed. This marked the end of oligarchy.  What happened next differed in the inner and outer parts of the planetary system. In the inner part, where the ratios of the escape velocities from the surfaces of the planets to the escape velocities from their orbits are smaller than unity, big bodies collided and coalesced after their random velocities became comparable to their escape velocities. In the outer part, where these ratios are larger than unity, the random velocities of some of the big bodies continued to rise until they were ejected. In both parts, the number density of the big bodies eventually decreased to the extent that gravitational interactions among them no longer produced large scale chaos. After that their orbital eccentricities and inclinations were damped by dynamical friction from the remaining small bodies.  The last and longest stage in planet formation was the clean up of small bodies.  Our understanding of this stage is fraught with uncertainty. The surviving protoplanets cleared wide gaps around their orbits that inhibited their ability to accrete small bodies. Nevertheless, in the inner planet system, all of the material in the small bodies ended up inside planets. Small bodies in the outer planet system probably could not have been accreted in the age of the solar system. A second generation of planetesimals may have formed in the disk of small bodies, either by collisional coagulation or gravitational instability.  In the outer solar system, bodies of kilometer size or larger would have had their random velocities excited until their orbits crossed those of neighboring protoplanets. Ultimately they would have either escaped from the sun or become residents of the Oort cloud. An important distinction is that growth of the inner planets continued through clean up, whereas assembly of the outer planets was essentially complete by the end of oligarchy. These conclusions imply that the surface density of the protoplanetary disk was that of the minimum solar mass nebula in the inner planet region but a few times larger in the outer planet region.  The timescale through clean up was set by the accretion rate at the geometrical cross section in the inner planet region, and by the ejection rate at the gravitationally enhanced cross section in the outer planet region. It was a few hundred million years in the former and a few billion years in the latter. However, since Uranus and Neptune acquired most of their mass by the end of oligarchy, they may have formed before Earth!  A few implications of the above scenario are worth noting. Impacts among protoplanets of comparable size were common in the inner planet system but not in the outer. Ejections from the outer planet system included several bodies with masses in excess of Earth after oligarchy, and an adequate number of kilometer size bodies to populate the Oort comet cloud during clean up. Except at the very end of clean up, collisions prevented Uranus and Neptune from ejecting kilometer size objects. Only Jupiter and, to a much lesser extent, Saturn were capable of populating the Oort cloud with comets of kilometer size. ",
        "introduction": "Modern scenarios for planet formation may be broken down into several stages.  The growth of the smallest gravitationally active bodies, planetesimals, is mired in controversy \\citep{Lis93,YS02}.  Orderly growth by the merging of planetesimals is followed by runaway accretion in which a small fraction of the bodies grow much larger than all the others \\citep{S72,WS89}.  When these big bodies are sparse enough so that each dominates viscous stirring in its feeding zone, runaway growth gives way to oligarchic growth during which the big bodies grow in lockstep maintaining similar masses and uniformly spaced orbits \\citep{KI98}.  As oligarchs grow, their orbital spacing increases and their number decreases. We investigate how oligarchy ends and what happens after it does.  The plan of our paper is as follows. We describe the conditions that pertain at the end of oligarchy in \\S \\ref{s:before}. We show in \\S \\ref{s:after} that at this stage dynamical friction from the small bodies is no longer able to balance the mutual stirring of the big bodies. \\S \\ref{s:completion} treats the regularization of the orbits of the big bodies and the clean up of small bodies.  We summarize our findings in \\S \\ref{ss:conclusions}.  A few definitions are in order. For simplicity, we consider two classes of bodies, big ones and small ones, each composed of material density $\\rho$.  We denote the surface mass density, random velocity dispersion, and radius of the former by $\\Sigma$, $v$, and $R$, and of the latter by $\\sigma$, $u$, and $s$. The distance from and angular velocity about the Sun are given by $a$ and $\\Omega$. In our numerical estimates, we set $\\rho=5.5\\gm\\cm^{-3}$, the density of Earth, at $1\\au$ and $\\rho=1.5\\gm\\cm^{-3}$, approximately the densities of Uranus and Neptune, at $25\\au$. For the condensible fraction of the protoplanetary nebula, we adopt the surface densities $\\sigma= 7\\gm\\cm^{-2}$ at $1\\au$ and $\\sigma=1.5\\gm\\cm^{-2}$ at $25\\au$. The former is just that appropriate to the minimum mass solar nebula \\citep{Hayashi81}, but the latter is enhanced sixfold relative to it. This enhancement is designed to make the isolation mass in the outer planet system comparable to the masses of Uranus and Neptune. This is necessary, since the timescale for the accumulation of the outer planets by coagulation of smaller isolation masses would exceed the age of the solar system. The particular value of six applies if half the mass had accreted into protoplanets by the end of oligarchy. Instead, if most of it had, three would be the appropriate enhancement factor.  In evaluating expressions containing the planet mass, $M_p$, we use $1 M_\\oplus$ for an inner planet and $15 M_\\oplus$ for an outer one, where $M_\\oplus\\approx 6.0\\times 10^{27}\\gm$ is Earth's mass.   It proves convenient to employ a symbol $\\alpha$ for the ratio between the radius of a body and that of its Hill sphere, $R_H$: \\begin{equation} \\alpha\\equiv \\left( {9\\over 4\\pi}{M_\\odot\\over \\rho a^3} \\right)^{1/3}\\approx \\cases{ 234^{-1} & for $a = 1\\au$ \\cr 3800^{-1} & for $a =25\\au$ } \\, . \\end{equation} Note that $\\alpha$ is approximately the angle subtended by the sun at $a$.  We also make use of the Hill velocity of the big bodies, $v_H\\sim \\Omega R_H$; $v_H\\sim \\alpha^{1/2}v_{\\rm esc}$, where $v_{\\rm esc}$ is the escape speed from the surface of a big body. ",
        "conclusions": "\\label{s:discussion}  \\subsection{Conclusions} \\label{ss:conclusions}  The scenario sketched in this letter addresses some of the basic problems in planet formation.  \\begin{itemize}  \\item The number and orbital spacing of the planets resulted from an evolution toward stability against large scale chaotic perturbations.  \\item After the cessation of chaotic perturbations, dynamical friction by the residual small bodies damped the orbital eccentricities and inclinations of the surviving protoplanets.  \\item Accretion during oligarchy involved small bodies created by a collisional fragmentation cascade. This stage probably lasted for less than $10^5\\yr$ in the inner solar system and less than $10^7\\yr$ in the outer solar system.  \\item The timescale for establishing the final configuration of planetary orbits was a few hundred million years for the inner planet system and a few billion years in the outer planet system. It was set by the accretion rate at the geometrical cross section in the former, and by the ejection rate at the gravitationally enhanced cross section in the latter.\\footnote{The timescale in the outer solar system could have been much shorter if all ejections were done by Jupiter and Saturn.}  \\item Clean up of small bodies is a complicated and poorly explored stage of planet formation. Small bodies in the inner solar system were incorporated into planets. Those in the outer solar system were probably ejected by Jupiter and Saturn, but that requires a second generation of planetesimal formation.  \\end{itemize}  \\subsection{Influence of Gas} \\label{ss:gasdrag}  We have neglected the influence of gas.  Observations of young stars indicate that protostellar disks dissipate in a few million years \\citep{HLL01,Strom93}.  We show below that although the presence of gas would alter some of our numerical results, it would not affect our picture qualitatively.  Gas drag can provide significant damping for the random velocities of small bodies in addition to that due to inelastic collisions. Relative to collisions, it is most effective in damping the random velocities of bodies that are smaller than the mean free path of the gas molecules.  For these, its damping rate obeys the same expression as that due to collisions, but with the surface density of the small bodies replaced by that of the gas.  We can account for the effects of gas drag by considering $s$ to be an effective size for the small bodies that can be less than their true size. This is a minor point for our story since, as we have emphasized, the true size of the small bodies is highly uncertain. Moreover, our main concern is with the stages of planet formation that follow velocity instability and these probably continue after the gas is gone.  \\cite{Raf03e} explored the fast accretion of protoplanetary cores in the presence of gas. His investigation runs parallel to the early phases of ours.  However, it terminates at the onset of velocity instability, when $v\\sim v_H$.  A potentially more significant effect of gas drag that was not considered by \\cite{Raf03e} is its role in damping the random velocities of the oligarchs. \\cite{Ward93} shows that the gas damping rate can be obtained from the damping rate due to small bodies by substituting the surface density of gas for that of the small bodies and the sound speed of the gas, $c_s$, for the random velocity of the small bodies, $u$.\\footnote{We note that at isolation $v_H<c_s<v_{\\rm es}$. Moreover we are assuming that $v<c_s$.} By stabilizing the oligarchs' random velocities, gas drag could have enabled them to consume all of the small bodies.  In the inner solar system, it is possible---though highly uncertain---that much of the gas survived until isolation. Then the full velocity instability of the oligarchs would have been delayed until the surface density of gas declined to match that contributed by oligarchs. After that the oligarchs would have excited their random velocities up to their escape speeds.  Although most of the small bodies would have been accreted before this happened, plenty of new ones created in glancing collisions could have damped the orbital eccentricities and inclinations of the planets that finally formed.\\footnote{Since it took the inner planets more than $100\\Myr$ to form (eqn. [\\ref{eq:tcoag}]), gas is unlikely to have contributed to regularizing their orbits.}  Outer solar system planets, Uranus and Neptune, are believed to have collected only a few earth masses of nebular gas. So it is likely that most of the gas had disappeared prior to isolation in the outer solar system.  Gas drag must have been more significant in the formation of Jupiter and Saturn. One might worry that the orbital decay of small particles, which are an integral part of our scenario, would have been too fast for them to have been accreted. Particles with stopping time comparable to their orbital time drift fastest. Their orbits decay on a timescale $\\Omega^{-1} (a\\Omega/c_s)^2 \\sim 10^3\\yr$.  By damping the random velocities of small bodies, gas drag can protect them from undergoing destructive collisions. This may result in larger bodies, for which the drift timescales are longer. Moreover, gas drag could have made the isolation timescale in the Jupiter-Saturn region as short as $\\sim 10^4\\yr$ (see equation [\\ref{eq:mintiso}])."
    },
    "0404/astro-ph0404394_arXiv.txt": {
        "abstract": "Microlensing is the only known direct method to measure the masses of stars that lack visible companions.  In terms of microlensing observables, the mass is given by $M=(c^2/4G)\\tilde r_\\e \\theta_\\e$ and so requires the measurement of both the angular Einstein radius, $\\theta_\\e$, and the projected Einstein radius, $\\tilde r_\\e$. Simultaneous measurement of these two parameters is extremely rare. Here we analyze OGLE-2003-BLG-238, a spectacularly bright ($I_{\\rm min} = 10.3$), high-magnification ($A_{\\rm max}= 170$) microlensing event.  Pronounced finite source effects permit a measurement of $\\theta_\\e=650\\,\\muas$. Although the timescale of the event is only $t_\\e = 38\\, {\\rm days}$, one can still obtain weak constraints on the microlens parallax: $ 4.4\\,\\au < \\tilde r_\\e <18\\,\\au$ at the $1\\, \\sigma$ level. Together these two parameter measurements yield a range for the lens mass of $0.36\\,M_\\odot < M< 1.48\\,M_\\odot$.  As was the case for MACHO-LMC-5, the only other single star (apart from the Sun) whose mass has been determined from its gravitational effects, this estimate is rather crude.  It does, however, demonstrate the viability of the technique. We also discuss future prospects for single-lens mass measurements. ",
        "introduction": "} When microlensing experiments were first proposed \\citep{pac86} and implemented \\citep{alcock93,aubourg93,udalski93}, it was not expected to be possible to measure the masses and distances of individual microlenses.  The only microlensing parameter that depends on the mass and that is routinely measured is the Einstein timescale $t_\\e$, which is a degenerate combination of the lens mass $M$, and the lens-source relative parallax, $\\pi_\\rel$, and proper motion, $\\mu_\\rel$. Specifically, \\begin{equation} t_\\e = {\\theta_\\e\\over\\mu_\\rel},\\qquad \\theta_\\e = \\sqrt{\\kappa M \\pi_\\rel},\\qquad \\kappa \\equiv {4 G\\over c^2\\,\\au} \\simeq 8.14 {{\\rm mas}\\over M_\\odot}, \\label{eqn:tedef} \\end{equation} where $\\theta_\\e$ is the angular Einstein radius.  However, \\citet{gould92} showed that if both $\\theta_\\e$ and the microlens parallax, \\begin{equation} \\pi_\\e = \\sqrt{\\pi_\\rel\\over \\kappa M}, \\label{eqn:piedef} \\end{equation} could be measured, then the mass and lens-source relative parallax could both be determined, \\begin{equation} \\label{eqn:massdist} M = {\\theta_\\e\\over \\kappa \\pi_\\e},\\qquad \\pi_\\rel = \\pi_\\e\\theta_\\e. \\end{equation}  Nevertheless, of the roughly 2000 microlensing events detected to date, there have been only of order a dozen for which $\\theta_\\e$ has been measured and a dozen for which $\\pi_\\e$ has been measured. Moreover, there is only one event, EROS-BLG-2000-5, with measurements of {\\it both} parameters, and so for which the microlens mass and distance have been reliably determined \\citep{jin}.  Since this one event was a binary, and since all the other stars with directly measured masses are components of binaries, it remains the case today that the only single star with a directly measured mass is the Sun.  The one partial exception is the microlens in MACHO-LMC-5.  \\citet{alcock01} were able to measure both $\\theta_\\e$ and $\\pi_\\e$ for this event and so measure the mass and distance. These estimates were completely inconsistent with photometry-based estimates of these quantities, but \\citet{gould04} resolved this puzzle by showing that the $\\pi_\\e$ measurement was subject to a discrete degeneracy for this event and that the alternate solution was consistent at the few $\\sigma$ level with the photometric evidence. Nevertheless, since the error in the mass estimate is about 35\\%, this mass determination must be regarded as very approximate.   Here we analyze OGLE-2003-BLG-238, the brightest microlensing event ever observed and only the fourth reported point-lens (i.e., non-binary) event with pronounced finite-source effects.  As with the other three such events \\citep*{alcock97,smith03,yoo}, these finite-source effects allow one to measure $\\theta_\\e$ with reasonably good ($\\sim 10\\%$) precision, where the error is typically dominated by the modeling of the source rather than the microlensing event. Hence, if $\\pi_\\e$ could also be measured, it would be possible to determine $M$.   Despite the event's short duration, it is still possible to detect parallax effects in OGLE-2003-BLG-238 because of its bright source and high magnification. For short events like this one, the Earth's acceleration can be approximated as uniform during the event. \\citet*{gmb} showed that under these conditions, the parallax effect reduces to a simple asymmetry in the lightcurve around the peak. The high magnification of OGLE-2003-BLG-238 permits a very accurate measurement of the peak time of the event, which in turn makes the fitting process very sensitive to this small asymmetry.  The brightness of the source allows high precision photometric measurements even in the wings of the event, which enable detection of these subtle deviations. Unfortunately, as also shown by \\citet{gmb}, the simplicity of the parallax effect for short events implies that only 1-dimensional parallax information can be effectively extracted, whereas the microlens parallax is intrinsically a 2-dimensional vector $\\bpi_\\e$.  That is, while one component of the vector parallax is well determined, the scalar parallax $\\pi_\\e$ is not well determined, and this degrades the mass determination through equation~(\\ref{eqn:massdist}). Nevertheless, this is still only the second single star (other than the Sun) for which any direct mass measurement at all can be made. ",
        "conclusions": ""
    },
    "0404/astro-ph0404488_arXiv.txt": {
        "abstract": "{ One hundred and five 15$\\mu m$ selected objects in three ISO ($Infrared~ Space~ Observatory$) deep survey fields (CFRS 3$^{h}$, UDSR and UDSF) are studied on the basis of their high quality optical spectra with resolution R$>$1000 from VLT/FORS2. Ninety two objects (88\\%) have secure redshifts, ranging from 0 to 1.16 with a median value of $z_{\\rm med}$=0.587. \\newline Considerable care is taken in estimating the extinction property of individual galaxy, which can seriously affect diagnostic diagrams and estimates of star formation rates (SFRs) and of metal abundances. Two independent estimates of the extinction have been made, e.g. Balmer line ratio and energy balance between infrared (IR) and H$\\beta$ luminosities. For most of the sources, we find a good agreement between the two extinction coefficients (within $\\pm$0.64 rms in $A_{V}$, the extinction in V band), with median values of $A_V$(IR) = 2.36 and $A_V$(Balmer)= 1.82 for z$>$0.4 luminous IR galaxies (LIRGs). At z$>$0.4, our sample show many properties (IR luminosity, continuum color, ionization and extinction) strikingly in common with those of local (IRAS) LIRGs studied by Veilleux et al. (1995). Thus, our sample can provide a good representation of LIRGs in the distant Universe.\\newline We confirm that most ($>$77\\%) ISO 15$\\mu m$ selected galaxies are dominated by star formation. Oxygen abundances in interstellar medium in the galaxies are estimated from the extinction corrected ``strong\" emission line ratios (e.g. \\ion{[O}{ii]}/H$\\beta$, \\ion{[O}{iii]}/H$\\beta$ and \\ion{[O}{iii]}/\\ion{[O}{ii]}). The derived 12+log(O/H) values range from 8.36 to 8.93 for the z$>$0.4 galaxies with a median value of 8.67. Distant LIRGs present a metal content less than half of that of the local bright disks (i.e. $L^*$). Their properties can be reproduced with infall models though one has to limit the infall time to avoid overproduction of metals at late time. The models predict that total masses (gas + stars) of the distant LIRGs are from $10^{11} M_{\\odot}$ to $\\le$ $10^{12} M_{\\odot}$. A significant fraction of distant large disks are indeed LIRGs. Such massive disks could have formed $\\sim$50\\% of their metals and stellar masses since z$\\sim$1. ",
        "introduction": "The IRAS all-sky survey detected tens of thousands of galaxies with far-infrared (far-IR) radiation luminosities from less than $10^{6} L_{\\odot}$ to $\\sim$$10^{13} L_{\\odot}$ up to a moderate redshift ($z$$\\sim$0.3). However, luminous infrared galaxies (LIRGs) are not typical of local galaxy population, and they account for only $\\sim$2\\% of the local bolometric luminosity density (Soifer et al. 1987; Sanders \\& Mirabel 1996).   However, the COsmic Background Explorer (COBE) observations imply  that there likely exists a very significant contribution of dust-obscured star formation at high redshifts  (Puget et al. 1996; Genzel \\& Cesarsky 2000). The ISO made it possible to study the infrared emission of galaxies at $z \\geq 0.5$, which plays an important role in understanding the co-moving star formation density evolution with look-back time. The ISO  mid infrared camera (ISOCAM) (Cesarsky et al. 1996)  is $\\sim$$10^{3}$ times more sensitive and has 60 times higher spatial resolution than IRAS. The mid-IR ISOCAM 15 $\\mu$m source counts provide evidence for strong IR light density evolution, as revealed by the strong excess of 15 $\\mu$m counts above the predictions of  non-evolution models at sub-mJy (Elbaz et al. 1999; Aussel et al. 1999). The cosmic infrared background resolved by ISOCAM shows that the co-moving density of  infrared light due to the luminous IR galaxies ($L_{\\rm IR}\\geq$10$^{11}L_{\\odot}$)  was more than 40 times larger at $z\\sim$1 than today (Elbaz et al. 2002).  The main driver for this evolution is the  luminous infrared starburst galaxies seen by ISO at $z$$>$0.4,  which form stars at a rate of more than $50M_{\\odot}$\\,yr$^{-1}$ (Flores et al. 1999).  Based on the correlation analysis of deep X-ray and mid-IR observations in Lockman Hole and Hubble Deep Field North (HDF-N), Fadda et al. (2002)  found that the active galactic nuclei (AGN) contribution to the 15 $\\mu$m background is only 17$\\pm$2\\%. They concluded that the population of IR luminous galaxies detected in the ISOCAM deep surveys, and the cosmic infrared background sources themselves,  are mostly dust-obscured starbursts (also see Elbaz et al. 2002). Reviews of extragalactic results from ISO can be found in Genzel \\& Cesarsky (2000),  Franceschini et al. (2001) and Elbaz \\& Cesarsky (2003).   Recently, Flores et al. (2004a) studied the interstellar extinction and SFRs of 16 luminous infrared galaxies in Canada-France Redshift Survey (CFRS) 3$^{h}$ and 14$^{h}$ fields using the spectra from the European Southern Observatory (ESO) Very Large Telescope (VLT) and Canada-France-Hawaii Telescope (CFHT). They found that the extinction coefficients obtained from H$\\gamma$/H$\\beta$ (using the VLT/FORS2 or CFHT spectra)  and H$\\alpha$/H$\\beta$ (combining the VLT/FORS2 and VLT/ISAAC spectra)   are in agreement, and that SFRs derived from H$\\alpha$ are consistent with those from infrared luminosities, except for the galaxies near the  ultra luminous IR galaxy (ULIRG) regime ($L_{\\rm IR}$ $>$ $10^{12}$ $L_{\\odot}$).   Spectrophotometric properties of IRAS galaxies have been studied in detail, providing a full diagnostic of their ISM properties (Veilleux et al. 1995, hereafter V95; Kim et al. 1995). However, at higher redshifts, studies of ISOCAM sources have mostly focused on source counts and SFRs. Very few is known about chemical properties of distant LIRGs, including about their metal content, and to fill this gap is the main objective of this paper.   In the local Universe, metallicity is well correlated with the absolute luminosity (stellar mass) of galaxies over a wide magnitude range (e.g. 7-9 mag)  (Zaritsky et al. 1994; Richer \\& McCall 1995; Telles \\& Terlevich 1997; Contini et al. 2002; Melbourne \\& Salzer 2002; Lamareille et al. 2004). Some results have been obtained on  the luminosity-metallicity (L-Z) relations in the intermediate-redshift Universe. Kobulnicky \\& Zaritsky (1999) found that the L$-$Z  relations of 14 intermediate-$z$  emission line galaxies  with $0.1<z<0.5$ are consistent with those   of the local spiral and irregular galaxies   studied by Zaritsky et al. (1994), Telles \\& Terlevich (1997) and Richer \\& McCall (1995). The 16 CFRS galaxies at $z\\sim$0.2 studied by Liang et al. (2004)  fall well in the region occupied by the local spiral galaxies (from Zaritsky et al. 1994).   At higher redshifts, Kobulnicky et al. (2003) have obtained the L$-$Z relations of 64 galaxies from the Deep Groth Strip Survey (DGSS) which have been separated into  three redshift ranges ($z$=0.2$-$0.4, 0.4$-$0.6, and 0.6$-$0.82).  In the highest redshift bin galaxies are brighter by $\\sim$1 mag relatively to those in the lowest redshift bin and  brighter by $\\sim$2.4 mag compared to the local ($z$$<$0.1) field galaxies  (from Kennicutt 1992a,b, hereafter K92, and Jansen et al. 2000a,b, hereafter J20). Such a result is confirmed by Maier et al. (2004).  These studies contrast with the results of Lilly et al. (2003), who have found that the L$-$Z relation of most of their 66 CFRS galaxies with $0.5<z<1$ is similar to that of the local galaxies from J20. However, Lilly et al. (2003) have assumed a constant $A_{V}$=1 for accounting for dust extinction. In this study, we investigate the L-Z relation for LIRGs in $z>0.4$ Universe, after a detailed account for their dust extinction properties.    This paper is organized as it follows.  In Sect.2,  we describe the sample selection, the observations and the data reduction and analysis, while the redshift distribution and the spectrophotometric properties are presented in Sect.3.  Sect. 2 and 3 are aiming at assessing whether our resulting sample can be used to test the properties of distant LIRGs. Flux measurements, interstellar extinction and SFRs  of the galaxies are shown in Sect.4. It includes a detailed comparison between extinction parameter deduced from Balmer line ratio to that derived from mid-IR luminosity.   In Sect.5, we discuss the diagnostic diagrams to test the AGN contribution as well as the ionization properties. In Sect.6, we present the luminosity-metallicity relation (based on oxygen abundances) of distant LIRGs which is compared to other galaxy samples.  Discussion and conclusion are given in Sections 7 and 8. Throughout this paper, a cosmological model with  $H_0$=70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega _M$=0.3 and $\\Omega _\\Lambda =0.7$ has been adopted. ",
        "conclusions": "A large sample (105) ISO/15$\\mu$m-selected sources in three ISO deep survey fields (CFRS 3$^{h}$, UDSR and UDSF), are studied on the basis of their high quality VLT/FORS2 spectra and the infrared data from ISO. Among the 92 redshift-identified objects, 75 (64 with $z>0.4$) are classified to be EL galaxies. 66 (55 with $z>0.4$) objects are EL galaxies out of the 77 ISOCAM 15 $\\mu$m detected sources. This is by far the largest sample of spectra of distant ISO galaxies. We present here their properties derived from the emission lines. The redshift distribution ($z_{\\rm med}$=0.587) is consistent with that of previous studies, and some galaxies belong to a $z$$\\sim$0.70 cluster (in the UDFS field) or to a $z$$\\sim$0.166 galaxy group (in the UDSR field). This study provides us:\\newline - an accurate calculation of the extinction coefficient from Balmer emission lines, which properly account for underlying Balmer absorption ($A_V$(Balmer), median value of 1.82 for $z$$>$0.4 sample galaxies)\\newline - an energy balance between IR and Balmer luminosities which provides us an original way to calculate an IR-based extinction coefficient ($A_V$(IR), median value of 2.36 for $z$$>$0.4 galaxies)\\newline - a good correlation between the above-mentioned two extinction coefficients, which, for the first time, allows us to reliably obtain the chemical properties of distant galaxies with considering interstellar extinction; this diagram could also be used as a diagnostic of their dust geometry and star formation history\\newline - the very similar general properties of $z$$>$0.4 ISO sources to those of the local IRAS LIRGs studied by V95, including IR luminosity distribution, continuum color, ionization and extinction properties, which ensures that our sample is representative of LIRGs in the distant Universe\\newline - an extinction corrected diagnostic diagram based on \\ion{[O}{ii]}/H$\\beta$ and \\ion{[O}{iii]}/H$\\beta$ ratio, which confirms the small fraction ($\\sim$23\\%) of Seyferts and LINERs for ISOCAM 15 $\\mu$m sources found in other studies (Flores et al. 1999; Fadda et al. 2002; Elbaz et al. 2002)\\newline - a determination of O/H metallicity, which ranges from 8.36 to 8.93 in 12+log(O/H) for the $z$$>$0.4 sample; no obvious trend is found between the IR luminosity and the metallicity. \\newline  This study shows that distant LIRGs draw a peculiar sequence in the $M_{B}$-O/H diagram, which is a broad L-Z relation at bright B absolute luminosities. When compared to local bright disks, at constant metallicity (12+log(O/H)=8.67, median value), LIRGs are from 2.5 to 5 magnitudes more luminous, and at constant $M_{B}$ ($M_{B}$= $-$21.24, median value), they are more than two times deficient in metal. These properties can be reproduced with infall models though one has to limit the infall time to avoid overproduction of metals in late time. Models predict a total mass ranging from $10^{11}$$M_{\\odot}$ to $\\le$$10^{12}$$M_{\\odot}$, which can be twice of the stellar masses of distant LIRGs derived by Zheng et al. (2003) on the basis of the K band luminosities. In a forthcoming paper (Hammer et al. 2004) we investigate whether LIRG properties could be related to a recent and significant star formation in massive galaxies, including spirals."
    },
    "0404/astro-ph0404441_arXiv.txt": {
        "abstract": "We present a detailed study of a simple scalar field model that yields non-singular cosmological solutions. We study both the qualitative dynamics of the homogeneous and isotropic background and the evolution of inhomogeneous linear perturbations. We calculate the spectrum of perturbations generated on super-Hubble scales during the collapse phase from initial vacuum fluctuations on small scales and then evolve these numerically through the bounce. We show there is a gauge in which perturbations remain well-defined and small throughout the bounce, even though perturbations in other commonly used gauges become large or ill-defined. We show that the comoving curvature perturbation calculated during the collapse phase provides a good estimate of the resulting large scale adiabatic perturbation in the expanding phase while the Bardeen metric potential is dominated by what becomes a decaying mode after the bounce. We show that a power-law collapse phase with scale factor proportional to $(-t)^{2/3}$ can yield a scale-invariant spectrum of adiabatic scalar perturbations in the expanding phase, but the amplitude of tensor perturbations places important constraints on the allowed initial conditions. ",
        "introduction": "\\let\\include\\input   Current cosmological observations, especially measurements of the spectrum of temperature anisotropies and polarisation of the cosmic microwave background, provide strong evidence for the existence of primordial density perturbations on large scales, far larger than the Hubble scale at the time of last-scattering. Observations are consistent with an almost scale-invariant, Gaussian distribution of adiabatic primordial perturbations~\\cite{WMAP}.  The standard explanation for the origin of the primordial perturbations is that the Universe underwent an accelerated, inflationary, expansion in the early universe~\\cite{LLbook}. Zero-point vacuum fluctuations in a massless scalar field naturally give rise to a scale-invariant distribution of effectively classical field fluctuations on super-Hubble scales during an exponential (de Sitter) expansion. If inflation is driven by a light, slowly-rolling, inflaton field then these field fluctuations can be directly related to a comoving curvature perturbation, $\\R$, which remains constant for adiabatic perturbations on super-Hubble scales until a given wavelength re-enters the horizon in the subsequent radiation- or matter-dominated eras.  It is important to consider whether or not inflation is the {\\em only} consistent model for the origin of large-scale structure in our Universe. An alternative proposal is that the hot big bang and the primordial perturbations originate not from an accelerated expansion, but from a preceding collapse. In classical general relativity a collapsing universe may be doomed to a singularity, but in a finite quantum theory of gravity, such as string theory, the spectrum of states going into the ``singularity'' would be expected to determine some out-going state. In particular the collapse phase could set the initial conditions for a subsequent post-big bang phase. Gasperini and Veneziano originally proposed \\cite{GV} a pre big bang scenario based on the superstring low energy effective action \\cite{pbbreviews}. More recently Khoury et al proposed the ekpyrotic scenario \\cite{ekpyrotic} motivated by colliding brane solutions in heterotic M-theory (see also \\cite{pyro,cyclic} for later developments of this scenario). Both models invoke a collapsing universe in the 4D Einstein frame where the energy density is dominated by minimally coupled scalar fields.  In a collapsing universe zero-point vacuum fluctuations with a given wavelength are squeezed outside the rapidly shrinking Hubble scale, also naturally giving rise to a spectrum of field fluctuations on super-Hubble scales. However in either pre big bang or ekpyrotic scenarios, vacuum fluctuations give rise to a steeply blue-tilted spectrum of comoving curvature perturbations during the collapse \\cite{BGGMV,Lyth}. If the tilt of the comoving curvature spectrum remains unchanged during the transition from pre to post big bang phase then this cannot seed the observed large scale structure in our Universe. One would then have to consider alternative sources for structure after a collapse phase, such as the curvaton mechanism \\cite{curvaton}, or consider a different collapse model. Collapse driven by matter or a scalar field with pressureless equation of state can produce a scale-invariant spectrum of comoving curvature perturbations~\\cite{Wands99,FB}.  Although very general physical arguments have been advanced that the comoving curvature perturbation, $\\R$, should be conserved on super-horizon scales for adiabatic density perturbations \\cite{Lyth85,DM,MS,WMLL,LW03} there are several possible ambiguities that may arise in the transition from collapse to bounce, not present in conventional models of an inflationary expansion. At any bounce the comoving Hubble scale diverges and all scales are inside the Hubble scale, at least for an instant~\\cite{MP}. And the Bardeen metric potential, $\\Phi$, does not have the same spectral dependence as the comoving curvature perturbation during the collapse \\cite{BF,Hwang,Khoury2,Lyth2,HwangNoh,DV}, even though the two tilts must coincide for adiabatic perturbations in the subsequent expanding universe. As a result it is impossible for both $\\R$ and $\\Phi$ to obey non-singular evolution equations through the bounce \\cite{CDC}.  In fact the Bardeen potential on super-Hubble scales grows rapidly in a collapsing phase.  This was noted in Ref.~\\cite{BGGMV} who realised that the divergence of perturbations in the Bardeen's longitudinal (or conformal Newtonian) gauge could be eliminated by a gauge-transformation to the spatially flat (or off-diagonal) gauge. However we will see that at the moment of bounce the spatially flat gauge itself becomes ill-defined. The comoving curvature perturbation may remain small during collapse but it too becomes ill-defined in the vicinity of the bounce in our model. The breakdown of perturbation theory in any given gauge does not necessarily signal the breakdown of perturbation theory. It is well known that the comoving curvature perturbation becomes singular at stationary values of the field even in single-field inflation models, but there are other gauge choices where perturbations remain small~\\cite{Kodama}.  Several authors have studied the evolution of perturbations through a bounce using some analytic form for the background scale factor \\cite{PPN,PPNG,Finelli:2003mc} and/or the evolution equations for perturbations \\cite{CDC}. But this has lead to apparently contradictory results depending on which functions are assumed to have a smooth regular form \\cite{CDC}.  In order to investigate some of these issues associated with the evolution of cosmological perturbations through the bounce we will consider a simple scalar field model for a non-singular bounce. We are following the general approach envisaged in the pre big bang scenario \\cite{GV} where higher-order corrections to the effective action will become significant at high curvatures or strong coupling and can yield non-singular bounce solutions \\cite{nonsing,Ed}.  In the ekpyrotic or cyclic models the transition is supposed to occur at a real singularity in the effective-4D theory. Not only background quantities but also perturbations diverge here. So our 4D analysis will not address arguments about the regularisation schemes at the singularity proposed so far \\cite{Khoury2,Schwarz,Tolley}.  There have already been a few numerical studies of the evolution of perturbations through non-singular bounce solutions including higher-order corrections to the string effective action \\cite{Tsujikawa:2001ad,Tsujikawa:2002qc,CartierHwang,Cartier}. It can be shown that that the null energy condition must be violated for a bounce to occur in general relativity in a spatially flat 4D FRW cosmology \\cite{Brustein}.  (This is quite different from non-singular classical bounces which have been studied in closed FRW models \\cite{ParkerFulling,Starobinsky78,Matzner,Hwangbounce,GordonTurok,Deruelle,DeruelleStreich,MP} which do not violate the null energy condition, but require fine-tuned initial conditions to obtain a sufficiently long collapse phase \\cite{Starobinsky78,Matzner}.) One problem with results based on higher-order corrections near the bounce is that the evolution close to the bounce becomes dominated by the ``corrections'' and one cannot be sure that even higher-order corrections would not also become important \\cite{Ed}. Another complication is that increasingly higher-order corrections are equivalent to introducing additional fields with non-canonical kinetic terms \\cite{TT,Wands93} so that the true dynamical degrees of freedom become obscured.  In order to make the physical degrees of freedom manifest we will work with scalar fields obeying canonical wave equations in four-dimensional Einstein gravity. The price we must pay for such simplicity, while still obtaining a non-singular bounce, is that one of the scalar fields will be a massless ``ghost'' field with negative kinetic energy, in order to violate the null energy condition near the bounce. Such a ghost field would lead to the presence of serious instabilities in any low energy theory \\cite{ghosts}. However it does not play any significant role in our model at early or late times. Instead it leads to precisely the sort of instability at high energies that we require to achieve a bounce in the high-energy regime, destabilising the pre big bang collapse phase and triggering the transition to an expanding cosmology. Our scalar field model is intended only as a simple model of a non-singular bounce and not as a realistic model of the low energy universe. The early- and late-time (low energy) evolution in our model is dominated by a conventional minimally-coupled scalar field with exponential potential. By choosing the dimensionless slope of the potential we are able to investigate different power-law solutions for the pre big bang collapse leading to different power spectra for the field and metric perturbations \\cite{Wands99}. A very similar bounce model was studied by Peter and Pinto-Neto \\cite{PPN}, but we find very different results.  In the original pre big bang scenario higher-order corrections to the effective action will become significant at high curvatures or strong coupling and can yield non-singular bounce solutions \\cite{nonsing,Ed}, at least for some initial conditions and some higher-order corrections. In particular it can be shown that the null energy condition must be violated for a bounce to occur in a 4D FRW cosmology \\cite{Brustein}.    We first study the background evolution of our model using a phase-plane analysis in Section \\ref{background}. With a particular choice of dimensionless phase variables the early or late time solutions correspond to critical points in the phase plane. However an alternative choice of variables is better suited to study evolution through the bounce. We identify a class of non-singular models which start in a power-law collapse in the asymptotic past, pass smoothly through a bounce and approach a power-law expansion.  We then define linear perturbations about this background model in Section \\ref{perturb}. We present the evolution and constraint equations for the perturbations in arbitrary gauge and then identify particular gauge choices that are well-suited to study the perturbations at early and late times, or through the bounce, analogously to the different choice of background variables used to study the background phase-plane in different regimes. The choice of gauge to take the perturbations through the bounce is vital as it must be ensured that the gauge itself is well defined throughout the bounce, and we identify problems with gauges commonly used in inflationary models.  In section \\ref{evolve} we evolve first-order evolution equations for the metric and field perturbations and use the constraint equations as a consistency check of our calculation. We assume perturbations begin in the quantum vacuum state on small scales and use this to calculate the spectrum of perturbations generated on super-Hubble scales during collapse. We evolve this numerically through the bounce in order to calculate the resulting spectrum of curvature perturbations. We present our results in section \\ref{results} in terms of the gauge-invariant variables ${\\cal R}$ and $\\Phi$. We find that it is possible to produce a scale-invariant spectrum of adiabatic perturbations on super-Hubble scales from a collapse phase in our simple model. Finally we present our conclusions in Section \\ref{conc}. ",
        "conclusions": "\\label{conc}  We have presented a simple model of a four-dimensional FRW universe that describes non-singular bouncing solutions. To do so within the context general relativity it is well-known that the null-energy condition must be violated. To achieve this we have introduced a massless ``ghost'' field with negative kinetic energy. A similar role is played by higher-order corrections to the low-energy string effective action which have recently been studied by Cartier~\\cite{Cartier} who found qualitatively similar results. In order to keep the dynamics as transparent as possible we have restricted ourselves to a canonical (second-order) scalar field Lagrangian albeit with the wrong sign. Like the higher-order loop corrections, our ghost field only plays a significant dynamical role in the high energy regime close to the bounce.  We have performed a phase-plane analysis of the qualitative behaviour of expanding and contracting FRW models with both a canonical scalar field (with positive energy density) and an exponential potential plus the massless ghost field. Away from the bounce we have usual power-law expansion or contraction for a universe dominated by a scalar field with exponential potential \\cite{LM85,Halliwell,BB}. The dimensionless slope of the potential, $\\lambda$, determines the asymptotic power-law index $p=2/\\lambda^2$ and the equation of state $w=(\\lambda^2-3)/3$ for sufficiently flat potentials with $\\lambda^2<6$.  It is known that for $\\lambda^2<6$ these power-law collapse solutions describe the generic early time behaviour but are unstable at late times to kinetic-dominated collapse with $w=1$ \\cite{Heard,Gratton} or to shear anisotropy~\\cite{Kunze}. Some kind of instability is a necessary to allow an escape from the singular power-law collapse. For steep potentials with $\\lambda^2\\geq6$ we find no non-singular solutions. But for flat potentials with $\\lambda^2<6$ we find a class of non-singular solutions that begin in the asymptotic past as a power-law collapse, approach a minimum value of the scale factor and then approach a power-law expansion at late times.  This class of non-singular solutions provides a simple bounce model in which we then study the evolution of linear perturbations. We assume that these perturbations begin in the flat-spacetime vacuum state at early times and use this to predict the amplitude of large scale (super-Hubble) adiabatic density perturbations after the bounce.  Although both the Bardeen potential $\\Phi$ and comoving curvature perturbation ${\\cal R}$ evolve smoothly at the instantaneous bounce, they both have potential problems. The Bardeen potential grows rapidly during the collapse phase and is much greater than unity for the modes of interest, calling into question the validity of linear perturbation theory in the longitudinal gauge~\\cite{Lyth2,PPN1}. On the other hand the comoving gauge becomes singular whenever the total kinetic energy of the scalar fields vanishes~\\cite{Kodama}, which occurs shortly before and after the bounce in our model.  We can however find a gauge in which the metric perturbations remain small and linear perturbation theory should be valid. In our model the massless ghost field is monotonic and provides a suitable time variable, defining a well-defined choice of gauge (the uniform $\\chi$-field gauge) for perturbations. We note that we solve the regularized unconstrained evolution equations for the perturbations, using the constraint equations as a check of our numerical integration. This avoids potential problems that could arise from using the constraint equations to eliminate variables whose coefficients vanish at the bounce~\\cite{Cartier}. We can then reconstruct the usual gauge-invariant variables from calculations in this gauge.  We find that the comoving curvature perturbation calculated during the collapse phase is a good estimator of the resulting adiabatic density perturbation on super-Hubble scales in the expanding phase. Note that the value at the end of the collapse phase is not in general the same as the that at Hubble-crossing during collapse (unlike the usual case in an inflationary expansion). This is because the dominant mode during collapse is a non-adiabatic field perturbation. However the comoving curvature perturbation remains constant on super-Hubble scales after the bounce.  By contrast the Bardeen potential calculated in the collapse phase is not a good indicator of the subsequent large scale adiabatic perturbation in the expanding phase. It is dominated by the shear in the uniform $\\chi$-field gauge, which grows rapidly in the collapse phase and then decays away after the bounce, until the Bardeen potential settles down to a constant value on super-Hubble scales, related to the comoving curvature perturbation by the usual expression $\\Phi=3(1+w){\\cal R}/(5+3w)$. This basic point has been made by many authors \\cite{Lyth,BF,Hwang,LW03} and here we demonstrate it explicitly for our simple bounce model.  In particular we find that the spectral tilt of the final adiabatic curvature perturbation on super-Hubble scales in the expanding phase is given by the spectral tilt of the initial comoving curvature perturbations (or equivalently the scalar field fluctuations in the spatially flat gauge). We therefore have given an explicit model for a power-law collapse with $\\lambda=\\sqrt{3}$ which yields a scale-invariant of adiabatic perturbations on super-Hubble scales in the expanding phase.  The amplitude of the comoving curvature perturbation on the smallest scale (that is the minimum value of the Hubble scale during collapse) is set by the energy scale relative to the Planck scale (where vacuum fluctuations in the metric become of order unity) \\be {\\cal P}_{\\cal R} (k=aH_{\\rm max}) \\sim \\left( \\frac{H_{\\rm max}}{M_{\\rm Pl}} \\right)^2 \\,. \\ee This is very similar to the amplitude of gravitational waves (tensor perturbations) produced during collapse. The tensor-scalar ratio is an important constraint on conventional inflation models and turns out to be a severe constraint on our collapse model when $\\lambda=\\sqrt{3}$.  Pre big bang or collapse models have been criticised for the degree of fine-tuning required to set up the initial conditions at the start of any collapse phase~\\cite{Turner,Linde,Gratton}. The dimensionless parameter $C$, defined in Eq.~(\\ref{defC}), is a first integral of the equations of motion which selects the trajectory in the homogeneous phase-space. Although any value for $|C|<1/\\sqrt{2}$ yields a non-singular bounce for models with $\\lambda=\\sqrt{3}$, only trajectories with $|C|>0.67$ give an acceptably small tensor-scalar ratio. This provides a quantitative measure of the degree of fine-tuning of initial conditions required to get an acceptable primordial density perturbation within the class of models considered in this paper."
    },
    "0404/astro-ph0404327_arXiv.txt": {
        "abstract": "\\targ~is a new X-ray transient, discovered by \\inte~during an observation of \\GRS. The source presents strong variations on timescales from seconds to days. We present results of multiwavelength observations, including spectral analysis of \\inte~observations, and propose that \\targ~is a Galactic X-ray binary. Further classification of the source is also discussed. ",
        "introduction": "X-ray transients are a subclass of X-ray binaries whose brightness can vary by a factor of more than 100. The main reason for those variations are changes in the mass accretion rate on to the compact object. The variations of the mass accterion rate can be related to an eccentric orbit, changes in the companion star, or instabilities of the accretion disk.  \\targ~(SIMBAD corrected name IGRJ19140+0951) is an X-ray transient that was detected by \\inte~\\citep{diana03a} on March 6th during an observation of \\GRS. It has also been observed later by Rossi X-ray Timing Explorer (RXTE) and ground-based optical and radio telescopes such as the Nordic Optical Telescope (NOT) and the Giant Meterwave Radio Telescope (GMRT). ",
        "conclusions": "A new X-ray source, \\targ~has been discovered with \\inte. Analysis of \\inte~and RXTE data shows that the spectral and temporal variability of \\targ~are best explained by a Galactic X-ray binary, preferably one with a low-mass donor and a black hole. However, further observations are needed to determine the binary component types."
    },
    "0404/astro-ph0404111_arXiv.txt": {
        "abstract": "We present an attempt to constrain deviations from the gravitational inverse-square law on large-scale structure scales.  A perturbed law modifies the Poisson equation, which implies a scale-dependent growth of overdensities in the linear regime and thus modifies the power spectrum shape. We use two large-scale structure surveys (the Sloan Digital Sky survey and the Anglo-Australian Two-degree field galaxy redshift survey) to constrain the parameters of two simple modifications of the inverse-square law.  We find no evidence for deviations from normal gravity on the scales probed by these surveys ($\\sim 10^{23}$m.) ",
        "introduction": "Einstein's theory of general relativity, which generalizes Newton's law of gravity, has been extensively tested and verified.  But precision tests of the inverse-square force law have been performed only on scales $\\le 10^{13}$ m (\\cite{Adelberger} and references therein). An enormous extrapolation is required to apply this law to cosmological scales $> 10^{22}$m \\cite{Peebles03}.  Some recent theories propose a deviation from Newtonian gravity at large distances, perhaps due to extra dimensions, branes, or new particles (e.g.\\cite{DGP}).  These theories, mostly motivated by finding an alternative explanation to dark energy for the cosmic acceleration (e.g.\\cite{Arkani-Hamed02,DDG02, Carroll03, Lue1}), modify gravity on scales comparable to the horizon scale $\\sim 10^{26}$m.  Here we set out to constrain theories in which gravity deviates from the inverse-square law on scales of $\\sim$10 Mpc or $10^{23}$m. We evolve the power spectrum (and bispectrum) under a perturbed law of gravity in the linear (and mildly non-linear) regimes, from recombination to the present day.  The perturbed law modifies the evolution of density fluctuations, altering, in particular, the power spectrum shape.  We compare the prediction for the modified power spectrum to that measured from two galaxy redshift surveys: the Anglo-Australian Two-degree Field Galaxy Redshift Survey (2dFGRS; \\cite{Colless2dF,Percival01}) and the Sloan Digital Sky Survey (SDSS; \\cite{GunnSDSS,YorkSDSS,Teg03}).  In the models we consider, gravity is consistent with the inverse-square law on scales smaller than $\\sim$1 Mpc, so the deviation does not affect early universe physics through recombination.  Thus, we assume that the standard computation of the cosmic microwave background (CMB) anisotropies holds, and that the CMB observations provide the mass power spectrum (and bispectrum) at recombination.  Moreover, as explained later, we marginalize over reasonable priors for the power spectrum amplitude and the primordial power spectrum slope, thus making our analysis insensitive to reasonable changes in the other cosmological parameters. We expect deviations from the inverse-square law to modify the integrated Sachs-Wolfe (ISW) effect.  However, as this affects the CMB on large scales where cosmic variance is large, we do not discuss ISW here.  In the context of Newtonian gravity, a similar approach has been widely used in the literature to constrain cosmological parameters (e.g., \\cite{Spergel03}) and biasing, the relation between clustering of galaxies and clustering of mass (e.g.,\\cite{Lahav02, Verde02}). Here, we fix the background cosmology \\cite{Spergel03} and assume a scale-independent bias \\footnote{Justification of this assumption, used in all analyses of large-scale structure, relies on the statistical properties of initially gaussian random fields \\cite{Kai86b, Weinberg97}.}. The deviation from Newtonian gravity is parameterized by a length-scale and a strength on which we place constraints. We find no compelling evidence for deviations from the inverse-square force law on large-scale structure scales. This null result leaves open the possibility that both Newton's gravity and our weak assumptions about scale-independent bias and the primordial power spectrum are incorrect in some contrived way so that their effects conspire to exactly cancel.  We regard such a scenario as exceedingly unlikely. This is the first attempt to use cosmological data to constrain gravity on these scales, but for related theoretical work, see \\cite{Uzan, Peebles02, Barrow98, Nusser04}. ",
        "conclusions": "Figure (\\ref{fig:yukawa}) shows the Yukawa-like model's parameter constraints for the 2dFGRS and SDSS data, and figure (\\ref{fig:N2}) shows the PL model's constraints.  The ``ringing'' in the $\\hat d$ function in the PL model, creates ``wiggles'' in the large scale structure power spectrum which are responsible for the geometry of the iso-likelihood contours in fig 3.  The left plots show the likelihood contours for one-sigma marginalized, one-sigma joint, and two-sigma joint confidence levels for one survey in $\\l$,$\\a$ (or $R$,$\\e$) space. The center plots show the difference in chi square ($\\Delta \\chi^2$) as a function of $\\a$ or $\\e$, marginalized over $\\l$ or $R$.  We used data up to $k \\sim 0.15$ h/Mpc for both surveys.  The results are consistent with each other and with a strict inverse-square law. Since the constraints are weak and the Taylor expansion analysis breaks down at large values of $|\\a|$ and $|\\e|$, we extrapolated the error bars from the graphs. We find $\\a =0.025 ^{+1.7}_{-1.7}$, $\\e = 0.05 ^{+1.3}_{-0.4}$ for 2dFGRS and $\\a = -0.35^{+0.9}_{-0.9}$, $\\e = 0.45 ^{+1.4}_{-1.4}$ for SDSS, at the one-sigma level.  Non-linearities become increasingly important as $k$ increases and might introduce systematic errors, but the signal-to-noise for the power spectrum increases with increasing $k$.  Therefore we explore the dependence of the result as a function of the maximum $k$ included in the analysis ($k_{\\rm max}$) on the right panel of the figures.  Since our analysis probes deviations in the {\\it shape} of today's power spectrum, the strongest constraints on $\\a$ and $\\e$ occur at scales where the {\\it slope} of $\\hat{d}$ (with $a=1$) is greatest. This typically corresponds to $\\lambda$ and $R$ towards the smallest scales of the data, so the smallest comoving scales probed have just reached the physical scale of the force-law deviation today, and the larger comoving scales have passed it.  (A deviation near a physical scale $\\lambda$ or $R$ has only affected comoving scales that have grown larger than $\\lambda$ or $R$.) The weakness of our constraints arises primarily from the need to marginalize over a wide range of power spectrum amplitudes and the spectral slope. Future weak lensing surveys will measure the power spectrum for dark matter directly.  If $a(t)$ can also be measured independently, this data will eliminate the need to marginalize over the amplitude and shrink our error bars for the PL model.  We forecast that this would roughly halve the error on $\\e$.  The constraint does not improve significantly for the Yukawa-like model."
    },
    "0404/astro-ph0404261_arXiv.txt": {
        "abstract": "We report on the eclipse mapping analysis of an ensemble of light curves of the dwarf nova V2051 Oph with the aim to study the spatial distribution of its steady-light and flickering sources. The data are combined to derive the orbital dependency of the steady-light and the flickering components at two different brightness levels, named the 'faint' and 'bright' states. The differences in brightness are caused by long-term variations in the mass transfer rate from the secondary star. The white dwarf is hardly affected by the long-term changes. Its flux increases by only 10 per cent from the faint to the bright state, whereas the disk flux raises by a factor of 2. Eclipse maps of the steady-light show asymmetric brightness distributions with enhanced emission along the ballistic stream trajectory, in a clear evidence of gas stream overflow. The comparison between the steady-light maps of the faint and bright states suggests that the quiescent disk responds to changes in mass transfer rate in an homologous way. The ability to separate the orbital dependency of the low- and high-frequency flickering components allowed us to identify the existence of two different and independent sources of flickering in V2051~Oph. The low-frequency flickering arises mainly in the overflowing gas stream and is associated to the mass transfer process. It maximum emission occurs at the position of closest approach of the gas stream to the white dwarf, and its spatial distribution changes in response to variations in the mass transfer rate. High-frequency flickering originates in the accretion disk, showing a radial distribution similar to that of the steady-light maps and no evidence of emission from the hot spot, gas stream or white dwarf. This disk flickering component has a relative amplitude of about 3 per cent of the steady disk light, independent of disk radius and brightness state. If the disk flickering is caused by fluctuations in the energy dissipation rate induced by magneto-hydrodynamic turbulence, its relative amplitude lead to a viscosity parameter $\\alpha_{cool} \\simeq 0.1 - 0.2$ at all radii for the quiescent disk. This value seems uncomfortably high to be accommodated by the disk instability model. ",
        "introduction": "\\label{introduction}  One of the long standing unsolved problems in accretion physics is related to the cause of flickering. The term is used to name the intrinsic, random brightness fluctuations of 0.01--1 magnitudes on timescales from seconds to dozens of minutes that are seen in the light curves of T~Tauri stars (e.g., Herbst \\& Shevchenko 1999), mass-exchanging binaries (e.g., Augusteijn et~al. 1992; Baptista, Bortoletto \\& Harlaftis 2002; Sokoloski, Bildsten \\& Ho 2001; Bruch 2000 and references therein) and AGNs (e.g., Garcia et~al. 1999). Flickering is considered a basic signature of accretion (Warner 1995). Although the study of flickering (a time dependent phenomenon) may yield crucial information to the understanding of the viscosity is accretion disks, flickering is still the least studied and one of the most poorly understood of the accretion phenomena.  A promising step towards a better understanding of flickering consists in using indirect imaging techniques such as eclipse mapping (Horne 1985; Baptista 2001) to probe the surface distribution of the flickering sources in Cataclysmic Variables (CVs) [e.g., Welsh \\& Wood 1995]. In these close binaries mass is fed to a white dwarf (the primary) by a Roche lobe filling companion star (the secondary) via an accretion disk (for weakly magnetized white dwarfs) or column (for magnetic white dwarfs).  A hot spot is expected to form where the gas stream from the donor star hits the outer edge of the accretion disk.  In a mass-exchanging binary, the orbital dependency of the flickering amplitude can be obtained by measuring the random brightness variations caused by flickering in an ensemble of light curves (i.e., the scatter curve) as a function of the orbital phase. Two different approaches exist for this purpose.  In its original form, the `ensemble' method defines the scatter curve as the root-mean-square (rms) deviations of the set of light curves with respect to the mean orbital light curve (Horne \\& Stiening 1985). This definition yields good estimates of the flickering amplitude when the gross features of the light curve are stable in time and all the variability is caused by flickering. In the presence of long term brightness variations (i.e., changes on scales longer than the orbital period), the rms curve overestimates the flickering amplitude because the scatter produced by flickering is mixed with that caused by the long term brightness changes (e.g., Welsh \\& Wood 1995). A critical investigation of this problem is presented in Bruch (2000). Bennie, Hilditch \\& Horne (1996) devised a way to disentangle the scatter induced by long term variations from that caused by flickering. They associated a `reference' out-of-eclipse flux level to each light curve in the data set. For a given phase bin, a plot of the flux of the data points versus the reference flux shows that these quantities are correlated. The slope of the correlation measures the long term change in the flux at that phase (a unity slope means that the flux at that phase varies in an 1:1 proportion with the reference flux) whereas the rms with respect to the best linear fit gives an unbiased measurement of the flickering component.  In the `single' method one quantifies the flickering in a set of light curves by measuring the scatter of each individual light curve with respect to a smoothed version of the light curve, and by combining the results from all light curves to derive a mean scatter curve (Bruch 1996). This is equivalent of a data filtering process, yielding the high-frequency components of the flickering. An intrinsic limitation of this method is the difficulty in separating the high-frequency flickering from rapid brightness changes, p.ex., caused by the eclipse, which may lead to the introduction of artifacts around sharp mid-ingress and mid-egress eclipse phases. A careful choice of the smoothing filter cut-off frequency is needed in order to avoid this problem.  The two techniques are complementary.  The `ensemble' method samples flickering at all frequencies. But because the power spectrum density of the flickering can be reasonably well described by a power-law $P(f) \\propto f^{-\\alpha}$, with $\\alpha\\simeq 1-3$ (Bruch 1992), an ensemble curve will normally be dominated by the low-frequency flickering components.   On the other hand, the `single' method produces curves which sample only the high-frequency flickering, with the cut-off frequency being defined by the smoothing filter width. Thus, the combination of both methods opens the possibility not only to probe the location of the flickering sources but also to separate low- and high-frequency components of the flickering. This may be very useful in the case where different flickering mechanisms lead to variability with distinct frequencies (e.g., see section~\\ref{results}).  V2051~Ophiuchi is a dwarf nova, a sub-class of CVs comprised by low-mass transfer binaries showing recurrent outbursts (caused by a sudden increase in the mass inflow in the disk) in which the brightness increases by factor of 5--100 on timescales of weeks to months (Warner 1995). V2051~Oph is distinguished among other ultra-short period eclipsing dwarf nova (i.e., Z~Cha, OY~Car and HT~Cas) by its remarkable flickering activity (amplitude of $\\ga 30$ per cent), which is responsible for a variety of eclipse morphologies and usually masks the orbital hump associated to anisotropic emission from the hot spot (e.g., Warner \\& Cropper 1983; Cook \\& Brunt 1983). Based on their light curves, Warner \\& Cropper (1983) concluded that the flickering in V2051~Oph arises mainly from the inner disk regions, with only a minor contribution from the hot spot. Baptista et~al. (1998a) caught the star in an exceptionally low brightness state during 1996 in which the flickering activity was mostly suppressed. They benefited from the clean view of the eclipses of the white dwarf and the faint hot spot at that epoch to derive basic parameters of the binary (e.g., masses, radii, orbital inclination).  Bruch (2000) applied the `single' method to a large set of light curves of V2051~Oph to find that the flickering curve shows an eclipse `coincident with the white dwarf eclipse and definitely narrower than the disk eclipse' as well as an orbital hump indicative that the bright spot contributes to the flickering in this star. However, the phase resolution of his scatter curve, $\\Delta\\phi=0.01$~cycles, is not enough to adequately resolve the shape of the eclipse.  In this paper we present flickering curves of V2051~Oph of high phase resolution ($\\Delta\\phi=0.002$~cycles) obtained with the `ensemble' and `single' methods, and we apply eclipse mapping techniques to these curves to map the flickering sources in this binary. The observations are described in section~\\ref{observa}. The data analysis is reported in section~\\ref{analysis}. The results are presented in section~\\ref{results} and discussed in section~\\ref{discuss}. A summary of the conclusions is presented in section~\\ref{conclui}. ",
        "conclusions": "\\label{discuss}  \\subsection{A comparison with previous results}  The eclipse in the flickering (`single') curve of Bruch (2000) is narrower and deeper than that of our `single' curves and led him to conclude that the flickering in V2051~Oph originates in the innermost disk regions.  This is in contrast with our results. Here we explore a possible explanation for this difference.  Bruch's (2000) smoothing filter cut-off frequency of 0.018~Hz implies that only flickering occurring on time scales shorter than 1 minute are included in his `single' curve, whereas our `single' curves comprise flickering of time scales shorter than 2.8 min.  Taking into account the power-law dependency of the flickering distribution, this means that his flickering curve samples typically higher frequencies than our 'single' curves. Both the disk turbulence and the magnetic reconnection flickering models predict a spatial segregation of flickering events, with higher frequency flickering occurring closer to disk center (Geertsema \\& Achterberg 1992; Kawaguchi et~al. 2000).  Thus, a `single' flickering curve sampling higher frequencies should lead to a narrower and deeper eclipse than a curve sampling lower frequencies as the distribution of the flickering sources in the former case is more concentrated towards the disk center. This could explain the observed difference in shape of `single' curves with different cut-off frequencies.  We tested this idea by creating 'single' curves using the same cut-off frequency of Bruch (2000).  Although the resulting curves are too noisy to allow drawing any firm conclusion, we found marginal evidence that the increase in cut-off frequency indeed lead to a narrower and deeper eclipse shape.  A larger data set (to increase the S/N ratio of the `single' curves) is required in order to provide a more conclusive test of this prediction.  It is worth mentioning that V2051~Oph was perceptibly brighter at the epoch sampled by Bruch's (2000) data (out-of-eclipse magnitude of $B\\simeq 15.0$~mag) than at the epoch of our observations ($B\\simeq 15.7-16.2$~mag).  We remark that if the flickering distribution (and the relative importance of each of the flickering sources) changes significantly with brightness level, the comparison made in this section may be misleading. In addition, we note that the flickering curve of Bruch (2000) was derived from a data set containing a mix of white light observations obtained with photomultipliers of different effective wavelengths (e.g., the `blue' Amperex 56\\,DVP vs. the `red' RCA 31034A) including both data in quiescence and in outburst (Warner \\& Cropper 1983; Warner \\& O'Donoghue 1987) and, therefore, should be looked with some caution.   \\subsection{Estimating the viscosity parameter $\\alpha$}  Geertsema \\& Achterberg (1992) investigated the effects of magneto-hydrodynamic (MHD) turbulence in an accretion disk. They found that the rate of energy dissipated per unit area at the disk surface shows large fluctuations which could be a source of flickering in CVs and x-ray binaries.  Here we use their model and the observed relative amplitude of the `single' flickering to estimate the viscosity parameter $\\alpha$ (Shakura \\& Sunyaev 1973) of the quiescent accretion disk in V2051~Oph.  In the model of Geertsema \\& Achterberg (1992), the rms value of the fluctuations $\\sigma(D)$ in the average rate of energy dissipated per unit area $\\langle D \\rangle$ is given by, \\begin{equation} \\frac{\\sigma(D)}{\\langle D \\rangle} \\simeq \\frac{2.5}{\\sqrt{N}} \\;\\;\\; , \\end{equation} where, \\begin{equation} D(r) = \\frac{3 G M_1 \\hbox{\\.{M}}}{8 \\pi r^3} \\left( 1 - \\sqrt{\\frac{r_1}{r}} \\right)  \\;\\;\\; , \\end{equation} $M_1$ and $r_1$ are, respectively, the mass and radius of the white dwarf, and $N$, the number of turbulent eddies that contribute to the local fluctuation, is given by, \\begin{equation} N(r) = 4 \\pi \\frac{r}{H} \\left( \\frac{H}{L} \\right)^2 \\simeq 1.25 \\times 10^3 \\left( \\frac{H}{L} \\right)^2   \\;\\;\\; , \\end{equation} where $H$ is the disk scale height, $L$ is the size of the largest turbulent eddies, and we assumed $H/r \\simeq 10^{-2}$, typical of thin disks (e.g., Frank, King \\& Raine 1992). The $\\alpha$-parameter of Shakura \\& Sunyaev (1973) can be written as, \\begin{equation} \\alpha \\simeq 0.9 \\left( \\frac{L}{H} \\right)^2  \\simeq 1130 \\; \\frac{1}{N} \\simeq 180 \\left[ \\frac{\\sigma(D)} {\\langle D \\rangle} \\right]^2 \\;\\; . \\label{eq.flick} \\end{equation}  If the disk flickering is caused by fluctuations in the energy dissipation rate induced by MHD turbulence, the relative amplitude of this flickering component is a direct measurement of the ratio $\\sigma(D)\\,/\\langle D \\rangle$. Replacing the observed relative amplitude of the `single' flickering in the above expressions we find $N\\simeq 7000$, $L/H \\simeq 0.42 \\; (10^2\\, H/r)^{-1/2}$, and a viscosity parameter $\\alpha \\simeq 0.16 \\; (10^2\\, H/r)^{-1}$. Because the relative amplitude of the 'single' flickering is constant with radius, the inferred viscosity parameter is the same at all disk radii.  The inferred viscosity parameter is significantly higher than expected for a quiescent, cool disk in the disk instability model ($\\alpha_{cool} \\simeq 10^{-2}$).  In fact, it is of the order of the expected viscosity of hot disks during dwarf novae outbursts ($\\alpha_{hot} \\simeq 0.1-0.2$, see e.g., Lasota 2001). Thus, if our assumption that MHD turbulence is the cause of the disk flickering is correct, the estimated viscosity parameter is uncomfortably high to be accommodated by the disk instability model.  One way to avoid this inconsistency is if the accretion disk of V2051~Oph is geometrically thick ($H/r \\ga 0.05$). However, there is a problem with this thick disk scenario. It would require a disk half-opening angle of $\\beta \\ga 3\\degr$. At the high inclination of V2051~Oph, this would lead to an apparent brightness contrast between the emission from the disk hemisphere closest to (the `front' side) and farthest from (the `back' side) the secondary star, with the `back' side being a factor of at least 2.5 brighter than the `front' side \\footnote{We note that this exercise does not take into account possible limb darkening effects, which would make the brightness ratio ever larger than assumed.}. There is no evidence of front-back brightness contrast in our steady-light eclipse maps.  An important implication of the derived high value of the viscosity parameter $\\alpha$ is that the quiescent disc of V2051~Oph should be in a steady-state (because the resulting viscous timescale, $\\simeq 30$ days, is short in comparison with the typical timescale between outbursts, $\\simeq 10^{3}$ days).  This is apparently in contradition with the flat radial temperature profile derived in section~\\ref{trad}. However, the optical spectrum of V2051~Oph is dominated by strong Balmer lines and a Balmer jump in emission, indicating important contribution from optically thin gas. In this case, the assumption that the B-band blackbody brightness temperature reflects the gas effective temperature probably fails and the comparison of the radial temperature distribution with theoretical disk models (Fig.~\\ref{fig7}) is possibly meaningless."
    },
    "0404/astro-ph0404057_arXiv.txt": {
        "abstract": "The proposed field-of-view of the {\\it Kepler} mission is at an ecliptic latitude of $\\sim55^\\circ$, where the surface density of scattered Kuiper Belt Objects (KBOs) is a few percent that in the ecliptic plane.  The rate of occultations of \\kep\\ target stars by scattered KBOs with radii $r\\ga10~\\km$ is $\\sim 10^{-6}-10^{-4}~{\\rm star^{-1}}~{\\rm yr^{-1}}$, where the uncertainty reflects the current ignorance of the thickness of the scattered KBO disk and the faint-end slope of their magnitude distribution.  These occultation events will last only $\\sim 0.1\\%$ of the planned $\\texp=15$ minute integration time, and thus will appear as single data points that deviate by tiny amounts.  However, given the target photometric accuracy of {\\it Kepler}, these deviations will nevertheless be highly significant, with typical signal-to-noise ratios of $\\sim 10$.  I estimate that $\\sim 1-20$ of the $10^5$ main-sequence stars in {\\it Kepler}'s field-of-view will exhibit detectable occultations during its four-year mission.  For unresolved events, the signal-to-noise of individual occultations scales as $\\texp^{-1/2}$, and the minimum detectable radius could be decreased by an order of magnitude to $\\sim 1~\\km$ by searching the individual $3~{\\rm sec}$ readouts for occultations. I propose a number of methods by which occultation events may be differentiated from systematic effects. {\\it Kepler} should measure or significantly constrain the frequency of highly-inclined, $\\sim 10~\\km$-sized KBOs. ",
        "introduction": "}  The Kuiper belt (see \\citealt{lj02} for a review) constitutes one of the few fossil records of the formation of planets and planetesimals in the early solar system.  It is therefore of great interest to determine the ensemble properties of its denizens.  These properties provide clues to the physical processes operating in the protoplanetary solar disk, as well as in the subsequent evolution of the solar system.  Kuiper Belt Objects (KBOs) are currently discovered in deep optical imaging surveys that are primarily limited to latitudes within $\\sim 10\\deg$ of the ecliptic or invariable plane (see, e.g., \\citealt{millis02}). The apparent $R$-magnitude of a KBO at a distance of $a\\simeq 40~\\au$ with radius $r \\simeq 100~\\km$ is $m_R\\simeq 23.4$, assuming an albedo of $4\\%$.  Since the flux of a KBO arises from reflected sunlight, it is proportional to $r^{2}$ and $a^{-4}$.  Therefore, imaging surveys are only able to detect relatively large and nearby KBOs with reasonable expenditure of resources. Furthermore, in order to have the best chance of success, they do not target high ecliptic latitudes, and are therefore biased against KBOs with large inclinations $i$.  Large KBOs are known to fall into at least three dynamical classes. Classical KBOs are relatively dynamically cold, with eccentricities $e\\la 0.3$ and $i\\la 5\\deg$.  Resonant KBOs are in mean motion resonances with Neptune, and were likely trapped there and subsequently excited to high $i$ and $e$ by the outward migration of Neptune \\citep{malhotra95}.  Scattered KBOs have very broad $e$ and $i$ distributions, and have probably been excited by some as-yet unknown process.  The size distribution of large KBOs is known to follow the form of $r^{-q}$, with $q\\simeq 4$, and a surface density of objects of $1~{\\rm deg^{-2}}$ at $r\\simeq 100~\\km$.  For $r\\la 100~\\km$, collisions will preferentially destroy KBOs, so that below a certain radius $r<\\rb\\la 100~\\km$, the size distribution is expected to flatten.  This break radius is set by the dynamical and physical properties of the KBOs, and is roughly where the collision time is equal to the age of the system.  The first estimate of this break radius, based on deep Hubble Space Telescope (HST) observations, yields $\\rb \\sim 70~\\km$ (\\citealt{bern04}, hereafter B04).  Unfortunately, it will be difficult, in the near future, to learn about the properties of small KBOs beyond what is already known using direct imaging, and especially so for those with high inclinations. Further progress will require next-generation synoptic surveys (Pan-STARRS, \\citealt{kaiser02}; The Discovery Channel Telescope, \\citealt{mds03}; LSST, \\citealt{tyson02}), or new methods.  One such method, first proposed by \\citet{bailey76}, is to search for stellar occultations by foreground KBOs.  This technique has been explored by several authors \\citep{dyson92,bw97,cf04}, and is being implemented by \\citet{roques03}, and the Taiwanese-American Occultation Survey (TAOS; \\citealt{liang03}). However, due to the low expected event rate, the first occultation surveys will likely focus on low-latitude fields.  Therefore, the statistics of small, high-inclination KBOs will remain meager.  Here I explore the possibility of measuring the surface density of small KBOs with high inclinations using the NASA space mission \\kep. \\kep's main goal is to detect transiting planets around normal stars, and it is far from ideally suited to detect occultations by foreground KBOs, especially given the high latitude ($55\\deg$) of its target field, and its long ($15~{\\rm min}$) exposure time. However, due to the large number of target stars and extraordinary photometric precision, I demonstrate that it may nonetheless detect a significant number of occultation events.  If these can be separated from systematic artifacts, {\\it Kepler} should be able to determine or severely constrain the frequency of small, high-inclination KBOs.  \\clearpage ",
        "conclusions": "\\label{sec:discussion}  The shaded area in Figure \\ref{fig:one} shows roughly the 95\\% confidence limit on the standard deviation $\\sigma_i$ of the inclination distribution from B01, and the faint-end slope $\\alpha_2$ from B04.  In this region, the predicted number of events varies from $\\ll 1$ to nearly $20$.  Marginalizing $\\ndet$ over $\\sigma_i$ and $\\alpha_2$, assuming Gaussian uncertainty distributions in both parameters with mean and dispersion of $20\\deg$ and $4\\deg$ for $\\sigma_i$, and $-0.5$ and $0.6$ for $\\alpha_2$, yields a best estimate of the number of expected events of $\\left< \\ndet \\right>\\sim 3$.  If no events are detected during the mission lifetime, a 95\\% confidence upper limit can be placed that should correspond to the $\\ndet=3$ contour, which would exclude roughly half the currently allowed region of $\\sigma_i$ and $\\alpha_2$.  Although the expected number of events in the 95\\% c.l.\\ region of $\\sigma_i$ and $\\alpha_2$ is relatively small, it is important to keep in mind that these parameters might be considerably in error, as all previous surveys have had relatively strong selection biases against scattered KBOs.  In particular, the constraint on $\\alpha_2$ from B04 is applicable to their `excited' sample, the majority of which are probably resonant KBOs. Therefore, the faint-end slope of scattered KBOs may differ considerably from that measured by B04; this might even be expected if scattered KBOs were excited early in the formation of the solar system.  For $\\alpha_2\\ga 0.4$, the number of expected events can be quite substantial.  In addition to uncertainties in the properties of the known scattered KBO population, there may be additional populations of KBOs at larger distances that are currently undetectable, but may be detected or constrained by \\kep.  For unresolved events, the signal-to-noise of occultations scales as $\\texp^{-1/2}$.  Therefore, the signal-to-noise of individual events can be increased by simply decreasing the \\kep\\ exposure time. In the break-limited regime ($\\alpha_2 \\la 0.4$), where small objects with $m_R>m_b$ do not contribute much to the optical depth, this does not significantly increase the number of detectable events, although it does improve the reliability of already detectable occultations with $m_R\\la m_b$. In the signal-to-noise limited regime $\\alpha_2 \\ga 0.4$, decreasing the exposure time can increase the number of detectable occultations dramatically.   As currently planned, \\kep\\ will only transmit integrations of $\\texp=15~{\\rm min}$, due primarily to bandwidth limitations. However, the detectors will in fact be read out every 3 seconds in order to avoid saturation by bright stars.  These $3~{\\rm sec}$ readouts will then be filtered and combined into $15~{\\rm min}$ integrations. It may be possible, using in-flight software, to search for occultations in the individual $3~{\\rm sec}$ exposures, and flag and transmit significant events.   With this possibility in mind, in Figure \\ref{fig:one}(b) I show $\\ndet$ as a function of $\\sigma_i$ and $\\alpha_2$ for $\\texp=3~{\\rm sec}$.  As expected, deep in the break-limited regime ($\\alpha_2\\la 0$), the number of detectable occultations is similar for $\\texp=3~{\\rm sec}$ and $\\texp=15~{\\rm min}$. However, for $\\alpha_2\\ga 0$, the number of expected event increases significantly for $\\texp=3~{\\rm sec}$.  In the allowed region of parameter space, the number of events can be as large as $\\ndet\\sim 600$.  Marginalizing over $\\sigma_i$ and $\\alpha_2$ yields $\\left< \\ndet \\right> \\sim 12$.  For $\\alpha_2=0.6$, and $\\sigma_i =20\\deg$, $\\ndet \\simeq 10^3$.  Whether or not it is efficient or feasible to search the individual $3~{\\rm sec}$ integrations for occultations remains to be seen. However, given the large potential increase in the number of expected events, and the fact that the subsequent constraints will be hard to acquire with other methods, the possibility should be considered.  Given the fact that the occultation events are unresolved, and have a small amplitude, systematic effects that could mimic occultation events are a serious concern.  For a large sample of candidate events, there are some signatures that may differentiate between real occultations and systematics.  Occultation events should show a factor of $\\sim 2$ gradient in the optical depth over the $\\sim 10\\deg$ in ecliptic latitude spanned by the \\kep\\ FOV.  In addition, the relative velocity also varies over the course of the year; for the latitude of the \\kep\\ FOV, this variation is $10~\\kms$ peak-to-peak, which corresponds to a variation in both the signal-to-noise and event rate of amplitude $\\sim 30\\%$.  Unfortunately, given the intrinsic velocity dispersion of KBOs, and the variation in signal-to-noise noise due to non-equatorial occultations, it will be difficult to measure these deviations unless there are a large number of events. For individual events, it may be possible to distinguish between occultations and systematic effects based on the detailed shape and centroid of the stellar point-spread function for high signal-to-noise events \\citep{jd03}. Ground-based follow-up observations can be used to eliminate contamination from background stars that happen to be blended with the primary star.  It may also be possible to directly detect the KBOs causing occultation events by imaging them in reflected light. Provided that the data are analyzed in real time, candidate occultation events could be alerted within a day of the event.  The signal-to-noise of the occultations should be correlated with their size (e.g., Eq.\\ \\ref{eqn:snexp}), and therefore brighter events can be preferentially selected for follow-up.  The brighter candidates will have $m_R \\sim m_B \\simeq 26~{\\rm mag}$, within the range of large, ground-based telescopes.  Typical proper motions are a few arcseconds per hour.  Therefore, deep imaging taken at several epochs over a few days after the occultation should reveal a faint object moving away from the bright target star.  Provided that systematic effects can be controlled, and a reliable sample of occultation events can be acquired, then the observed distribution of signal amplitudes can be used to statistically infer the distribution of event time scales.  This can then be converted to an estimate of the optical depth.  If no events are detected, then it should be possible to place limits on the surface density of KBOs in the \\kep\\ FOV.  Therefore, \\kep\\ should measure or significantly constrain the number of highly-inclined, $\\sim 10~\\km$-sized KBOs."
    },
    "0404/astro-ph0404138.txt": {
        "abstract": "The Virgo Cluster is the dominant mass concentration in the Local Supercluster and the largest collection of elliptical and lenticular galaxies in the nearby universe. In this paper, we present an introduction to the ACS Virgo Cluster Survey: a program to image, in the F475W and F850LP bandpasses ($\\approx$ Sloan $g$ and $z$), 100 early-type galaxies in the Virgo Cluster using the {\\it Advanced Camera for Surveys} on the {\\it Hubble Space Telescope}. We describe the selection of the program galaxies and their ensemble properties, the choice of filters, the field placement and orientation, the limiting magnitudes of the survey, coordinated parallel observations of 100 ``intergalactic\" fields with WFPC2, and supporting ground-based spectroscopic observations of the program galaxies. In terms of depth, spatial resolution, sample size and homogeneity, this represents the most comprehensive imaging survey to date of early-type galaxies in a cluster environment. We briefly describe the main scientific goals of the survey which include the measurement of luminosities, metallicities, ages, and structural parameters for the many thousands of globular clusters associated with these galaxies, a high-resolution isophotal analysis of galaxies spanning a factor of $\\sim$ 450 in luminosity and sharing a common environment, the measurement of accurate distances for the full sample of galaxies using the method of surface brightness fluctuations, and a determination of the three-dimensional structure of Virgo itself. ",
        "introduction": "\\label{sec:introduction}  Clusters of galaxies are among the largest gravitationally bound structures in the universe. Although they contain a small fraction of all galaxies in the universe ($e.g.$, Turner \\& Tyson 1999), they have long provided invaluable constraints on structure formation and galaxy evolution. During the past few decades, it has become abundantly clear that the formation of galaxy clusters is a highly protracted process: most clusters show some degree of substructure (Geller \\& Beers 1982; Dressler \\& Shectman 1988;  Fitchett \\& Merritt 1988; Mohr, Fabricant \\& Geller 1993; Bird 1994) and there is mounting evidence that, even among nearby clusters, the process of virialization may not be fully complete ($e.g.$, Adami, Biviano \\& Mazure 1998; Drinkwater, Gregg \\& Colless 2001).  Rich clusters may contain several thousands of member galaxies within a region $\\sim$ 1 Mpc in diameter. Collectively, these galaxies usually comprise only a few percent of the total gravitating mass ($e.g.$, Carlberg et al. 1996). Instead, cluster masses are dominated by dark matter, and to a lesser extent, by the X-ray emitting gas which fills their gravitational potential wells (typically $\\gtrsim 90$\\% and 10--30\\% of the total mass, respectively; B\\\"ohringer 1995). Yet despite their relatively minor contribution to the cluster mass budget, and to that of the universe in general, much of our understanding of galaxy formation and evolution is based on observations of galaxies in cluster environments.  Virgo is probably the most thoroughly studied cluster of galaxies, being the cluster nearest to our own Galaxy. In most respects, its properties are typical of large clusters. It is a relatively populous system (Abell richness class I; Girardi et al. 1995), consisting of $\\approx$ 2000 cataloged members brighter than $B_T = 18$ (Binggeli, Sandage \\& Tammann 1987). Mass estimates vary substantially ($e.g.$, 0.15--1.5$\\times10^{15}M_{\\odot}$; B\\\"ohringer et al. 1994; Schindler, Binggeli \\& B\\\"ohringer 1999; McLaughlin 1999a; Tonry et al. 2000; Fouqu\\'e et al.  2001), but generally fall within the range typical of rich clusters (B\\\"ohringer 1995). It contains vast quantities of X-ray emitting gas (Takano et al. 1989; B\\\"ohringer et al. 1994;  Schindler et al. 1999), and shows clear evidence for both substructure and non-virialized motions ($e.g.,$ Huchra 1985; Binggeli et~al. 1987; Binggeli, Popescu \\& Tammann 1993). Indeed, the property which sets Virgo aside from other clusters is proximity --- at a distance of $\\approx 17$ Mpc (Tonry et al. 2001), it may be studied in a level of detail which will never be possible with more distant systems.  The discovery of the cluster itself dates back to Messier and M\\'echain, who noted that 13 of the entries in Messier's catalog lie in the constellation of Virgo (Messier 1781).\\footnote{Messier writes ``The constellation Virgo and especially the northern wing is one of the constellations which encloses the most nebulae. This catalog contains 13 which have been determined, viz. Nos. 49, 58, 59, 60, 61, 84, 85, 86, 87, 88, 89, 90 and 91.... Most of these nebulae have been pointed out to me by M. M\\'echain.\" It is now recognized that three additional galaxies from Messier's catalog (M98, M99 and M100) are also Virgo members.} The first systematic studies of the cluster were those of Shapley \\& Ames (1926; 1929a-e), Hubble \\& Humason (1931) and Smith (1936). Subsequent studies by Zwicky (1942; 1957) and Holmberg (1958) produced both a better understanding of its spatial structure and a more complete census of cluster members. The cluster richness and its large dynamical mass led de Vaucouleurs (1956; 1961) to identify it as the nucleus of the Local Supercluster. Surveys of the Virgo Cluster culminated in the 1980s, with the heroic efforts of Bruno Binggeli, Allan Sandage and Gustav Tammann ($e.g.,$ Binggeli, Sandage \\& Tammann 1984; 1985; 1987; Sandage, Binggeli \\& Tammann 1985). Their Virgo Cluster Catalog (VCC) of 2096 galaxies, based on deep photographic images covering an area of $\\approx$ 140 deg$^{2}$, still stands as the most complete census of Virgo galaxies brighter than $B_T = 18$.  Continuing studies of Virgo and its constituent galaxies at X-ray (Schindler et al. 1999; B\\\"ohringer et al. 2001; Allen et al. 2001), ultraviolet (Lieu et al. 1996; Cortese et al. 2003), optical (Trentham \\& Hodgkin 2002; Sabatini et al. 2003), infrared (Boselli et al. 1997; Tuffs et al. 2002), and radio (van Driel et al. 2000; Solanes et al. 2002) wavelengths --- as well as combinations thereof (Gavazzi et al. 2002; 2003; Boselli, Gavazzi \\& Sanvito 2003) --- have largely confirmed the view of Virgo as an environment that is, in most respects, typical of rich clusters (see, $e.g.$, Binggeli 1999 and references therein). It therefore seems likely that Virgo will continue to play a central role in the study of galaxy formation and evolution. This is particularly true in the case of early-type galaxies since Virgo contains, by far, the largest collection of early-type giant and dwarf galaxies in the local universe.  Not surprisingly, this enormous reservoir of galaxies has been a favorite target of the {\\it Hubble Space Telescope} (HST) since its launch in 1990. To date, thousands of high-resolution images for fields in Virgo have been obtained with FOC, WFPC, NICMOS, STIS and, especially, WFPC2. The scientific output from the first decade of Virgo research with HST touches upon a remarkably diverse range of astrophysical topics including the extragalactic distance scale, accreting and pulsating stars, the cores of spiral and early-type galaxies, the stellar populations of galaxies, extragalactic globular cluster systems, dwarf galaxies, the nuclear and circumstellar structure of active galaxies, intergalactic stars and diffuse light, the ultraviolet/optical/IR morphologies of nearby galaxies, galaxy collisions and interactions, and pre- and post-main-sequence stages of stellar evolution.  With the installation of the {\\it Advanced Camera for Surveys} (ACS; Ford et al. 1998) during Servicing Mission 3B in March 2002, the imaging capabilities of HST have improved dramatically ($i.e.$, by nearly an order of magnitude in discovery efficiency). Relative to WFPC2, the Wide Field Channel (WFC) mode on ACS boasts a factor of $\\sim$ two increase in imaging area and a typical gain of 3--4 times in throughput over the wavelength range 0.4--1.0~$\\mu$m. ACS thus opens the possibility of obtaining deep, high-resolution optical images for a large sample of Virgo Cluster galaxies in a relatively modest allocation of observing time.  In this paper, the first in a series, we present an introduction to such a program: the ACS Virgo Cluster Survey. The program, which was awarded 100 orbits in Cycle 11, aims to obtain deep F475W ($\\approx$ Sloan $g$) and F850LP ($\\approx$ Sloan $z$) ACS/WFC images for 100 early-type galaxies that are confirmed members of the Virgo Cluster. The paper is organized as follows: we present a brief description of the primary scientific goals of the survey and the anticipated data products in \\S\\ref{sec:motivation}, the selection of sample galaxies in \\S\\ref{sec:properties}, and the adopted observing strategy, including choice of filters, exposure times, limiting magnitudes and field placement, in \\S\\ref{sec:strategy}. Since the survey also includes supplemental WFPC2 parallel observations of 100 intergalactic fields in Virgo and ground-based, long-slit spectroscopy of the program galaxies, brief descriptions of these observations are given in \\S\\ref{sec:parallel} and \\S\\ref{sec:spectroscopy}, respectively. Scientific results from the survey will be presented in future papers in this series. ",
        "conclusions": "\\label{sec:summary}  This paper has provided a brief introduction to the ACS Virgo Cluster Survey (GO-9401). In terms of depth, spatial resolution, sample size and homogeneity, this represents the most comprehensive imaging survey to date of early-type galaxies in a cluster environment. Scientific results from the survey will be presented in future papers in this series. More information on the survey, including object catalogs and other data products, are available from the program website: {\\sf http://www.physics.rutgers.edu/$\\sim$pcote/acs}."
    },
    "0404/astro-ph0404043_arXiv.txt": {
        "abstract": "We consider Double Quintessence models for which the Dark Energy sector consists of two coupled scalar fields. We study in particular the possibility to have a transient acceleration in these models. In both Double Quintessence models studied here, it is shown that if acceleration occurs, it is necessarily transient. We consider also the possibility to have transient acceleration in two one-field models, the Albrecht-Skordis model and the pure exponential. Using separate conservative constraints (marginalizing over the other parameters) on the effective equation of state $w_{eff}$, the relative density of the Dark Energy $\\Omega_{Q,0}$ and the present age of the universe, we construct scenarios with a transient acceleration that has already ended at the present time, and even with no acceleration at all, but a less conservative analysis using the CMB data rules out the last possibility. The scenario with a transient acceleration ended by today, can be implemented for the range of cosmological parameters $\\Omega_{m,0}\\gtrsim 0.35$ and $h\\lesssim 0.68$. ",
        "introduction": "The release of type Ia Supernovae data independently by two groups, the Supernovae Cosmology Project and the High-z Survey Project \\cite{SN1}, confirmed in more recent work \\cite{SN2}, indicating that our Universe might be presently accelerating, has profound implications on the current paradigm in cosmology. It was found that, assuming flatness, the best-fit Universe with a cosmological constant $\\Lambda$ is given by the set of cosmological relative densities $(\\Omega_{\\Lambda,0},\\Omega_{m,0})=(0.72, 0.28)$. Here, $m$ stands for usual pressureless matter including (cold) dark matter and baryonic matter. These data are also in surprising agreement with the location of the first acoustic (Doppler) peak of the Cosmological Microwave Background temperature anisotropy multipoles. If these data are confirmed in the future, they imply a radical departure from usual textbooks Friedman-Lema\\^{\\i}tre-Robertson-Walker (FLRW) cosmology \\cite{SSPR}. Indeed, a perfect isotropic fluid cannot lead to accelerated expansion unless it has a sufficiently negative pressure while the data seem to imply that such a kind of smoothly distributed matter, called Dark Energy, constitute about two thirds of the whole matter budget of our Universe.  Obviously, a pure comological constant $\\Lambda$ could be responsible for this acceleration. However, its amplitude has to be exceedingly small, about $123$ orders of magnitude too small in order to be explained in a ``natural'' way. While this possibility is actually in good agreement with observations and attractive as it seems to make many pieces of our present understanding of structure formation fit into a consistent picture, in view of the above mentioned theoretical problems research has focused on other, more elaborate, Dark Energy candidates that would mimic a pure cosmological constant.  A $\\Lambda$-term is equivalent to a perfect isotropic fluid with constant equation of state $p_{\\Lambda}= -\\rho_{\\Lambda}$, equivalently $w_{\\Lambda}=-1$. Provided enough of the total energy density is stored in this component, the expansion will be accelerated. A first generalization could be some phenomenological scaling Dark Energy with constant equation of state and $w<-\\frac{1}{3}$. Some interesting insight can be gained by considering most of the universe energy stored in such a component (see e.g. \\cite{CPol}).  Another more elaborate alternative to a pure cosmological constant $\\Lambda$ is some effective, slowling varying, cosmological constant term that will start driving the universe expansion at low redshifts. The prominent candidate in this respect is some minimally coupled scalar field $\\Phi$, often termed quintessence, slowly rolling down its potential so that it has a negative pressure, $p_{\\Phi}=w_{\\Phi}\\rho_{\\Phi}$, $w_{\\Phi}<-\\frac{1}{3}$ \\cite{RP,Wetterich,FJ,CDS}. Later on tracking solutions were introduced \\cite{ZWS}, a substantial improvement with respect to the initial conditions problem, however without solving the cosmic coincidence problem. Actually, this is precisely the mechanism which drives the inflationary stage in the Early Universe and here too, one can consider several potentials and investigate how well they fit the observational data. If a scenario of this kind is the correct one, then it is possible in principle to reconstruct (the relevant part of) its potential $V(\\Phi)$ and the corresponding equation of state characterized by $w_{\\Phi}$, using luminosity distance measurements in function of redshift $z$. This procedure that can be extended to more elaborate models of gravity like so-called generalized, or extended, quintessence models in the framework of scalar-tensor theories of gravity \\cite{BEPS,EP}. Note that the latter case can represent a possible physical realization of so-called phantom energy, with $w <- 1$ \\cite{BEPS} and it has been intensively investigated (see e.g. \\cite{PM03}). One can also consider other modifications to the theory of gravity or to the coupling of scalar fields (see e.g. \\cite{GG}).  A Chaplygin gas is another example of more exotic Dark Energy candidates that was proposed, in the course of the universe expansion, this gas undergoes a transition from a dust-like to a cosmological constant-like equation of state \\cite{Chaplygin}.  Actually, as the Dark Energy sector is still unknown, it is interesting to explore all possible models that can pass successfully the observational constraints and we want to investigate Double Quintessence models for which this sector consists of two coupled scalar fields. Two scalar fields models give rise to exchange of kinetic energy between the two fields and produce a non-trivial time evolution of the Dark Energy equation of state. Some models of two fields quintessence have been introduced and studied in the literature \\cite{Halyo,Fujii,MPR,BBS}. We show with simple two-fields models how the introduction of an auxiliary field can bring accelerated expansion of our Universe to an end. Though our models have natural parameters, the cosmological coincidence however is not solved here. We note that many-fields inflationary models can have signatures \\cite{PS95} which distinguish them from single-field inflationary models, in the first place the possibility to have a characteristic scale in the primordial fluctuations spectrum (see e.g. \\cite{PS92}). In this respect, the situation is here more contrasted and we will return to this point in the Conclusion.  The existence of an event horizon in the case of a de Sitter phase is an obstacle to the implementation of string theory because the S-matrix formulation is no longer possible, hence eternal acceleration, leading to a de Sitter space in the asymptotic future is problematic \\cite{HKS}. Therefore, transient acceleration of our universe is certainly a wellcome feature in this respect. Even more, we will show that we can have, in a fairly natural way, a transient acceleration which has already ended at the present time and still in accordance with the observations. We explore the window in parameter space for which this intriguing scenario is realized. It is even possible, though this is a marginal possibility, to construct scenarios where there is no acceleration at all, neither at the present time nor in the past of our Universe. Of course, agreement with the observational data still require that the Dark Energy sector dominate from some time on and has an effective equation of state satisfying $w<0$. We will see that in our models the quantity $w$, which will be called eos (equation of state) parameter, is not constant and actually strongly varying on the redshift interval $0<z<2$ relevant for (past) luminosity distance measurements. Actually, it was already found earlier that luminosity-distance measurements are rather insensitive to large variations of the eos parameter at low redshifts \\cite{MBS,CPol}. This is why, very interestingly, in our Double Quintessence models, it is possible to construct scenarios where the standard interpretation of a presently accelerating Universe can be challenged.  In Section II, we give the basic equations of our system and introduce the relevant quantities and notations. In Section III, we list the observational constraints on our models and the associated quantities. In Section IV, our models are presented and investigated numerically in Section V and regions in parameter space leading to viable models are explored. We will take $8\\pi G=M_p^{-2}=1$ (where $M_p$ is the reduced Planck mass) and $\\hbar=c=1$ in the following. ",
        "conclusions": "Recent observational data suggest that some unknown component, called Dark Energy, contribute to about two thirds of the present total energy density filling our universe. Though the accuracy of the existing observations allow to constrain already at the present stage to some extent the possible Dark Energy candidates, a very large number of models are still permitted. It is therefore interesting to investigate all possible scenarios and we have also investigated in this work specific models where the Dark Energy sector is made of two coupled scalar fields. We have studied numerically two one-field models, the pure exponential in Section V.A and the AS model in Section V.B, and two Double Quintessence models, which can be seen as extensions of the corresponding one-field model, in Section V.C and V.D. We were interested in particular in scenarios for which the (recent) stage of accelerated expansion is transient.  Constraining the models with the observations, we have found the allowed window in the corresponding parameter space of each model. Investigation of these models has revealed that the following three possibilities can arise: the present acceleration is transient and still going on; some accelerated expansion did take place in the recent past but is already finished by today; finally no acceleration at all, this latter possibility being marginal. The two first scenarios, i.e. a transient acceleration either finished or not by today, can be obtained in all four models. A scenario with no acceleration at all is obtained only for the two one-field models V.A and V.D provided the cosmological parameters take their values at the edge of the allowed range, $\\Omega_{m,0}\\gtrsim 0.4$ and $h\\lesssim 0.65$. We summarize the possible dynamics of our models in Table \\ref{TS}.  In both Double Quintessence models studied in Sections V.C and V.D, for the allowed window in parameter space where accelerated expansion takes place, the dynamics of the model is such that the acceleration is necessarily transient, hence for these two models acceleration if it takes place will eventually come to an end either before today or in the future. In the model of Section V.C, in addition to the tracking field $\\Phi$, an auxiliary field $\\Psi$ is introduced which controls the presence of a minimum for the $\\Phi$ field and the duration of the Universe transient acceleration. In the model of Section V.D, it is the auxiliary field $\\Psi$, and hence $\\lambda_{eff}$, which induces accelerated expansion when it reaches its minimum. In contrast to our Double Quintessence models where acceleration is necessarily transient, for the one-field models we have studied the acceleration is typically eternal and covers most of the allowed parameter space. In view of the theoretical problems posed by eternal acceleration, all viable scenarios with transient acceleration constitute a wellcome alternative. On the other hand, the cosmological coincidence is not solved here and requires some amount of fine tuning on one of the free parameters. We insist that all the scenarios studied here are in agreement with observations, in particular with the Hubble diagram $H(z)$, or the luminosity distances $d_L(z)$ as a function of redshift, as reconstructed from the Supernovae data leading to the possible interpretation of a flat universe with $\\Omega_{\\Lambda,0}\\simeq 0.72$.  As mentioned in the Introduction, it is well known that many-fields inflationary models can produce primordial fluctiations spectra with a characteristic scale. The question arises naturally whether similar effects can be produced here. In fact, the luminosity distances can exhibit a characteristic scale if the equation of state of dark energy undergoes a phase-transition. Such cases were considered in \\cite{BKSU}, where the quantity $w(z)$ was taken with a step-like structure. We have checked numerically that such models could indeed exhibit a characteristic scale in their luminosity distances. In our Double Quintessence models where accelerated expansion is ended at the present time, we have also a large variation of $w(z)$ at low redshifts, typically from $w(z)\\simeq -1$ up to some higher value. However, as can be seen from Figures \\ref{eosAz} and \\ref{eosqB} this variation takes place at very low redshifts $0\\leq z\\lesssim 0.3$ and, though significant, this variation is not sharp enough. This is the reason why no characteristic scale is seen in the corresponding luminosity distances. Actually, it was already emphasized that the luminosity distances are not very sensitive to large, smooth, variations of the equation of state, or equivalently of the quantity $w(z)$, at low redshifts $z$. Results obtained here are a particular illustration of this property.  In the light of our results, it is clear that eternal acceleration could be challenged in two ways, either the acceleration is transient and will end at some time in the future, either it has already ended by today. Results obtained here show that the second possibility must be taken seriously if the observations allow a rather high matter content today, viz. $\\Omega_{m,0}\\gtrsim 0.35$ and also large variations of the eos parameter $w(z)$ at low redshifts. In particular, if it turns out that the data require $w_{Q,0}\\lesssim -0.5$, this possibility is ruled out."
    },
    "0404/astro-ph0404275_arXiv.txt": {
        "abstract": "It is well known that the clustering of galaxies depends on galaxy type. Such {\\it relative bias\\/} complicates the inference of cosmological parameters from galaxy redshift surveys, and is a challenge to theories of galaxy formation and evolution. In this paper we perform a joint counts-in-cells analysis on galaxies in the 2dF Galaxy Redshift Survey, classified by both colour and spectral type, $\\eta$, as early or late type galaxies. We fit three different models of relative bias to the joint probability distribution of the cell counts, assuming Poisson sampling of the galaxy density field. We investigate the nonlinearity and stochasticity of the relative bias, with cubical cells of side $10\\Mpc \\leq L \\leq 45\\Mpc$ ($h=0.7$). Exact linear bias is ruled out with high significance on all scales. Power law bias gives a better fit, but likelihood ratios prefer a bivariate lognormal distribution, with a non-zero `stochasticity' -- i.e. scatter that may result from physical effects on galaxy formation other than those from the local density field. Using this model, we measure a correlation coefficient in log-density space ($r_{\\LN}$) of 0.958 for cells of length $L=10\\Mpc$, increasing to 0.970 by $L=45\\Mpc$. This corresponds to a stochasticity $\\smash{\\sigma_b/\\hat{b}}$ of $0.44\\pm0.02$ and $0.27\\pm0.05$ respectively. For smaller cells, the Poisson sampled lognormal distribution presents an increasingly poor fit to the data, especially with regard to the fraction of completely empty cells. We compare these trends with the predictions of semianalytic galaxy formation models: these match the data well in terms of overall level of stochasticity, variation with scale, and fraction of empty cells. ",
        "introduction": "The question of whether galaxies trace the matter distribution of the universe has many implications for cosmology and galaxy formation theories. Since \\citet{1931ApJ....74...43H} it has been known that galaxies of different type cluster differently, and as such it cannot be possible for all galaxies to trace the matter distribution exactly. This observation has been reconfirmed many times, traditionally by comparisons of the correlation functions of different subgroups. For example, early type (or passive) galaxies are more strongly clustered than late type (or actively  star-forming) galaxies (e.g. \\citealt{1976ApJ...208...13D,1980ApJ...236..351D}; \\citealt*{1990ApJ...350..119L}; \\citealt{1996MNRAS.283..709H,2002MNRAS.332..827N,2002ApJ...571..172Z,2003MNRAS.344..847M}) and luminous galaxies cluster more strongly than faint galaxies (e.g. \\citealt*{1998AJ....115..869W}; \\citealt{2001MNRAS.328...64N,2002MNRAS.332..827N,2002ApJ...571..172Z,2004ApJ...608...16Z}).  Any difference in the distribution of galaxies relative to mass has become known as {\\it galaxy bias\\/}. This assumed a central importance in cosmology via the attempts to rescue the $\\Omega_m = 1$ universe after observations of cluster mass-to-light ratios suggested values closer to $\\Omega_m = 0.2$. Such bias could occur if the galaxy formation efficiency were increased in overdense regions of space, the so called `high peak scenario' \\citep{1985ApJ...292..371D,1986ApJ...304...15B}. Although these efforts ultimately proved fruitless, understanding of bias remains important. In recent years much effort has been put into investigating galaxy bias through theory and numerical modelling, while observational results have been restricted by small survey volumes. With the advent of large galaxy redshift surveys such as the 2dF Galaxy Redshift Survey (2dFGRS: Colless et al. 2001, 2003)\\nocite{2001MNRAS.328.1039C}, the Sloan Digital Sky Survey (SDSS: Strauss et al. 2002) \\nocite{2002AJ....124.1810S}and the Deep Extragalactic Evolutionary Probe (DEEP: Davis et al. 2002) \\nocite{2003SPIE.4834..161D}it is becoming possible to quantify the galaxy distribution as never before, and provide detailed descriptions with which to compare theoretical and numerical models.  In principle the form of bias should be derivable from the fundamental physical processes involved in galaxy formation; until we understand these, bias remains a description of our ignorance. The simplest model of galaxy biasing is the linear biasing model: $\\delta_g({\\bf x}) = b\\delta_m({\\bf x})$ where $\\delta_g$ is the galaxy overdensity perturbation, $\\delta_m$ the mass overdensity perturbation and $b$ a constant bias parameter. This model is unphysical for $b>1$, as it allows negative densities. Alternative models in the literature fall into several basic classes: linear or non-linear, local or non-local, deterministic or stochastic. Locally biased galaxy formation \\citep[e.g.][]{1993MNRAS.262.1065C,1998ApJ...504..607S,1993ApJ...413..447F} depends only on the properties of the local environment, and the galaxy density is assumed to be a universal function of the matter density: \\be \\delta_g = f(\\delta_m). \\ee Because galaxies are discrete objects, this prescription is normally supplemented by the {\\it Poisson Clustering Hypothesis\\/}, in which galaxies are modelled as random events, whose expectation number density is specified via $\\delta_g$. This model for discreteness can only be an approximation, but there is no simple alternative. We therefore assume Poisson sampling in what follows; for consistency, theoretical predictions are treated in the same way as the real data.  Non-local models \\citep[e.g.][]{1993ApJ...405..403B,1999ApJ...525..543M} arise when the efficiency of galaxy formation is modulated over scales larger than those over which the matter moves, for example by effects of quasar radiation on star formation. Stochastic bias (Pen 1998\\nocite{1998ApJ...504..601P}, Dekel \\& Lahav 1999\\nocite{1999ApJ...520...24D}, hereafter DL99) allows a range of values of $\\delta_g$ for a given $\\delta_m$, above the Poissonian scatter caused by galaxy discreteness. Stochasticity is a natural part of non-local models, but some stochasticity is always expected to arise from physical processes of galaxy formation \\citep{1999ApJ...522..590B}. Throughout this paper we follow the general framework for non-linear stochastic biasing of DL99, in which  the overdensity of one field can be related to that of a second field contained in the same volume of space through \\be \\delta_1 = b(\\delta_2) \\delta_2 + \\epsilon. \\ee The scatter (stochasticity) in the relation is given by \\be \\epsilon \\equiv \\delta_1 - \\langle\\delta_1\\rangle. \\ee In principle the bias parameter $b$ can be any function of $\\delta_2$; a constant value of $b$ and $\\epsilon=0$ represents deterministic linear biasing.  Galaxy bias is clearly of astrophysical interest in relation to an understanding of galaxy formation and evolution. Bias is also a major practical source of uncertainty in deriving cosmological constraints from galaxy surveys. Some particular examples are the measurement of $\\beta = \\Omega_m^{0.6} / b$ \\citep{2001Natur.410..169P,2003MNRAS.346...78H}, where DL99 showed that stochastic effects could explain large discrepancies between results from different methods \\citep[for a review see][Table 7.2]{1999fsu..conf.....D}. Power spectrum measurements require constant bias as a fundamental assumption \\citep{2001MNRAS.327.1297P}, and constraints placed on neutrino mass also assume scale independent biasing \\citep{2003JCAP...04..004E}. \\citet{1998ApJ...504..601P} calculate the effect of nonlinear stochastic bias on the measurement of the galaxy power spectrum on large scales, showing how the galaxy variance, bias and galaxy-dark matter cross correlation coefficient can be calculated from velocity distortions in the power spectrum. The importance of biasing has increased still further with the release of the WMAP first year results \\citep{2003ApJS..148..175S}. In order to combine CMB and 2dFGRS data to give tighter constraints on cosmological parameters, a model for galaxy bias is required \\citep{2003ApJS..148..195V}.  Three independent methods have been used to investigate galaxy biasing in the 2dFGRS catalogue. \\citet{2002MNRAS.333..961L} combined pre-WMAP CMB and 2dFGRS datasets to measure the average bias over a range $0.02 < k< 0.15\\, h\\,{\\rm Mpc}^{-1}$, concluding that galaxies are almost exactly unbiased on these scales.  \\citet{2002MNRAS.335..432V} found the bias parameter to be consistent with unity over scales $0.1<k<0.5\\,h\\,{\\rm Mpc}^{-1}$ through measurements of the 2dFGRS bispectrum.  A more direct method of studying the relation between mass and light is to map the dark matter using gravitational lensing. This field has made great progress in recent years, and it has been possible to measure not only the absolute degree of bias, but also its nonlinearity and stochasticity \\citep{2000AJ....120.1198F,2002ApJ...577..604H,2003ApJ...594...33F,2003MNRAS.346..994P}. For example, \\citet{2002ApJ...577..604H} combine the Red-Sequence Cluster Survey and VIRMOS-DESCART survey to find an average bias $b=0.71$ and linear correlation coefficient of $r \\simeq 0.57$ on scales of $1-2\\mpcoh$. However, current weak lensing measurements are dominated by non-linear and quasi-linear scales in the power spectrum, and it is not yet possible to say a great deal about bias in the very large-scale linear regime. This is of course the critical region for the interpretation of redshift surveys, where we want to know the relation between the power spectra of mass and light on $>100\\Mpc$ scales.  This question will be settled by future weak lensing surveys. In the meantime, we can address a related simpler problem: the {\\it relative\\/} bias between subsets of galaxies. The morphological differences between galaxies and the link to their environments has been discussed for many decades as a potential clue to the nature and evolution of galaxy clustering \\citep[e.g.][]{1951ApJ...113..413S,1972ApJ...176....1G,1976ApJ...208...13D,2001ApJ...558..520Y}. Modern galaxy redshift surveys allow us to split the galaxy population into a variety of subdivisions such as spectral type, colour and surface brightness. We can look at relative bias as a function of scale, and weighted by luminosity. This should yield important insights into the absolute degree of bias that may exist. \\citet{2001MNRAS.328...64N} measured bias as a function of luminosity in the 2dFGRS, finding a bias relative to $L^*$ galaxies of  $b/b^* = 0.85+ 0.15L/L^*$. We concentrate on the natural bimodality of the galaxy population, between red early types with little active star formation, and the blue late type population \\citep[e.g.][]{2003IAUS..216E.201B}. Lahav \\& Saslaw (1992) measured bias as a function of morphological type and scale using the UGC, ESO and IRAS catalogues. The Las Campanas Redshift Survey has already provided some observational evidence against the linear deterministic model from splitting galaxies by their spectral types \\citep{1999ApJ...518L..69T,2000ApJ...544...63B}, and we present here a more extensive analysis of this type.  There are several complementary methods for the measurement of galaxy clustering, although most previous studies of the relative bias between galaxy types have concentrated on a relative bias parameter defined as the square root of the ratio of the correlation functions for the types under study. \\citet{2003MNRAS.344..847M} used this method to measure the relative bias in the 2dFGRS, finding $b$(passive/active) ranging from about 2.5 to 1.2 on scales $0.2\\mpcoh < r < 20\\mpcoh$. However, even within such a large survey as the 2dFGRS the correlation functions become noisy beyond about $10\\mpcoh$. A second method is counts-in-cells, which can be directly related to the correlation function \\citep{1980lssu.book.....P}, and is optimised for the study of larger scales. It is this latter method that we employ in this paper. Conway et al. (2004) have also investigated the relative bias of different galaxy types using a counts-in-cells analysis of the 2dFGRS, but they use magnitude limited samples, and consider only deterministic bias models, whereas our present analyses use volume limited samples, and consider stochastic bias models.  The counts-in-cells method has also been used to calculate the variance and higher order moments of galaxy clustering in the 2dFGRS (Conway et al. 2004; Croton et al. 2004a,b; \\citealt{astro-ph/0401405}). \\nocite{astro-ph/0404276,astro-ph/0401434,astro-ph/0401406}  Many theoretical results on biasing from numerical models have also been reported. There are two main approaches to modelling galaxy distributions: semianalytic (e.g. \\citealt*{1997MNRAS.286..795K}; \\citealt{2000MNRAS.316..107B, 1999MNRAS.310.1087S}) and hydrodynamic \\citep[e.g.][]{1992ApJ...399L.113C,1999ApJ...522..590B,2000ApJ...538...83C,2001ApJ...558..520Y}. Comparisons are given by \\citet{2003MNRAS.338..913H} and \\citet{2002MNRAS.335..762Y}. Several studies have been made of galaxy biasing in these numerical simulations \\citep[e.g.][]{2001MNRAS.320..289S,2001ApJ...558..520Y}, but none provide results in sufficient detail to allow an easy comparison with the 2dFGRS. We therefore analyse a large new semianalytic calculation which is capable of yielding mock results that can be analysed in an identical manner to the real data.  In this paper we concentrate on a few aspects of relative bias mentioned above, splitting galaxies by spectral type and colour. We investigate the nonlinearity, stochasticity and scale dependence of the biasing relation through comparison with three models. Section 2 summarises the DL99 framework for biasing, and presents the bias models used in this paper. Section 3 describes the 2dFGRS catalogue, the derivations of the galaxies spectral types and colours and Section 4 explains the counts-in-cells method. In Section 5 we show the methods used for model fitting and error estimation. Section 6 gives the results and we compare our results with simulations in Section 7.  Throughout, we adopt a cosmological geometry with $\\Omega_m=0.3$, $\\Omega_v=0.7$ in order to convert redshifts and angles into three-dimensional comoving distances. We define our cells with $h=0.7$, and all cell lengths are quoted in $\\Mpc$, instead of the standard $\\mpcoh$. ",
        "conclusions": "We have presented fits of three relative biasing schemes to joint counts-in-cells distributions of 2dFGRS galaxies, separated by both colour and spectral type $\\eta$. Each scheme is convolved with a Poisson distribution to account for statistical `shot noise'. Our first two models present two alternative types of deterministic biasing: linear and power law bias. Linear bias is an important concept in cosmology and many results are linked to it, but it is not physically plausible as it allows negative densities. Power law bias presents a simple cure for this problem, but still has little physical motivation. With the advent of large semianalytic and hydrodynamic simulations, interest has grown in `stochastic' bias models. Bias could be determined by parameters other than the local overdensity of the dark matter, and considerable scatter could occur in the relation. Galaxy distributions have previously been measured to be well approximated as lognormal, therefore a bivariate lognormal distribution seems a natural model for relative bias between galaxy types. This model incorporates stochasticity and nonlinearity in a well defined manner, which is mathematically simple and consistent with observation.  To account for the discrete nature of galaxies, the {\\it Poisson sampling hypothesis} is assumed, and all models are convolved with a Poisson distribution. On small scales where our cell counts become shot noise dominated, we find this hypothesis to fail causing overestimates of variance compared to direct estimation methods. The main symptom of the discrepancy is a number of completely empty cells that exceeds the Poisson sampled lognormal prediction. This is found not to affect our results for stochasticity, and the same effect is seen in the simulations, but it emphasises the need for a greater understanding of Poisson statistics in relation to galaxy clustering.  We have detected a significant deviation from  $r_\\LN=1$ in the 2dFGRS and confirmed this detection of stochasticity through likelihood ratio tests, Kolmogorov-Smirnoff probability testing, and Monte Carlo simulations. We have measured stochasticity at a level of $\\sigma_b/\\hat{b}=0.44 \\pm 0.02$ or $r_\\LN=0.958 \\pm 0.004$ on the smallest scales ($10\\Mpc$), declining with increasing cell size. The nonlinearity of the biasing relation is less than $5\\%$ on all scales. The small measured values of stochasticity and nonlinearity support the use of galaxy redshift surveys for studies of the large scale distribution of matter in the universe, and the measurement of cosmological parameters. However, as precision in cosmology increases and new techniques are developed, the effects of stochastic bias on parameter estimation should be understood. For example, studies of cosmology through weak gravitational lensing requires knowledge of nonlinear and stochastic bias \\citep{astro-ph/0403698}. Our results for $r_{\\rm lin}$ on 10~Mpc scales are consistent within the (large) errors with galaxy-mass correlations measured by weak lensing surveys \\citep{2002ApJ...577..604H} on the largest scales probed.  A comparison with semianalytic simulations shows a similar variation of nonlinearity and stochasticity with scale. The amplitude of stochasticity appears a little lower than in the true data, particularly on large scales, and the nonlinearity is slightly greater on small scales. Nevertheless, given the known imperfections of the current generation of semianalytic calculations, the general agreement is certainly encouraging. We hope that this work will stimulate the investigation of more detailed biasing models. Through the linking of new simulations to observations, a more thorough understanding of the processes of galaxy formation and evolution should be within our reach."
    },
    "0404/astro-ph0404105_arXiv.txt": {
        "abstract": "The first eclipsing binary pulsar, PSR~B1957+20, was discovered in 1987. Since then, 13 other eclipsing low-mass binary pulsars have been found, 12 of these are in globular clusters. In this paper we list the known eclipsing binary pulsars and their properties, with special attention to the eclipsing systems in 47 Tuc. We find that there are two fundamentally different groups of eclipsing binary pulsars; separated by their companion masses. The less massive systems ($M_c \\, \\sim \\, 0.02 \\, \\rm M_{\\odot}$) are a product of predictable stellar evolution in binary pulsars. The systems with more massive companions ($M_c \\, \\sim \\, 0.2 \\, \\rm M_{\\odot}$) were formed by exchange encounters in globular clusters, and for that reason are exclusive to those environments. This class of systems can be used to learn about the neutron star recycling fraction in the globular clusters actively forming pulsars. We suggest that most of these binary systems are undetectable at radio wavelengths. ",
        "introduction": "The first eclipsing binary pulsar, PSR~B1957+12, was discovered in 1987 (\\cite{fst88}). Since its discovery, 13 other eclipsing binary pulsars with low-mass companions have been found, these are listed in Table \\ref{tab:binaries}, with relevant references. Of these systems, 12 are to be found in globular clusters (GCs). To these we add a list of binary systems with companion masses smaller than 0.02~M$_{\\odot}$ that are, as we will see, intimately related to a subset of the eclipsing binary systems.  The systems that have been known longest are also the ones for which the pulsar has the largest flux density, it is therefore understandable that such systems have been discussed with great detail in a large number of publications. A review of all these results is beyond the scope of this work, mainly because of limitations in available space. We will concentrate instead on the general properties of the population of eclipsing binaries and infer some general trends. In section \\ref{sec:47tuc}, we do concentrate specifically on the eclipsing binary pulsars in 47~Tuc; this is done for the sake of completeness, as most of the remaining eclipsing binary systems are discussed elsewhere in these proceedings.  In Table \\ref{tab:binaries}, we included mostly references that concentrate on the timing of the eclipsing pulsars. From these references, one concludes that for the 5 eclipsing binaries with a published timing baseline longer than 4 years (PSR~B1957+20 and PSR~J2051$-$0827, Terzan 5 A and 47~Tuc~J and O) random variations of the orbital period have invariably been detected.  \\begin{table} \\caption{\\label{tab:binaries} A set of eclipsing binary pulsars and related objects in GCs and in the Galactic disk. Those for which a timing solution is not known are followed by an asterisk, the names could change in the future when that solution is published. 47~Tuc~W has a position determined by optical astrometry. Those that are known to eclipse are followed by letter ``e''. The minimum companion masses ($m_c$) are calculated assuming a pulsar mass of 1.4 M$_{\\odot}$. References are: rh: Ransom, Hessels, these proceedings, dn: David Nice, these proceedings, st: Stappers, these proceedings, po: Possenti, these proceedings 1:~\\cite{clf+00}, 2:~\\cite{pdm+03}, 3:~\\cite{lbhb93}, 4:~\\cite{dpm+01,fpds01}, 5:~\\cite{lmbm00,nat00}, 6:~\\cite{rsb+04}, 7:~\\cite{fck+03}, 8:~\\cite{cha03}, 9:~\\cite{dlm+01,rgh+01}, 10:~\\cite{dma+93}, 11:~\\cite{fst88,arz95}, 12:~\\cite{sbl+96,sbm+98} } \\begin{tabular}{ l l c c c l } & & & & & \\\\ \\hline \\hline Pulsar & GC, letter & P(ms) & $P_b$ (days) & $m_c (M_{\\odot})$ & Ref. \\\\ \\hline J0024$-$7204V* & 47~Tuc~V (e?) & 4.81  & $\\sim$0.2 & $\\sim$0.3 & 1 \\\\ J0024$-$7204W* & 47~Tuc~W (e)  & 2.35  & 0.133     & 0.127 & 1 \\\\ J1701$-$3006B & M62~B (e)      & 3.59  & 0.145     & 0.124 & 2 \\\\ B1718$-$19    & NGC~6342A (e)  & 1004  & 0.258     & 0.117 & 3 \\\\ J1740$-$5340  & NGC~6397A (e)  & 3.65  & 1.354     & 0.188 & 4,po \\\\ J1748$-$2446A & Terzan 5A (e)  & 11.56 & 0.076     & 0.089 & 5,dn \\\\ J2140$-$2310A & M30~A (e)      & 11.02 & 0.174     & 0.101 & 6,rh \\\\ \\hline J0024$-$7204I & 47~Tuc~I       & 3.48  & 0.230     & 0.013 & 7 \\\\ J0023$-$7203J & 47~Tuc~J (e)   & 2.10  & 0.121     & 0.021 & 7 \\\\ J0024$-$7204O & 47~Tuc~O (e)   & 2.64  & 0.136     & 0.022 & 7 \\\\ J0024$-$7204P* & 47~Tuc~P      & 3.64  & 0.147     & 0.018 & 1 \\\\ J0024$-$7204R* & 47~Tuc~R (e)  & 3.48  & 0.066     & 0.026 & 1 \\\\ J1518+0204C*   & M5~C (e)      & 2.48  & 0.087     & 0.038 & rh \\\\ J1701$-$3006E* & M62~E (e)     & 3.23  & 0.16      & 0.03  & 8 \\\\ J1701$-$3006F* & M62~F         & 2.30  & 0.20      & 0.02  & 8 \\\\ J1807$-$2459A* & NGC~6544~A    & 3.06  & 0.071     & 0.009 & 9 \\\\ B1908+00*      & NGC~6760~A    & 3.62  & 0.141     & 0.018 & 10,dn \\\\ J1953+1846A*   & M71~A (e)     & 4.89  & 0.177     & 0.032 & rh \\\\ B1957+20       & (Galaxy) (e)  & 1.61  & 0.382     & 0.022 & 11,dn \\\\ J2051$-$0827   & (Galaxy) (e)  & 4.51  & 0.099     & 0.027 & 12,st \\\\ \\hline \\hline \\end{tabular} \\end{table} ",
        "conclusions": "We have listed the set of observational properties that lead us to believe that VLMBPs, LMBPs and ELMBPs are different sorts of objects. The first two classes are likely products of stellar evolution in binary systems, and can therefore be observed in GCs and the Galaxy. The second kind is likely the result of exchange encounters, such systems should only be observed in GCs. Many of them are unlikely to be undetectable, because of the effect of shrouding suggested by the existence of 47~Tuc~V. The rate of occurrence of such systems in GCs can be used to estimate neutron star recycling for the GCs that produce more pulsars."
    },
    "0404/astro-ph0404333_arXiv.txt": {
        "abstract": "Molecular hydrogen line radiation emitted in formation events of first-generation stars are evaluated in a discussion of its detectability by future observational facilities. H$_2$ luminosity evolution from the onset of prestellar collapse until the formation of a $\\sim 100 {\\rm M_{\\odot}}$ protostar is followed. Calculations are extended not only to the early phase of the runaway collapse but also to the later phase of accretion, whose observational features have not been studied before. Contrary to the runaway collapse phase, where the pure-rotational lines are always dominant, in the accretion phase rovibrational line emission becomes prominent. The maximum luminosity is also attained in the accretion phase for strong emission lines. The peak intensity of the strongest rovibrational line reaches $\\sim 10^{-29}{\\rm W/m^2}$, corresponding to the flux density of $10^{-5}{\\rm \\mu Jy}$, for a source at the typical redshift of first-generation star formation, $1+z=20$. Although the redshifted rovibrational H$_2$ emission from such an epoch falls in the wavelength range of the next-generation infrared satellite, {\\it Space Infrared Telescope for Cosmology and Astrophysics}, for exceeding the detection threshold $10^7$ such protostars are required to reach the maximum luminosity simultaneously in a pregalactic cloud. It is improbable that this condition is satisfied in a realistic scenario of early structure formation. ",
        "introduction": "First-generation stars can be a source of both the observed metals in the early universe and ultraviolet photons that cause the cosmological reionization. In fact, substantial metal enrichment has been observed in Ly$\\alpha$ clouds as early as $z\\sim 5$ (Cowie et al. 1995; Songaila \\& Cowie 1996; Songaila 2001). Furthermore, the WMAP satellite has measured the electron scattering optical depth of $\\tau_{e}=0.17 \\pm 0.04$, which corresponds to very early reionization at $z\\simeq 17\\pm 5$ (Bennett et al. 2003). The growing evidence, though still indirectly, suggest that first-generation stars did play a significant role in the early evolution of galaxies.  Undisputedly the most decisive evidence for the first star formation is the direct detection of ongoing formation events. First star-forming clouds are considered to be bright in H$_2$ line emission because the first stars are formed by gravitational collapse of primordial clouds, which is induced by H$_2$ line cooling (Saslow \\& Zipoy 1967; Peebles \\& Dicke 1968; Matsuda, Sato, \\& Takeda 1969).  In this paper, we evaluate the H$_2$ line emission associated with first-generation star formation. Similar work includes that by Kamaya \\& Silk (2002) and by Ripamonti et al. (2002). They, however, have followed only the evolution in the earlier phase of star formation, namely the runaway collapse phase, and have not extended their calculations well into the later phase, namely the main accretion phase. Here, we study the observational appearance in both evolutionary phases. We have found that the maximum value of H$_2$ line luminosity is indeed attained in the main accretion phase.  Currently, several next-generation large observational facilities, including {\\it James Webb Space Telescope} ({\\it JWST}) \\footnote{http://www.jwst.nasa.gov/}, {\\it Space Infrared Telescope for Cosmology and Astrophysics} ({\\it SPICA}) \\footnote{http://www.ir.isas.ac.jp/SPICA/}, and {\\it Atacama Large Millimeter Array} ({\\it ALMA}) \\footnote{http://www.alma.nrao.edu/}, are awaited. One of their major scientific goals is to explore the formation and evolution of pregalactic objects in the high-redshift universe. Among them, {\\it SPICA} is the most relevant to our purpose since it has the best sensitivity in the mid- to far-infrared ranges, where the strongest H$_2$ emission lines from the typical epoch for first star formation ($z \\sim 20$) are redshifted. In this paper, we therefore discuss the detectability of those H$_2$ line photons by {\\it SPICA}.  The outline of this paper is as follows: In section 2, the method of calculation is described. In section 3, the results for H$_2$ emission from a forming first star are presented. Finally, in section 4, a brief summary is given and the detectability of first-generation stars by {\\it SPICA} is discussed. ",
        "conclusions": "We have studied the evolution of H$_2$ line emission from a forming first star, starting from a prestellar dense core up to the formation of a $100{\\rm M_{\\odot}}$ protostar. After the formation of the central protostar, the H$_2$ line luminosity increases and luminosities of strong lines reach their maximum when the central protostellar mass is $\\sim 13{\\rm M_{\\odot}}$. The four strongest ${\\rm H_2}$ lines are the rovibrational lines of  1-0Q(1) ($2.34{\\rm \\mu m}$), 1-0O(3) ($2.69{\\rm \\mu m}$), the pure-rotational lines, 0-0S(4) ($8.27{\\rm \\mu m}$), and 0-0S(3) ($10.03{\\rm \\mu m}$). During the accretion phase, vibrational lines of 1-0Q(1) and 1-0O(3) are the strongest, while in the collapse phase the strongest lines are the pure-rotational lines.  Strong H$_2$ line emission is expected also from diffuse regions in low-metallicity ($Z < 0.01Z_{\\odot}$) forming galaxies (e.g., Shchekinov \\& \\'{E}nt\\'{e}l' 1985; Shchekinov 1991; Omukai \\& Kitayama 2003). However, the dominant lines from those sources are pure-rotational lines because of the lower density ($n \\sim 10^{0-3} {\\rm cm^{-3}}$). On the other hand, in the late stage of the first star formation, ${\\rm H_2}$ rovibrational lines are dominant, because of both high density ($n \\sim 10^{10-12} {\\rm cm^{-3}}$) and temperature ($T> 1000$K) in the H$_2$ emitting region. Therefore, detection of ${\\rm H_2}$ rovibrational lines can be a useful method to identify first-generation star forming clouds.  Several large observational projects using next-generation facilities are now in progress, including {\\it JWST}, {\\it SPICA} and {\\it ALMA}. The wavelength ranges of these three facilities are presented in Table II. \\begin{center} \\begin{tabular}{c|c|c} \\multicolumn{3}{c}{\\bf Table II. The wavelength ranges of future telescopes}\\\\ \\hline\\hline {\\it JWST}      & {\\it SPICA}    & {\\it ALMA} \\\\ \\hline $0.6-28{\\rm \\mu m}$  &$5-200{\\rm \\mu m}$  &$300{\\rm \\mu m}-3{\\rm mm}$\\\\ \\hline \\end{tabular} \\end{center} For typical formation epochs of the first stars $z\\sim 20$ (e.g., Abel et al. 2002), the emitted H$_2$ rovibrational line photons are redshifted to the far-infrared wavelength. As is seen in Table II, {\\it SPICA} has the deepest detection limit in this range and thus is the most suitable to our purpose. {\\it SPICA} is a next-generation infrared satellite, which is planned to have a cooled (4.5K) and large (3.5m) telescope and to be highly sensitive at the mid- and far-infrared wavelengths. The line detection limit of {\\it SPICA} is about $10^{-20}{\\rm W/m^2}$ in the far-infrared wavelength $\\lambda\\sim 40{\\rm \\mu m}$ with a spectroscopic resolution $R=3000$ and observation time of one hour. Since a signal to noise ratio is proportional to $\\sqrt{t}$, the sensitivity improves, of course, with longer integration time. Nonetheless, in the mid- and far-infrared wavelengths, the high background of the zodiacal light limits the sensitivity of {\\it SPICA} to $10^{-(21-22)}{\\rm W/m^2}$ (private communication H. Matsuhara).  The observed wavelengths of the four strongest lines redshifted from the emission at $1+z=20$, the typical epoch of first star formation, are given in Table III. \\begin{center} \\begin{tabular}{l|c|c|c|c} \\multicolumn{5}{c}{\\bf Table III. The Observed Wavelengths of Strong H$_2$ Lines}\\\\ \\hline\\hline line & 1-0Q(1) & 1-0O(3) & 0-0S(4)  & 0-0S(3) \\\\ \\hline rest frame &$2.34{\\rm \\mu m}$  &$2.69{\\rm \\mu m}$  &$8.27{\\rm \\mu m}$ &$10.03{\\rm \\mu m}$ \\\\ \\hline redshifted ($1+z=20$) &$46.8{\\rm \\mu m}$  &$53.8{\\rm \\mu m}$ &$165.4{\\rm \\mu m}$  &$200.6{\\rm \\mu m}$\\\\ \\hline \\end{tabular}. \\end{center} The wavelengths of the redshifted pure-rotational lines shown above are outside the range of any future telescope given in Table III. On the other hand, the wavelengths of the rovibrational lines, which are intrinsically stronger than pure-rotational lines, fall in the range where the sensitivity of {\\it SPICA} is excellent.  By using the value of peak luminosity $L_{\\rm peak}$ in Table I, the peak intensity of the strongest rovibrational lines from $1+z=20$ is \\begin{eqnarray} F_{\\rm peak}=\\frac{L_{\\rm peak}}{4\\pi D_{1+z=20}^2} =\\frac{1.73\\times10^{35}{\\rm erg}/{\\rm s}}{4\\pi(2.13\\times 10^{11}{\\rm pc})^2} \\simeq 0.32 \\times 10^{-28}{\\rm W}/{\\rm m}^2, \\end{eqnarray} where $D_{1+z=20}=2.13\\times 10^{11}{\\rm pc}$ is the luminosity distance to $1+z=20$ given by the formula \\begin{eqnarray} D_{1+z}=\\frac{c (1+z)}{H_0}\\int_{1}^{1+z} \\frac{dy}{[\\Omega_{\\Lambda}+\\Omega_{\\rm M}y^{3}]^{1/2}}, \\end{eqnarray} where $H_0$ is the Hubble parameter, with the values $H_0=70{\\rm km}{\\rm s}^{-1}{\\rm Mpc}^{-1}$, $\\Omega_{\\Lambda}=0.7$, $\\Omega_{\\rm M}=0.3$. This peak intensity corresponds to the peak flux density of $\\sim 10^{-5}{\\rm \\mu Jy}$ using the line width $\\nu_0 v/c \\sim 10^{10}{\\rm Hz}$, where $\\nu_0$ is the line-center frequency for the observer and $v$ is the velocity of the mass shell. Note that in considering a cluster of star-forming cores, a larger value of the line width should be taken because relative velocity between the cores is usually larger than the thermal velocity in a core. In this case, the peak flux density per unit core will be lower than the value given above.  Taking into account the limiting sensitivity $10^{-(21-22)}{\\rm W/m^2}$ and the maximum rovibrational line flux $\\sim 10^{-29} {\\rm W/m^2}$, we conclude that {\\it SPICA} is able to observe a cluster of forming first-generation stars if it is contains more than $10^{7}$ sources with their luminosity near the peak value.  Next, we consider whether or not this situation is possible in a realistic cosmological context. The epoch of the peak luminosity lasts only for $\\Delta t_{\\rm peak} \\sim 10^{3} {\\rm yrs}$ (see Fig 2), which is by far shorter than the typical free-fall time in the pregalactic clouds (say, $t_{\\rm ff} \\gsim 10^{7}$yrs). Even in the extreme case of starburst, star formation is a prolonged event that continues at least a free-fall timescale of the pregalactic cloud. Thus, at a given time, at most only $\\Delta t_{\\rm peak} /t_{\\rm ff} \\lsim 10^{-4}$ of the total gas mass of the pregalactic cloud is a total mass of star forming cores approximately at peak luminosity. To reach the luminosity peak, the first-star forming cores must be as massive as $13{\\rm M_{\\odot}}$ (\\S 3). The mass scale of actual cores is probably more massive, typically $10^{2-3}{\\rm M_{\\odot}}$, according to numerical simulations of the fragmentation of primordial clouds (e.g., Bromm et al. 2002). Thus the total mass of 10$^{7}$ cores is $\\sim 10^{9}{\\rm M_{\\odot}}$. Since only $10^{-4}$ of all the cores are at peak luminosity, in order for 10$^{7}$ cores to reach peak luminosity at the same time, the total cloud mass must be as high as $10^{13}{\\rm M_{\\odot}}$. Note also that in this case the star formation rate is required to be an improbably high value of $10^6{\\rm M_{\\odot}/yr}$. It is true that for pure-rotational lines, i.e., 0-0S(3), the epoch in which the luminosity is higher than $10^{34}$erg/s continues for as long as $10^5$yrs. Although this long emission period reduces the required star formation rate, it is still improbably high ($10^5{\\rm M_{\\odot}/yr}$). On the other hand, in order to emit radiation mainly in H$_2$ lines during the prestellar collapse, especially for the very high density region, the metallicity in the pregalactic clouds must be lower than $\\sim 10^{-4}Z_{\\odot}$ (Omukai 2000). However, formation of Galactic scale objects with such low metallicity is probably excluded in the context of standard cosmology (Scannapieco, Schneider, \\& Ferrara 2003). In conclusion, it is improbable that detection of H$_2$ line emission from forming first stars can be accomplished with {\\it SPICA}.  In this paper, the accretion flow has been assumed to be spherical. In reality, however, an accretion disk should form owing to an angular momentum barrier. In this case, the total energy emitted in ${\\rm H}_2$ lines might increase, since the released gravitational energy is converted more efficiently into radiation in the accretion disk. At the same time, H$_2$ is more susceptible to photodissociation owing to direct irradiation from the star. To estimate the H$_2$ luminosity more precisely, it is necessary to study effects of the disk in the accretion phase. We leave this problem for future work. Recently, Bromm \\& Loeb(2004) performed a three-dimensional numerical simulation and derived the mass accretion rate onto the protostar. Their accretion rate is similar to our value until $\\sim 10^3$ yr after the protostar formation. After this time, their accretion rate drops more rapidly than ours. This difference is probably caused by more realistic evolution that they calculated, e.g. quasi-static contraction in the early evolution and/or the complex accretion flow owing to 3D effects etc. In our calculation the maximum values of the H2 luminosities are reached at $\\sim 300$ yr after the protostar formation. At this epoch, our accretion rate does not differ much from their value. Therefore, the peak value of H2 luminosity is expected to be almost the same even if their accretion rate is adopted. After $\\sim 10^3$ yr of protostar formation, we are probably overestimating the accretion rate and thus the H2 luminosity. Nevertheless, the H2 luminosity at this time is already weak and does not important for our discussion on observability by future facilities.  Contrary to our claim, Kamaya \\& Silk (2002) and Ripamonti et al.(2002) have concluded that future observational facilities, e.g., {\\it JWST}, {\\it ALMA}, {\\it ASTRO-F}, are able to detect H$_2$ emission from the first star-formation sites. Here we stress that our estimated value of the peak luminosities of ${\\rm H_2}$ lines is larger than their value, which is attained for the earlier phase. Moreover, in the mid- to far-infrared wavelengths, the sensitivity of {\\it SPICA} that we use here to estimate the detectability exceeds that of the other facilities that these authors considered. The reason for the apparent contradiction is as follows: in converting the emitted H$_2$ luminosity to the observed flux, they used the luminosity distance underestimated by a factor of $1+z$, where $z$ is the redshift of the source. This error propagates to the final value of the observed flux, which is overestimated by $(1+z)^2$. Note that, anyway, our value of the intrinsic luminosity in the runaway collapse phase is consistent with Kamaya \\& Silk (2002) and Ripamonti et al.(2002).  Finally, we touch on another future telescope, {\\it the Single Aperture Far-Infrared Observatory} ({\\it SAFIR}) \\footnote{ http://safir.jpl.nasa.gov/whatIs/index.asp}. Thanks to its larger mirror size (8 m), the line detection limit of {\\it SAFIR} is better than that of {\\it SPICA} and the background limit of $\\sim 10^{-22}{\\rm W/m^2}$ at $\\lambda \\sim 40{\\rm \\mu m}$ can be reached in one hour with spectroscopy of $R=1000$. Recall that the detection limit we used in this paper is also the background limit $\\sim 10^{-22}{\\rm W/m^2}$, which is set by the zodiacal light. Therefore, even with {\\it SAFIR} we conclude that detecting ${\\rm H}_2$ line emission from forming first stars is difficult.  The authors would like to thank Tetsu Kitayama, Hajime Susa and Masayuki Umemura for valuable comments, and Ken-ichi Oohara and Kazuya Watanabe for continuous encouragement. This work was supported in part by the Ministry of Education, Culture, Sports, Science, and Technology of Japan, Grants-in-Aid for Scientific Research No. 14540227 (RN) and Research Fellowship of the Japan Society for the Promotion of Science for Young Scientists, grant 6819 (KO).  \\clearpage \\newpage  \\begin{figure} \\begin{center} \\FigureFile(80mm,80mm){figure1a.eps} \\FigureFile(80mm,80mm){figure1b.eps} \\end{center} \\caption{(a) The temperature evolution and (b) the evolution of the H$_2$ (bold line) and H (thin line) concentration against the number density. For the runaway collapse phase, the evolutionary trajectory of the central part is drawn with a solid line. For the accretion phase, the evolutionary trajectories of the mass shells of $M/ {\\rm M}_{\\odot} = $ 1(long-dashed lines), 13(short-dashed lines) and 100(dotted lines) are shown.} \\label{fig:TH2} \\end{figure}  \\newpage  \\begin{figure} \\begin{center} \\FigureFile(120mm,80mm){figure2.eps} \\end{center} \\caption{The luminosity evolution of the four strongest ${\\rm H_2}$ lines, the rovibrational lines of 1-0Q(1) (solid lines), and 1-0O(3) (long-dashed lines), and the pure rotational lines of 0-0S(4) (short-dashed lines), and 0-0S(3) (dotted lines). The horizontal axis is the time since the onset of the collapse. The evolution up to the formation of a $10^2 {\\rm M_{\\odot}}$ protostar is shown. The formation of the central protostar corresponds to $t=0$. Note that all lines attain their maximum luminosity at $t>0$, i.e., in the main accretion phase. Also the luminosities of rovibrational lines become stronger than that of pure-rotational ones in this phase. } \\label{fig:L(time)2} \\end{figure}  \\newpage  \\begin{figure} \\begin{center} \\FigureFile(80mm,60mm){figure3.eps} \\end{center} \\caption{The evolution of the liberation rate of chemical binding energy (dotted line) and H$_2$ line emissivity (solid line) for the mass shell (13${\\rm M}_{\\odot}$). The time $t=0$ is set when the mass shell (13${\\rm M}_{\\odot}$) is supposed to reach the center ($r=0$). Because of the finite radius of the star, the mass shell reaches the stellar surface before $t=0$.} \\label{fig:energy} \\end{figure}  \\newpage  \\begin{figure} \\begin{center} \\FigureFile(120mm,80mm){figure4.eps} \\end{center} \\caption{Liberation rate of chemical binding energy, at the final stage of the collapse phase (solid line), when the mass shells of $M_{\\rm r}/ {\\rm M}_{\\odot} = $ 5 (long-dashed line), 13 (short-dashed line) and 50 (dotted line) fall on the center, respectively. The horizontal axis is the distance from the center. The area surrounded by each line and the horizontal axis indicates the chemical energy liberated per unit time.} \\label{fig:energy} \\end{figure}  \\newpage  \\begin{figure} \\begin{center} \\FigureFile(80mm,40mm){figure5a.eps} \\FigureFile(80mm,40mm){figure5b.eps} \\FigureFile(80mm,40mm){figure5c.eps} \\end{center} \\caption{The evoluton of (a) number density, (b) temperature, (c) ${\\rm H_2}$ concentration of selected mass shells against the radial distance from the central protostar in the main accretion phase. The values for the mass shell ($13{\\rm M}_{\\odot}$) are shown by the thick solid line. Recall that the radii of the mass shells decrease with time since the envelope matter is falling onto the protostar. For comparison, we also plot the values of other two mass shells that contain $1{\\rm M}_{\\odot}$ (dashed line) and $100{\\rm M}_{\\odot}$ (dotted line) inside. } \\label{fig:nTH2} \\end{figure}   \\clearpage \\newpage"
    },
    "0404/astro-ph0404455_arXiv.txt": {
        "abstract": "In order to investigate discrepancies between recent published estimates of the the HI mass function (HIMF), we explore the impact of distance uncertainties on the derivation of the faint end slope of mass and luminosity functions of galaxies in the local volume by deriving HIMFs from mock HI surveys. We consider various survey geometries and depths and compare the HIMFs measured when using ``real'' distances, distances derived by assuming pure Hubble flow and distances assigned from parametric models of the local velocity field.  The effect is variable and dependent on the exact survey geometry, but can easily lead to incorrect estimates of the HIMF, particularly at the low mass end. We show that at least part of the discrepancies among recent derivations of the HIMF can be accounted for by the use of different methods to assign distances. We conclude that a better understanding of the local velocity field will be necessary for accurate determinations of the local galaxy luminosity and mass functions. ",
        "introduction": "One of the perplexing discrepancies between observations and CDM theory revolves around the ``missing satellite'' problem. Numerical simulations in $\\lambda$CDM predict a value for the logarithmic slope of the faint end of the halo mass function, $\\alpha$, which is close to the value of $-1.8$ that arises analytically from the Press-Schechter formalism (Press \\& Schechter 1974), but most recent determinations of the optical luminosity function (LF) in the local volume yield values of $\\alpha$ that are significantly flatter. Complementary to the results on the optical LF, several determinations of the local HI mass function (HIMF) have likewise yielded relatively flat faint-end slopes. The over-prediction of the number of low mass objects relative to those actually observed is considered one of the last remaining stumbling blocks for $\\lambda$CDM, and points towards a new understanding of the baryon physics of galaxy formation. It is obviously desirable that all of the uncertainties in the observed luminosity and mass functions be well understood.  To derive a luminosity (or mass) function from observational data, the sample selection effects must be considered carefully. In most surveys, the lowest luminosity (or mass) objects are only visible nearby, and their contribution must be weighted accordingly. Fractional distance errors can be large for relatively nearby objects, and can thus have a very strong impact on determinations of the faint end characteristics of luminosity (or mass) functions. When those functions are evaluated using data obtained in clusters, the principal source of uncertainty lies in the assessment of membership; assuming that uncertainty is to first order independent on luminosity (or mass), it can have a small impact on the estimate of LF slopes. Determinations of the HIMF, however, have to rely on relatively shallow, wide area surveys to detect low HI mass objects over a sufficient volume, and are thus most susceptible to uncertainties associated with distance errors in the nearby population.  Primary distances (eg. cepheids, the tip of the red giant branch or surface brightness fluctuations, the best of which give distances to $\\pm$ 10\\%) are available for $\\sim 600$ galaxies, most within $\\sim 30$ Mpc. Secondary distances (which rely on the above methods for calibration) carry larger errors ($\\sim$ 15--30\\%) but do account for peculiar velocities for some 5000 objects. For all other galaxies, the only indication of distance is their redshift. Outside of $cz=6000$ \\kms, pure Hubble flow works well, since a typical peculiar velocity of a few hundred \\kms ~creates an uncertainty on the redshift distance of $<$ 10\\%. For closer galaxies, however, the relative error is much larger, and, for the nearest objects, whose recessional velocities are dominated by peculiar motions, the error can be a factor of two or more. Furthermore, peculiar velocities are not random but exhibit coherent flows towards local over-densities. In some directions, the redshift-distances of all (nearby) galaxies will be systematically under- or over- estimated if peculiar velocities are ignored.  The effects of peculiar velocities on distance determinations can be mitigated by the use of parametric models of the local velocity field. While clearly oversimplifications of the real flows, such models give a direction-dependent redshift-distance prescription. A recent example of such a model is the multi-attractor model of Tonry \\etal ~(2000; hereafter TB00: also see Willick \\& Batra ~2001).  In order to investigate the effects of distance uncertainties on local luminosity or mass functions, we consider here the derivation of the HIMF. Recent determinations of the local HIMF have yielded discrepant results on the abundance of low HI mass objects, namely: Zwaan \\etal ~(1997; Z97), Rosenberg \\& Schneider (2002; RS02) and Zwaan \\etal ~(2003; Z03). The Z03 HIMF is based on the HI Parkes All Sky Survey (HIPASS) Bright Galaxy Catalogue (BGC), while the Z97 and RS02 HIMFs are both based on drift scan surveys conducted at Arecibo during the period of its recent upgrade (respectively called AHISS, the Arecibo HI Strip Survey, and ADBS, the Arecibo Dual Beam Survey). A summary of the results of those surveys is provided in Table 1, parameterized as Schechter functions, \\be \\phi(x) = \\frac{dn}{d\\log M} = \\ln 10 ~\\phi_\\star~ x^{(\\alpha+1)}~ e^{-x}, \\ee where $x=M/M_\\star$, $M$ is the HI mass, and $n$ is the total number density of galaxies. The faint end slopes of the HIMF derived by the three surveys vary between $-1.20$ and $-1.53$, yielding extrapolations below $M_{HI} = 10^7$ \\msun~ that disagree by an order of magnitude, the RS02 HIMF having the steeper slope.  \\begin{deluxetable}{lccc} \\tablecaption{Parameters in published HIMFs} \\tablewidth{0pt} \\tablehead{ \\colhead{Reference} &\\colhead{Z97} & \\colhead{RS02} & \\colhead{Z03} \\\\ \\colhead{Survey} &\\colhead{AHISS} & \\colhead{ADBS} & \\colhead{HIPASS} } \\startdata $\\phi_\\star [h_{70}^3 {\\rm Mpc}^{-3}]$ & 0.0048 & 0.0047 & 0.0070  \\\\ $\\log(M_\\star/M_\\odot) + 2 \\log h_{70}$ & 9.86 & 9.94   & 9.85  \\\\ $\\alpha$ & -1.20 & -1.53  & -1.30  \\\\ Survey area & 65 deg$^2$ & 430 deg$^2$  & 20,600 deg$^2$ \\\\ $D_{\\rm lim}$ for $10^8$\\msun & 100 Mpc & 13 Mpc & 7 Mpc \\\\ \\# galaxies & 61 & 265 & 1000 \\\\ \\enddata \\end{deluxetable}  Previous studies of the impact of distance errors (from ignoring peculiar velocities) on HIMFs have modeled the effect, often called the `Eddington effect', by adding random Gaussian noise to the distances. Using this model, RS02 conclude that the effect is negligible in their survey volume until velocity dispersions of 600 \\kms ~are reached, while Z03 find that for an assumed velocity dispersion of 50-100 \\kms ~in the local universe ($cz < 3000$ \\kms) the slope of the low mass end of the HIMF probably steepens by $\\Delta \\alpha < 0.05$. Since the peculiar velocity field in the local universe is not random but rather dominated  by in-fall onto the Virgo cluster, these models could easily hide the main bias. In this paper, we investigate whether all or part of the discrepancy among the three HIMF determinations can be accounted for by the effects of distance uncertainties introduced by the different methods used by Z97, RS02 and Z03 to assign distances.  Here, we construct mock catalogs designed to mimic the HI surveys conducted by Z97, RS02 and Z03 and then derive mock HIMFs under different distance assumptions. $H_\\circ = 70 ~{\\rm km~s}^{-1}{\\rm Mpc}^{-1}$ is assumed throughout, and the dependence on $H_\\circ$ is explicitly listed using $h_{70} = H_\\circ/(70~ {\\rm km~s}^{-1}{\\rm Mpc}^{-1}$). ",
        "conclusions": "We have used mock HI survey data to demonstrate that distance uncertainties can create serious biases in the low mass end of the HIMF; similar arguments could be applied to studies of the field galaxy optical LF. When the effects of the local peculiar velocity field are ignored, incorrect estimates of the HIMF will be derived. The exact nature and size of the bias will depend both on the survey geometry and its depth.  In a survey constructed to simulate the HIPASS BGC, the low mass HIMF slope is {\\bf underestimated} if peculiar velocities are neglected. We suggest that the HIMF reported by Z03 suffers from this problem, which when corrected, would bring the Z03 and RS02 HIMFs into agreement within their errors. The adjusted value of the Z03 low mass slope is $\\alpha=-1.4\\pm0.1$, which is to be compared with the RS02 value of $\\alpha=-1.5\\pm0.2$. This change does not affect the implied HI mass density since it is dominated by $M_\\star$ galaxies.  For a survey where low mass galaxies are seen only in the foreground of Virgo, the distances (and therefore masses) of the nearby low mass galaxies will be overestimated and the low mass end of the HIMF {\\bf seriously underestimated} if in-fall onto Virgo is neglected. In deeper surveys, the effect becomes negligible, so that differences in the adopted distances alone cannot explain the discrepancy between the Z97 and RS02 HIMFs.  A simulated survey in the anti-Virgo region points to the limitations of currently available models for the local velocity field. In such a survey, use of distances taken from a multi-attractor model for the local velocity field leads to an underestimate of the low mass HIMF slope. Currently available flow models are not well constrained in the anti-Virgo direction because of a lack of galaxies with primary distances in the region. Velocity field models based on larger numbers of tracer galaxies will be required to solve this problem.  The use of a good model for the local peculiar velocity field is found to be essential to provide distances for nearby galaxies for which no redshift-independent distances are available. Without such a model, it is difficult to derive an unbiased HIMF, and in particular, the low mass end may be in error. Disentangling the effects of environment and evolution of the HIMF will not be possible unless the local flow model is well understood. Other applications where distances to galaxies are needed will also benefit from a better understanding of local flows."
    },
    "0404/astro-ph0404349_arXiv.txt": {
        "abstract": "We study the potential of weak lensing surveys to detect clusters of galaxies, using a fast Particle Mesh cosmological N-body simulation algorithm specifically tailored to investigate the statistics of these \\emph{shear selected} clusters. In particular, we explore the degree to which the radial positions of galaxy clusters can be determined \\emph{tomographically}, by using photometric redshifts of \\emph{background} source galaxies.  We quantify the errors in the tomographic redshifts, $\\Delta z \\equiv z_{\\rm tomography}-z_{\\rm real}$, and study their dependence on mass, redshift, detection significance, and filtering scheme. For clusters detected with signal to noise ratio $S/N\\gtrsim 4.5$, the fraction of clusters with tomographic redshift errors $|\\Delta z| \\leq 0.2$ is $45\\%$ and the root mean square deviation from the real redshifts is $\\langle(\\Delta z)^2\\rangle^{1/2}=0.41$, with smaller errors for higher $S/N$ ratios.  A Tomographic Matched Filtering (TMF) scheme, which combines tomography and matched filtering, is introduced which optimally detects clusters of galaxies in weak lensing surveys.  The TMF exploits the extra information provided by photometric redshifts of background source galaxies, neglected in previous studies, to optimally weight the sources. The efficacy and reliability of the TMF is investigated using a large ensemble of mock observations from our simulations and detailed comparisons are made to other filters. Using photometric redshift information with the TMF enhances the number of clusters detected with $S/N\\gtrsim 4.5$ by as much as $76 \\%$, and it increases the dynamic range of weak lensing searches for clusters, detecting more high redshift clusters and extending the mass sensitivity down to the scale of large groups. Furthermore, we find that coarse redshift binning of source galaxies with as few as three bins is sufficient to realize the gains of the TMF. Thus, the tomographic filtering techniques presented here can be applied to current ground based weak lensing data in as few as three bands, since two colors and magnitude prior are sufficient to bin source galaxies this coarsely. Cosmological applications of shear selected cluster samples are also discussed. ",
        "introduction": "\\label{sec:intro}  Mapping the distribution of dark matter in the universe is now one of the primary goals of observational cosmology.  The traditional approach has been to survey the sky for a population of luminous objects, assuming a direct relationship between luminosity and dark matter. While large galaxy surveys such as the 2dF Galaxy Redshift Survey \\citep{Coll01} and the Sloan Digital Sky Survey \\citep{York00} are perhaps the first examples that come to mind, future wide field X-ray \\citep{Romer01} and Sunyaev-Z'eldovich (SZ) surveys \\citep{Carl02,Koso03} will usher in an era where maps of the universe will be available at multiple frequencies.  Generically, all objects that emit radiation will be biased tracers of the underlying dark matter distribution. Hence, the interpretation of this data is limited by our ability to quantify the relationship between dark and luminous matter.  Weak gravitational lensing, or the coherent distortion of images of faint background galaxies by the foreground matter distribution, provides a unique opportunity to map the dark matter directly without making any assumptions about how baryons trace dark matter (see e.g., \\citet{BaSc01} or \\citet{vWM03} for a review).  Recently, weak lensing by large scale structure, or ``cosmic shear'', has been independently detected by several groups (see Table 1 of \\citet{vWM03} for a compilation of references).  The primary focus of these observations has been to measure the two point statistical properties of the shear field, which probes the matter power spectrum in projection.  However, as first suggested by \\citet{Kaiser95} and \\citet{Sch96}, weak lensing surveys can also be used to detect individual mass concentrations, allowing one to construct a \\emph{shear selected} sample of dark halos. Typically these will be galaxy cluster size $M \\gtrsim 10^{13.5} \\ \\msol$ objects, as the finite number of background sources and their intrinsic ellipticity place a lower limit on the size of a halo that can be detected.  To date, there are only two such cases of spectroscopically confirmed cluster detections from `blank field' weak lensing observations \\citep{Witt01,Witt03}. Although, recently \\citet{Schirm03,Schirm04} have used weak lensing to confirm several color-selected optical cluster candidates.  A shear selected sample of galaxy clusters would be of great astrophysical and cosmological interest. Probable biases in optical, X-ray, and SZ selected cluster samples with respect to richness, baryon fraction, morphology, and dynamical state, would be revealed \\citep{Hughes04}.  Furthermore, if there does exist a population of high mass to light clusters that has hitherto gone undetected, this would cast serious doubt upon the ``fair sample'' hypotheses \\citep{White93}, invoked to determine the matter density parameter, $\\Omega_{\\rm m}$, from comparisons of cluster baryon fraction to total cluster mass measurements (see e.g., \\citet{Allen03} or \\citet{VV04} for recent examples).  Perhaps a more contentious issue is the possible existence of a population of truly dark clusters. Recently, several groups have purportedly detected dark mass concentrations from weak lensing \\citep[though see Erben \\etal 2003]{Fis99,Erben00,UF00,Mir02,Dahle03}.  If these objects are actually `dark' virialized mass concentrations as opposed to non-virialized objects \\citep{Wein02} or projections of large scale structure \\citep{MWL01,WvWM02,Paddy03,HTY03}, this would pose serious challenges for our current structure formation paradigm and theories of galaxy formation. The absence of galaxies would require a complex form of feedback that suppresses galaxy formation in clusters.  If such a dark cluster were X-ray or SZ faint, one would have to invoke some exotic mechanism to segregate dark matter from baryons on Mpc scales.  It is well known that a large sample of clusters of galaxies out to moderate redshift $z \\sim 0.5$ can be used to impose stringent constraints on cosmological parameters.  Several authors have recently suggested using the number count distribution of galaxy clusters, from X-ray and SZ observations \\citep{HMH01,WB03}, deep optical surveys \\citep{New02}, and weak lensing \\citep{BPB02,Wein03}, to probe the nature of the dark energy believed to be causing the acceleration of the expansion of the universe.  Besides probing the dark energy, these cluster samples would constrain the matter density and amplitude of mass fluctuations \\citep[see e.g.,][]{BF98,HHM01,Pier03}. In addition, the clustering of such cluster samples would provide another means to measure the matter power spectrum \\citep{TED98,Schue01} which would provide complimentary parameter constraints \\citep{MM03b} and might possibly allow for a measurement of the angular diameter distance-redshift relation \\citep{Cooray01} or baryon wiggles \\citep{HH03}.  Yet, a limitation of the `baryon selection' in optical, X-ray, and SZ cluster samples, is that their selection function depends on observables (e.g., richness, velocity dispersion, flux, temperature, SZ decrement) that serve as a proxy for mass. Attempts to model these mass-observable relationships depend on the uncertain astrophysics of galaxy formation and the state of baryons in clusters. While these can be determined empirically, scatter in the correlation between any two such observables is significant \\citep[see e.g.,][]{Koch03}; furthermore, these mass-observable relations surely evolve with redshift.  The mass function of galaxy clusters is steep and exponentially sensitive to changes in limiting mass; moreover, errors in the mass-observable relations mimic cosmological parameter changes \\citep{VL99,LSW02,MM03a,Hu03}.  Inclusion of uncertainties in the mass-observable relations significantly degrades constraints on cosmological parameters obtained from surveys that rely on a baryonic proxy for mass \\citep{LSW02,MM03a}. Although consistency checks provide additional leverage \\citep{Hu03}, expensive follow up mass measurements are in general required to empirically calibrate the relationships between light and mass (Majumdar \\& Mohr 2003a; though see Lima \\& Hu 2004).  In this regard, the great advantage of a shear selected sample of galaxy clusters is that the selection function can be predicted \\emph{ab initio}, given a model for structure formation.  Because on the scales of interest for weak lensing, only gravity is involved and the mass is dominated by dark matter, precision cosmological measurements will in principle be limited by instrumental systematics rather than unknown astrophysics.  An accurate determination of the statistics of shear selected clusters from weak lensing requires N-body simulations of cosmological structure formation \\citep[though see][for simple analytical treatments]{KS99,BPB02,Wein03}, which brings us to the subject of this work.  In this paper we consider the detection of shear selected clusters using fast Particle Mesh (PM) N-body simulations of cosmological structure formation specifically tailored to investigate the statistics of weak lensing by clusters.  Similar numerical studies carried out by other groups \\citep{Reb99,WvWM02,Paddy03,HTY03}, have focused on the completeness and efficiency of weak lensing cluster searches.  Devising an optimal filtering scheme to detect clusters has received little attention and furthermore, the additional information provided by photometric redshifts of background source galaxies has been neglected. In this paper we introduce the fast PM simulation algorithm, deduce the optimal cluster detection strategy, and investigate cluster tomography.   A fundamental obstacle for weak lensing studies of the matter distribution is the lack of radial information. All of the matter along the line-of-sight to a distant source contributes to the lensing and so the distortion reflects a two dimensional projection of the dark matter. This limitation can be overcome by using photometric redshifts of the background source population, which will be distributed in redshift, thus allowing for a \\emph{tomographic} reconstruction of the three dimensional foreground matter distribution.  A variant of this technique has already been applied to real weak lensing data. \\citet{Witt01,Witt03} tomographically determined the redshifts of two \\emph{shear selected} clusters at $z_{\\rm d}=0.30$ and $z_{\\rm d}=0.70$ respectively, which were both confirmed by follow up spectroscopy to be at $z=0.28$ and $z=0.68$. Recently, \\citet{Tay04} conducted a weak lensing analysis of the previously known Abell supercluster A901/2 at $z=0.16$, and tomographically confirmed its redshift from weak lensing alone. They even constructed the first 3-D map of the dark matter potential of this cluster using the inversion method introduced by \\citet{Tay01}.  \\citet{HK02} and \\citet{BT03} applied the \\citet{Tay01} tomographic inversion method to analytical models of clusters put in `by hand', with mixed success.  However, the reliability of these tomographic reconstruction techniques have yet to be characterized on realistic distributions of matter from N-body simulations; although it is clear that the fidelity of the reconstructions will be severely degraded by line of sight projections of large scale structure.  In this work we introduce a complementary tomographic technique which is effectively two dimensional. It is qualitatively similar to that used by \\citet{Witt01,Witt03} and \\citet{Tay04} and also to maximum likelihood methods developed to study cluster mass profiles \\citep{GS98,SKE00,KS01}. The efficacy and reliability of the method is investigated by applying it to our large ensemble of N-body simulations.  Additionally, a Tomographic Matched Filtering (TMF) scheme, which combines tomography and matched filtering, is introduced which takes advantage of the additional radial information provided by photometric redshifts of source galaxies. The TMF is similar in spirit to the matched filtering algorithms used to find clusters in optical surveys \\citep{Postman96,Kepner99,WK02,Koch03}.  Applying it to our ensemble of simulations, we find that it is superior to filtering techniques used previously \\citep{Reb99,WvWM02,Paddy03,HTY03}, detecting more clusters per square degree and probing lower masses and higher redshifts.  The outline of this paper is as follows. In \\S 2 we describe our implementation of a PM code to simulate weak lensing observations. The formalism behind our maximum likelihood tomographic technique and the adaptive matched filtering method is introduced in \\S 3. In \\S 4 we apply the TMF and other filters to a large ensemble of simulations and determine the `optimal' filter for cluster detection. The completeness of this optimal filter is discussed in \\S 5. We assess the reliability of tomographic redshifts in \\S 6 and conclude in \\S 7. ",
        "conclusions": "We have described a fast efficient simulation scheme for computing the statistics of shear selected galaxy clusters from weak lensing. By tiling the line of sight with 32 unique PM simulations, our algorithm is unique in that it densely samples the primordial Gaussian distribution of random phases of large scale structure along the light cone, providing an accurate representation of the statistics of rare events. We have confirmed that the two point statistical properties of the three dimensional density field (Figure \\ref{fig:tiling}) and the abundance of galaxy clusters in the light cone (Figure \\ref{fig:dndz}) are accurately reproduced.  Although the dynamic range of our PM simulations is limited, we argued that high resolution simulations are not necessary to study shear selected cluster because of the small scale shot noise limit of weak lensing observations.   A maximum likelihood tomographic technique was presented which utilizes the photometric redshifts of source galaxies to determine the radial position of a galaxy cluster from weak lensing alone. We introduced a tomographic matched filtering (TMF) shceme, which optimally incorporates this additional radial information to detect more clusters.  We applied the TMF to a large ensemble of simulations and found that it enhances the number of clusters detected with $S/N\\gtrsim 4.5$ by as much as $76 \\%$ (see Table \\ref{table:eff_tomo} and \\ref{fig:eff_dist}).  As illustrated by Figure \\ref{fig:eff_dist}, the TMF increases the dynamic range of weak lensing searches for clusters, detecting more high redshift clusters and extending the mass sensitivity down to the scale of large groups.  Furthermore, binning the sources coarsely using only three redshift bins was sufficient to reap the benefits of tomography.  Thus the filtering scheme developed here can be applied to ground based weak lensing observations in as few as three bands, since two colors and a magnitude provide enough information to bin sources this crudely.  For the case where the photometric redshifts of sources are known more precisely, as could be achieved with multicolor ground based data, we quantified the errors in the tomographic redshifts. For the densities of sources used in this paper, $45\\%$ of clusters detected with signal to noise ratio $S/N\\gtrsim 4.5$ had tomographic redshift errors $|\\Delta z| \\leq 0.2$ and the rms deviation from the real redshifts is $\\sigma_{\\Delta z}=0.41$, with smaller errors for higher $S/N$ ratios.      The most appealing property of shear selected cluster samples in comparison to other methods of detecting clusters, is that \\emph{the expected cluster distributions can be reliably determined for any cosmological model using simulations like those described in this paper.} Because on the scales of interest only gravity is involved, comparing theory to observations does not, in principle, depend on assumptions about the relationship between dark and luminous matter. However, in practice projection effects and incompleteness are likely will complicat the interpretation of shear selected samples.  Weak lensing searches for clusters are plagued by projection effects. In particular, we found that the maximum intrinsic efficiency of weak lensing cluster surveys is $\\sim 85\\%$.  Phrased in terms of our criteria for matching peaks with clusters (\\S 2.5), of order $\\sim 15\\%$ of the most significant peaks detected in noiseless weak lensing maps do not have a collapsed halo with $M >10^{13.5} \\ \\hmsol $ within a $3^{\\prime}$ aperture.  This intrinsic inefficiency, also noted by previous investigators \\citep{WvWM02,Paddy03,HTY03}, arises because of the broad weak lensing kernel and confusion from large scale structure fluctuations.  The calculated maximum intrinsic efficiency is sensitive to parameters. Lowering the limiting mass, $10^{13.5} \\ \\hmsol$, of our cluster catalog, increasing the size of the aperture, $3^{\\prime}$, used to correlate peaks with halos, or reducing the linking length, $b=0.2$, of the FOF group finder, would all increase the efficiencies deduced.  Nevertheless, some of the most significant peaks in the mass maps do not correspond to actual clusters.  While cluster finding in numerical simulations involves applying a group finding algorithm to 3-d dark matter particle distributions, observationally determining whether there is a galaxy cluster at the location of a given peak will depend on observational criteria like the overdensity of galaxies, X-ray flux, SZ decrement, or the presence of a brightest cluster galaxy. \\emph{Thus, any criteria employed to distinguish the real clusters from the projections will rely on detecting baryons}. Because our ability to detect these baryons is a function of mass and redshift, and because of the scatter between baryonic observables and the underlying mass of the dark matter halo, the projection effects in weak lensing searches will degrade the well defined selection function deduced from `dark matter only' simulations like those employed in this work.   In the context of using cluster counts to constrain cosmological parameters, the uncertainties in the selection function caused by projections are exacerbated by the high incompleteness of shear selected cluster samples.  Even for an efficiency cut of $e \\geq 60\\%$, the completeness is $\\gtrsim 50\\%$ only for the most massive clusters $M\\gtrsim 10.0^{14.3}$ which are near the peak of the lensing kernel at $z\\sim 0.3$, and lower everywhere else. The severe projection effects imply that these missed clusters are likely to scatter in or out of the detected sample. In light of the fact that this number of missed clusters is of the same order as the number detected, an accuracy of $x\\%$ in the final cluster sample requires a theoretical understanding the selection function to better than $\\sim x\\%$. This is likely to be very hard. One alternative is to throw away the number count information and use only the shape of the redshift distribution of shear-selected clusters to constrain cosmology since the shape will be much less sensitive to the precision to which the selection function is known.    Cross correlation of shear selected cluster samples with overlapping optical, SZ, or X-ray cluster catalogs is an important cosmological application (Schirmer \\etal 2003,2004; Hughes \\etal 2004).  For such a cross-correlation it is likely to be advantagous to work at lower efficiency thresholds, considering that the number of clusters per square degree for an efficiency $e\\geq 60\\%$ is five times larger than that with $e\\geq 75\\%$ (see Table \\ref{table:eff_tomo}).  These less significant peaks will span a larger mass and redshift range increasing the amount of overlap between the two cluster samples. Besides addressing probable baryonic biases for baryon selected clusters, cross correlation could be the most expedient route to precision cosmology with cluster surveys. In order to recover parameter sensitivity degraded by the scatter in the mass-observable relations, \\citet{MM03a} advocate calibrating these relations with extensive follow up mass measurements of a fraction of the clusters in a `baryon selected' sample.  An alternative to such a follow up campaign of individual clusters is to conduct a parallel weak lensing survey, and use the clusters detected from both weak lensing and, say SZ, to determine the mass-observable relation \\emph{statistically}. Because the mass and redshift distributions of the shear selected clusters is known for any cosmological model, the distribution of dual weak lensing-SZ detections can be used to simultaneously constrain cosmology and the relationship between mass and SZ decrement.  In addition to the statistical mass-SZ decrement calibration provided by the distribution of dual detections, stacked weak lensing mass measurements of the individual SZ clusters can be used to provide another constraint on this relation. Finally, it is worth emphasizing that an optical cluster search \\citep{Postman96,Kepner99,GY00,WK02,Koch03} could be conducted using the same deep imaging data used to measure weak lensing. Indeed, \\citet{Schirm03,Schirm04} have already combined shear selection with optical color selection, and confirmed several color-selected optical cluster candidates with weak lensing.  Such parallel analyses of shear-selected cluster samples and baryonic cluster samples will provide mass calibration, help break parameter degeneracies, and eliminate systematics inherent in cluster surveys.  The potential of shear selected clusters to provide precision constraints on the dark energy, either alone or in conjunction with another baryonic cluster survey, are topics which merit future investigation."
    },
    "0404/astro-ph0404380_arXiv.txt": {
        "abstract": "We identify 175 Seyfert galaxies from the Southern Sky Redshift Survey - SSRS2. We use the entire SSRS2 catalog to investigate the correlation between the presence of AGN with host environment. The AGN phenomenon is more strongly correlated with the internal host properties, than with the external environment. In particular, we find that Seyferts reside in more luminous galaxies, and are twice as frequent in barred galaxies and systems showing sign of advanced merger condition, when compared to a control sample. ",
        "introduction": "In order to study the dependence of AGN properties with local host environment and to derive reliable statistical results (see contribution of H. Schmitt in these proceedings), it is extremely important to select a suitable control sample. In this work we use the SSRS2 catalog \\cite{daC98}, which contains 5399 galaxies with $m_B \\le 15.5$ covering 1.69 sr of the southern sky. This catalog is currently 99.99\\% complete in redshift, and in our database we have optical spectra available for $\\sim$ 78\\% of the SSRS2 galaxies, while the remaining radial velocities were obtained from the literature. The spectra were inspected and AGNs identified using the diagnostic diagrams of line intensity ratios proposed by \\cite{Bal81}. New Seyfert galaxies were detected during the survey and were reported in the papers by Maia et al. (1987 and 1996). Additional AGN hosts in the SSRS2 were obtained inspecting the NASA Extragalactic Database, and the literature. For AGN host galaxies without spectra in the SSRS2 database, new observations were taken, therefore keeping the identification process homogeneous. Our AGN sample is made up of 175 Seyferts, which comprise ~3-4\\% of the parent sample. The ratio Sey-2/Sey-1 is 3:1. This fraction may contain an excess of Sey-2 which are in reality Sey-1 with broad lines that are partially obscured by the material in the disk of host galaxies seen almost edge-on. ",
        "conclusions": "We used the magnitude-limited SSRS2 to identify Seyfert galaxies, and investigated how the properties of their hosts relate to general population of non-active galaxies. We find that 175/5339 (3-4\\%) of the galaxies in the SSRS2 are Seyferts. The ratio of Seyfert 2s to 1s is $\\sim$3:1. Most Seyferts are of morphological types betwen S0a and Sb, and $\\sim$ 10\\% have hosts with indication of an ongoing merger. There is no difference between the morphology of Seyfert types 1 and 2. The AGNs are preferentially detected in high luminosity hosts, and are twice as frequent in barred hosts. No correlation with local density of galaxies. We find marginal evidence that Seyfert 1 galaxies have closer companions and/or are more susceptible to tidal effects than type 2. Additional details about the results presented here can be obtained in the paper by \\cite{Mai03}."
    },
    "0404/astro-ph0404039_arXiv.txt": {
        "abstract": "We report on preliminary results from our XMM-Newton observation of the Low Mass X-ray Binary 4U 1850-087, located in the galactic globular cluster NGC 6712. It is an ultracompact binary system, with an orbital period of 20.6 min. In previous low-resolution X-ray observations 4U 1850-087  displayed a soft excess  residual around 0.7 keV, possibly indicative of the presence of a high Ne/O ratio, as already found in other three ultracompact LMXBs containing a neutron star (4U 0614+091, 4U 0918-549, and 4U 1543-624). We discuss here our preliminary analysis of high resolution X-ray spectra of 4U 1850-087, and timing results. ",
        "introduction": "X1850-087 is an ultracompact Low Mass X-ray Binary (orbital period, P$_{orb}$, of 20.6 minutes, Homer et al. 1996) located in the globular cluster NGC6712 (at a distance of 6.8 kpc).  Ultracompact X-ray binaries (systems with orbital period less than 80 minutes) are believed to contain hydrogen--poor or degenerate companions (e.g. Verbunt \\& van den Heuvel, 1995).  Recent X-ray spectral evidence indicates that these degenerate companions contain anomalously high abundance of neutral Neon (e.g. Schulz et al., 2001; Juett et al., 2001). Indeed, a broad line-like feature at 0.7 keV, found in ASCA observations of ultracompact binaries (including X1850-087), is suggestive of the presence of an excess in the neutral neon absorption local to the systems (Juett et al., 2001). High resolution spectroscopy of the ultracompact LMXRBs 2S0918-549 and 4U 1543-624 confirmed an enhanced Ne/O number ratio with respect to the expected ISM ratio Ne/O=0.18 (Juett \\& Chakrabarty, 2003).    The aim of our XMM-Newton (Jansen et al. 2001) observation is to deeply investigate the presence of Neon excess suggested by previous ASCA observation, through  high spectral resolution observation with RGS (see den Herder eet al. 2001 for the description of the instrument, covering the energy range 0.3--2 keV). ",
        "conclusions": ""
    },
    "0404/astro-ph0404513_arXiv.txt": {
        "abstract": "Near-infrared, $JHKL$, photometry of $\\eta$~Car is reported covering the period 2000 to 2004. This includes the 2003 shell event which was the subject of an international multi-wavelength campaign. The fading that accompanied this event was similar to, although slightly deeper than, that which accompanied the previous one. The period between these events is $2023\\pm3$ days and they are strictly periodic. Their cause, as well as that of the quasi-periodic variations and secular brightening are discussed. It seems possible that all three types of variability are consequences of the binary nature of the star. ",
        "introduction": "$\\eta$ Carinae is one of the most luminous stars in the Galaxy, and despite being one of the best studied objects in the sky\\footnote{As of February 2004 NASA's ADS lists 844 publications mentioning $\\eta$ Car in the abstract, 250 of them published since 2000.0.}, it remains poorly understood. While most of its luminosity is emitted in the mid-infrared, there is flux from and variability at all wavelengths from hard X-rays to radio. $\\eta$ Car is in some senses the archetypal astronomical object - which can only be fully characterized through observations across the entire electromagnetic spectrum. Davidson \\& Humphreys (1997) provide an extensive review of $\\eta$ Car, though there have been many developments since its completion\\footnote{see, e.g., http://etacar.umn.edu.}.  It seems likely that we will soon be observing the details of similar objects in distant galaxies, so an understanding of this object potentially bears on many other topics.  Damineli (1996) discovered periodic variations in the equivalent width of HeI 1.083\\,$\\mu$m, and in the visibility of high excitation lines, which correlated with $JHK$ flux variations. He consequently suggested that $\\eta$ Car was in fact a binary system with an orbital period of about 5.5 years. Damineli predicted the next ``shell phase\" for 1997/98 and $\\eta$ Car was subjected to a period of intensified observations at all wavelengths (Morse et al. 1999). Not only did the ``shell phase\" occur on schedule, but it was accompanied by a number of unexpected phenomena, the most dramatic of which was a very sharply defined X-ray minimum (Corcoran et al. 2001). This was associated with, among other things, a very clear dip in the $JHKL$ light curves (Whitelock \\& Laney 1999), which we refer to as the ``eclipse-like event'' or simply the ``event'' in the following. Subsequent detailed examination of historical spectra (Feast et al. 2001) and monitoring at various wavelengths (e.g. 3\\,cm flux - Duncan \\& White 2003) established that a number of phenomena varied on this same timescale of 5.5 years.  The 5.5-year periodicity is now well established and most observers are convinced that $\\eta$ Car is a binary system. Some doubts do remain on the binary issue because, to date, none of the orbital parameter determinations have been convincing and there seem to be more questions about the underlying source(s) than ever before. The most recent ``shell phase\" occurred during 2003 and was observed even more intensively than the previous one. This paper is our contribution to the understanding of that event and we report broad-band near-infrared $JHKL$ photometry obtained over the last 5 years. It continues the series of papers which collectively cover 32 years of photometry of $\\eta$ Car (Whitelock et al. 1983 \\& 1994; Feast et al. 2001 - henceforth Papers I, II and III). We start with a detailed discussion in section 3 of the source of the near-infrared flux as it is somewhat different from that of the visible, mid-infrared or far-infrared radiation. It is crucial that we have some clarity on this before we go on to consider the variations in the $JHKL$ light curves and their various time scales in section 4.  In discussing the 5.5-year periodic phenomena, of which the 2003 shell event is the latest example, we refer our timings to the X-ray minimum observed by Corcoran et al. (2001) and Corcoran (2004). In particular we use Julian Day 2452819.7 (2003.49) as phase zero; this coincides with the start of the minimum (the faintest point) of the X-ray ``eclipse''. In doing this we make no particular comment on the correctness or otherwise of the Corcoran et al. model which has periastron at phase zero. The 2022 day period derived by Corcoran et al. is essentially identical to the 2023 days we derive from the $JHKL$ light curves (section 4.2).  In Paper III we interpreted the dip in the infrared light-curve, which repeats every 5.5 yrs, as an eclipse of a hot spot. We now believe that the balance of evidence does not favour what would normally be called an eclipse, as is discussed below in detail. Nevertheless, we continue to use terminology that is usually associated with eclipses (e.g. ``fourth contact'' in the caption to Fig.~5) simply because it provides a succinct description of what we observe. It is important that the reader is clear that the use of this terminology does not imply the existence of eclipses.  This may well be the last paper in this series as we anticipate the closure of the infrared photometer on the 0.75m telescope at Sutherland within the next year.  \\begin{table} \\begin{center} \\caption{Near-Infrared Photometry} \\begin{tabular}{rrrrr} \\hline \\multicolumn{1}{c}{JD}&\\multicolumn{1}{c}{$J$}& \\multicolumn{1}{c}{$H$}&\\multicolumn{1}{c}{$K$}&\\multicolumn{1}{c}{$L$}\\\\ \\multicolumn{1}{c}{--2450000}&\\multicolumn{4}{c}{(mag)}\\\\ \\hline 1709.00 &  2.490 &  1.599 &  0.454 & --1.708\\\\ 1712.22 &  2.487 &  1.597 &  0.450 & --1.696\\\\ 1714.25 &  2.474 &  1.598 &  0.449 & --1.702\\\\ 1737.23 &  2.499 &  1.615 &  0.466 & --1.685\\\\ 1738.21 &  2.507 &  1.619 &  0.463 & --1.694\\\\ 1743.21 &  2.514 &  1.607 &  0.458 & --1.681\\\\ 1808.66 &  2.543 &  1.643 &  0.479 &  \\\\ 1834.63 &  2.551 &  1.627 &  0.473 & --1.668\\\\ 1858.61 &  2.543 &  1.651 &  0.480 & --1.725\\\\ \\vspace*{2mm} 1862.61 &  2.547 &  1.643 &  0.483 & --1.692\\\\ 1869.62 &  2.508 &  1.619 &  0.470 & --1.723\\\\ 1881.61 &  2.533 &  1.620 &  0.453 & --1.705\\\\ 1890.62 &  2.518 &  1.614 &  0.448 & --1.730\\\\ 1926.56 &  2.572 &  1.672 &  0.498 & --1.740\\\\ 1929.54 &  2.574 &  1.683 &  0.504 & --1.717\\\\ 1962.47 &  2.571 &  1.664 &  0.495 & --1.687\\\\ 1965.45 &  2.559 &  1.670 &  0.489 & --1.734\\\\ 1976.45 &  2.580 &  1.666 &  0.481 & --1.710\\\\ 2032.34 &  2.487 &  1.586 &  0.448 & --1.758\\\\ \\vspace*{2mm} 2037.29 &  2.479 &  1.591 &  0.438 & --1.773\\\\ 2068.24 &  2.462 &  1.560 &  0.418 & --1.717\\\\ 2087.24 &  2.476 &  1.571 &  0.422 & --1.732\\\\ 2093.25 &  2.492 &  1.577 &  0.428 & --1.715\\\\ 2116.19 &  2.440 &  1.531 &  0.384 & --1.759\\\\ 2211.62 &  2.454 &  1.538 &  0.393 & --1.729\\\\ 2239.61 &  2.422 &  1.525 &  0.409 & --1.733\\\\ 2257.60 &  2.407 &  1.518 &  0.384 & --1.767\\\\ 2261.62 &  2.454 &  1.534 &  0.390 & --1.716\\\\ 2267.62 &  2.394 &  1.529 &  0.391 & --1.783\\\\ \\vspace*{2mm} 2283.58 &  2.426 &  1.512 &  0.371 & --1.750\\\\ 2284.61 &  2.464 &  1.551 &  0.405 & --1.710\\\\ 2285.58 &  2.456 &  1.551 &  0.410 & --1.727\\\\ 2287.57 &  2.456 &  1.547 &  0.411 & --1.751\\\\ 2288.59 &  2.489 &  1.586 &  0.466 & --1.693\\\\ 2289.64 &  2.466 &  1.562 &  0.426 & --1.725\\\\ 2318.52 &  2.468 &  1.568 &  0.419 & --1.738\\\\ 2320.49 &  2.430 &  1.538 &  0.399 & --1.752\\\\ 2324.49 &  2.462 &  1.575 &  0.427 & --1.760\\\\ 2346.43 &  2.474 &  1.555 &  0.415 & --1.747\\\\ \\vspace*{2mm} 2350.43 &  2.464 &  1.558 &  0.411 & --1.759\\\\ 2352.40 &  2.452 &  1.551 &  0.402 & --1.776\\\\ 2381.37 &  2.451 &  1.549 &  0.409 & --1.737\\\\ 2386.33 &  2.461 &  1.553 &  0.411 & --1.744\\\\ 2421.25 &  2.469 &  1.544 &  0.406 & --1.705\\\\ 2426.23 &  2.463 &  1.553 &  0.408 & --1.731\\\\ 2428.27 &  2.472 &  1.558 &  0.412 & --1.715\\\\ 2451.26 &  2.460 &  1.539 &  0.406 & --1.702\\\\ 2455.24 &  2.467 &  1.550 &  0.406 & --1.726\\\\ 2486.23 &  2.514 &  1.585 &  0.427 & --1.691\\\\ \\vspace*{2mm} 2490.19 &  2.512 &  1.582 &  0.420 & --1.695\\\\ 2492.20 &  2.513 &  1.587 &  0.421 & --1.692\\\\ 2525.68 &  2.487 &  1.561 &  0.397 & --1.711\\\\ 2529.67 &  2.512 &  1.580 &  0.413 &  \\\\ 2534.65 &  2.495 &  1.591 &  0.431 & --1.717\\\\ 2563.64 &  2.533 &  1.577 &  0.420 &  \\\\ 2564.63 &  2.488 &  1.585 &  0.419 & --1.725\\\\ 2568.64 &  2.478 &  1.557 &  0.399 & --1.747\\\\ 2573.63 &  2.473 &  1.566 &  0.412 & --1.744\\\\ 2593.62 &  2.490 &  1.563 &  0.413 & --1.752\\\\ 2603.60 &  2.479 &  1.562 &  0.411 & --1.737\\\\ \\hline \\end{tabular} \\end{center} \\end{table} \\setcounter{table}{0} \\begin{table} \\begin{center} \\begin{tabular}{rrrrr} \\hline \\multicolumn{1}{c}{JD}&\\multicolumn{1}{c}{$J$}& \\multicolumn{1}{c}{$H$}&\\multicolumn{1}{c}{$K$}&\\multicolumn{1}{c}{$L$}\\\\ \\multicolumn{1}{c}{--2450000}&\\multicolumn{4}{c}{(mag)}\\\\ \\hline 2661.57 &  2.490 &  1.554 &  0.382 & --1.761\\\\ 2664.57 &  2.485 &  1.560 &  0.387 & --1.783\\\\ 2665.59 &  2.481 &  1.550 &  0.394 & --1.770\\\\ 2667.56 &  2.480 &  1.554 &  0.390 & --1.781\\\\ 2675.56 &  2.483 &  1.557 &  0.393 & --1.761\\\\ 2676.51 &  2.483 &  1.562 &  0.386 & --1.783\\\\ 2680.52 &  2.512 &  1.560 &  0.377 & --1.747\\\\ 2689.49 &  2.481 &  1.561 &  0.404 & --1.760\\\\ 2690.44 &  2.497 &  1.548 &  0.399 & --1.749\\\\ \\vspace*{2mm} 2691.49 &  2.499 &  1.563 &  0.399 & --1.755\\\\ 2692.48 &  2.483 &  1.552 &  0.390 & --1.745\\\\ 2693.52 &  2.482 &  1.553 &  0.389 & --1.759\\\\ 2694.50 &  2.487 &  1.551 &  0.385 & --1.746\\\\ 2695.48 &  2.484 &  1.553 &  0.396 & --1.757\\\\ 2696.55 &  2.477 &  1.554 &  0.396 & --1.749\\\\ 2697.54 &  2.474 &  1.555 &  0.391 & --1.752\\\\ 2698.50 &  2.503 &  1.548 &  0.384 & --1.723\\\\ 2699.47 &  2.478 &  1.547 &  0.388 & --1.757\\\\ 2700.45 &  2.504 &  1.546 &  0.381 & --1.719\\\\ \\vspace*{2mm} 2701.52 &  2.481 &  1.539 &  0.389 & --1.751\\\\ 2702.46 &  2.489 &  1.542 &  0.384 & --1.727\\\\ 2710.44 &  2.476 &  1.556 &  0.383 & --1.747\\\\ 2711.49 &  2.477 &  1.540 &  0.378 & --1.756\\\\ 2712.54 &  2.462 &  1.529 &  0.368 & --1.757\\\\ 2715.53 &  2.488 &  1.529 &  0.327 & --1.784\\\\ 2716.43 &  2.455 &  1.511 &  0.359 & --1.785\\\\ 2752.34 &  2.426 &  1.454 &  0.298 & --1.806\\\\ 2753.33 &  2.428 &  1.458 &  0.300 & --1.793\\\\ 2754.31 &  2.402 &  1.452 &  0.294 & --1.806\\\\ \\vspace*{2mm} 2755.32 &  2.406 &  1.448 &  0.298 & --1.807\\\\ 2756.30 &  2.406 &  1.434 &  0.290 & --1.806\\\\ 2758.31 &  2.376 &  1.429 &  0.278 & --1.830\\\\ 2759.36 &  2.399 &  1.423 &  0.282 & --1.816\\\\ 2760.29 &  2.397 &  1.436 &  0.280 &  \\\\ 2761.29 &  2.384 &  1.424 &  0.282 & --1.806\\\\ 2762.29 &  2.396 &  1.428 &  0.283 & --1.814\\\\ 2764.42 &  2.382 &  1.430 &  0.273 & --1.809\\\\ 2773.39 &  2.355 &  1.400 &  0.270 & --1.825\\\\ 2774.26 &  2.324 &  1.392 &  0.256 & --1.842\\\\ \\vspace*{2mm} 2775.28 &  2.344 &  1.392 &  0.257 & --1.850\\\\ 2777.23 &  2.337 &  1.386 &  0.256 & --1.840\\\\ 2778.27 &  2.339 &  1.389 &  0.255 & --1.837\\\\ 2779.25 &  2.335 &  1.388 &  0.254 & --1.856\\\\ 2794.24 &  2.339 &  1.360 &  0.240 & --1.812\\\\ 2795.21 &  2.339 &  1.367 &  0.239 & --1.826\\\\ 2797.22 &  2.327 &  1.351 &  0.237 & --1.820\\\\ 2798.22 &  2.326 &  1.346 &  0.225 & --1.807\\\\ 2800.22 &  2.324 &  1.349 &  0.227 & --1.824\\\\ 2801.20 &  2.321 &  1.332 &  0.214 & --1.810\\\\ \\vspace*{2mm} 2802.23 &  2.305 &  1.317 &  0.209 & --1.817\\\\ 2803.23 &  2.311 &  1.338 &  0.222 & --1.827\\\\ 2804.21 &  2.299 &  1.315 &  0.217 & --1.833\\\\ 2805.21 &  2.293 &  1.319 &  0.211 & --1.831\\\\ 2806.20 &  2.305 &  1.318 &  0.211 & --1.821\\\\ 2807.20 &  2.292 &  1.307 &  0.208 & --1.829\\\\ 2809.28 &  2.291 &  1.310 &  0.207 & --1.843\\\\ 2810.27 &  2.291 &  1.301 &  0.200 & --1.815\\\\ 2811.26 &  2.282 &  1.301 &  0.203 & --1.836\\\\ 2813.30 &  2.283 &  1.308 &  0.221 & --1.800\\\\ 2814.29 &  2.281 &  1.307 &  0.215 & --1.813\\\\ \\hline \\end{tabular} \\end{center} \\end{table} \\setcounter{table}{0} \\begin{table} \\begin{center} \\caption[]{continued} \\begin{tabular}{rrrrr} \\hline \\multicolumn{1}{c}{JD}&\\multicolumn{1}{c}{$J$}& \\multicolumn{1}{c}{$H$}&\\multicolumn{1}{c}{$K$}&\\multicolumn{1}{c}{$L$}\\\\ \\multicolumn{1}{c}{--2450000}&\\multicolumn{4}{c}{(mag)}\\\\ \\hline 2815.22 &  2.287 &  1.312 &  0.217 & --1.785\\\\ 2816.33 &  2.267 &  1.318 &  0.238 & --1.785\\\\ 2817.25 &  2.289 &  1.318 &  0.239 & --1.773\\\\ 2818.19 &  2.301 &  1.324 &  0.244 & --1.756\\\\ 2819.27 &  2.300 &  1.328 &  0.249 & --1.748\\\\ 2820.21 &  2.288 &  1.325 &  0.261 & --1.728\\\\ 2821.26 &  2.307 &  1.340 &  0.274 & --1.726\\\\ 2823.23 &  2.303 &  1.352 &  0.297 & --1.713\\\\ 2825.19 &  2.332 &  1.384 &  0.324 & --1.695\\\\ \\vspace*{2mm} 2826.16 &  2.349 &  1.400 &  0.333 & --1.669\\\\ 2828.25 &  2.366 &  1.424 &  0.370 & --1.663\\\\ 2831.21 &  2.403 &  1.450 &  0.400 & --1.581\\\\ 2832.20 &  2.407 &  1.459 &  0.406 & --1.597\\\\ 2833.21 &  2.414 &  1.472 &  0.413 & --1.558\\\\ 2834.21 &  2.413 &  1.484 &  0.425 & --1.572\\\\ 2835.19 &  2.420 &  1.485 &  0.430 & --1.585\\\\ 2840.25 &  2.436 &  1.487 &  0.449 & --1.582\\\\ 2841.25 &  2.440 &  1.484 &  0.441 & --1.560\\\\ 2842.21 &  2.421 &  1.474 &  0.437 & --1.569\\\\ \\vspace*{2mm} 2843.20 &  2.429 &  1.473 &  0.436 & --1.586\\\\ 2844.20 &  2.423 &  1.469 &  0.437 & --1.560\\\\ 2845.20 &  2.408 &  1.451 &  0.418 & --1.562\\\\ 2846.19 &  2.406 &  1.459 &  0.433 & --1.561\\\\ 2847.19 &  2.412 &  1.449 &  0.416 & --1.561\\\\ 2848.19 &  2.399 &  1.438 &  0.407 & --1.570\\\\ 2849.19 &  2.394 &  1.436 &  0.395 & --1.559\\\\ 2850.19 &  2.363 &  1.431 &  0.398 & --1.585\\\\ 2852.19 &  2.346 &  1.413 &  0.379 & --1.625\\\\ 2854.19 &  2.354 &  1.411 &  0.384 & --1.595\\\\ \\vspace*{2mm} 2855.19 &  2.341 &  1.402 &  0.367 & --1.617\\\\ 2856.19 &  2.311 &  1.391 &  0.358 & --1.645\\\\ 2857.19 &  2.353 &  1.394 &  0.363 & --1.610\\\\ 2858.22 &  2.329 &  1.385 &  0.356 & --1.626\\\\ 2865.19 &  2.300 &  1.369 &  0.342 & --1.653\\\\ 2868.19 &  2.303 &  1.378 &  0.335 & --1.652\\\\ 2869.19 &  2.302 &  1.375 &  0.330 & --1.671\\\\ 2886.23 &  2.268 &  1.322 &  0.283 & --1.697\\\\ 2887.24 &  2.261 &  1.331 &  0.286 & --1.703\\\\ 2891.23 &  2.246 &  1.320 &  0.278 & --1.720\\\\ \\vspace*{2mm} 2893.23 &  2.256 &  1.300 &  0.270 & --1.729\\\\ 2894.22 &  2.261 &  1.326 &  0.269 & --1.704\\\\ 2896.23 &  2.232 &  1.333 &  0.274 & --1.767\\\\ 2918.64 &  2.240 &  1.319 &  0.272 & --1.721\\\\ 2931.64 &  2.238 &  1.322 &  0.260 & --1.696\\\\ 2932.63 &  2.252 &  1.329 &  0.265 & --1.691\\\\ 2933.64 &  2.244 &  1.336 &  0.274 & --1.731\\\\ 2955.62 &  2.265 &  1.351 &  0.298 & --1.727\\\\ 2959.62 &  2.255 &  1.358 &  0.301 & --1.744\\\\ 2961.61 &  2.271 &  1.362 &  0.294 & --1.718\\\\ \\vspace*{2mm} 2969.61 &  2.264 &  1.359 &  0.290 & --1.723\\\\ 2971.60 &  2.247 &  1.370 &  0.292 & --1.728\\\\ 2972.61 &  2.269 &  1.365 &  0.305 & --1.705\\\\ 2973.61 &  2.261 &  1.371 &  0.285 & --1.714\\\\ 2974.61 &  2.266 &  1.386 &  0.290 & --1.716\\\\ 2975.61 &  2.250 &  1.351 &  0.286 & --1.726\\\\ 2978.62 &  2.269 &  1.367 &  0.312 & --1.756\\\\ 3026.50 &  2.287 &  1.434 &  0.352 & --1.738\\\\ 3029.60 &  2.300 &  1.435 &  0.357 & --1.731\\\\ 3030.50 &  2.311 &  1.435 &  0.361 & --1.680\\\\ \\hline \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": "In view of the emphasis that Smith et al. (2000) put on the non-periodicity of the infrared flux variations (last paragraph of their section 4.2) it is important that we clarify this point. In section 4.2 we derived a period for the ``eclipse-like events\" of $2023\\pm3$ days; this is, as far as current information allows us to determine it, a ``precise period\". In contrast, the broad peaks seen in the $JHKL$ light curves over the last 30 years (see section 4.3) are quasi-periodic.  The similarity of the timescale of the quasi-periodic variations and the period between events suggests that the underlying cause of these phenomena is the same and that they are both orbitally modulated variations. Furthermore, the secular variations are characterized by discontinuous colour changes which occur around phase zero in the 2023 day cycle, which suggests that they too may be driven by something associated with the orbital period.  The morphology of the X-ray light curve during the 1997/8 event was explained by Corcoran et al.  (2001) in terms of a binary model with periastron passage at phase zero. The X-rays are produced by shock-waves in colliding stellar-winds and the specifics of the light curve can be reproduced if the mass-loss from $\\eta$ Car increases from $\\rm 3\\times 10^{-4}M_{\\odot}yr^{-1}$ by a factor of about 20 for a period of about 80 days following periastron. During the event the X-rays are absorbed in the dense wind resulting from the enhanced mass loss. Zanella et al. (1984) first suggested that the low-excitation spectroscopic phase and enhanced near-infrared emission might be caused by a shell ejection; although we note that they attributed the increased near-infrared flux to dust formation which is contrary to our results. When the 5.5\\,yr binary period started to look plausible various people conjectured that the close passage of the secondary in an elliptical orbit might trigger increased mass-loss at critical phases (e.g. Davidson 1999; Ishibashi et al. 1999).  Other evidence, cited in section 4.2, points to enhanced ultraviolet line opacity during the event, resulting in decreased excitation of emission lines at many wavelengths. Surprisingly there is no evidence for changes in the temperature or luminosity of the visual continuum from the central source and it is possible that the optical depth is sufficiently high that what is observed at high spectral resolution is well away from where the action is taking place. We have insufficient information to be certain what is causing the event observed at $JHKL$, but we might conjecture that free-free emission, coming from a rather broad volume around the central source, will greatly decrease if its source of ultraviolet excitation is quenched when the secondary is enveloped in the material pulled from the primary or from the passage of the secondary through dense pre-existing circumstellar material.  In this picture the quasi-periodic changes would be explained in a similar way to the radio variations (Duncan \\& White 2003) if the secondary star is largely responsible for ionizing the HII region surrounding $\\eta$~Car. Orbital modulation of the source of ionizing radiation as it moves in and out of the denser parts of the circumstellar disk will drive the quasi-periodic variations.  Much of the nebula must be optically thick at radio wavelengths but not in the near-infrared.  Thus the secular variations may also be the consequence of enhanced free-free emission as the newly ejected mass joins the material around $\\eta$ Car and emits at $JHK$ wavelengths. It may well be possible to use the information from the multi-wavelength campaign to establish where the material is building-up. We might speculate that if conditions become suitable for dust formation in these ejecta then radiation pressure could cause the dust and gas to be expelled with great efficiency."
    },
    "0404/astro-ph0404272.txt": {
        "abstract": "{Dust formation {in brown dwarf atmospheres} is studied by utilising a model for driven turbulence in the mesoscopic scale regime. We apply a pseudo-spectral method where waves are created and superimposed {within} a {limited} wavenumber interval. The turbulent kinetic energy distribution follows the Kolmogoroff spectrum which is assumed to be the most likely value. Such superimposed, stochastic waves may occur in a convectively active environment. They cause nucleation fronts and nucleation events and thereby initiate the dust formation process which { continues until} all condensible material is consumed. Small disturbances {are found to} have a large impact on the dust forming system. An initially dust-hostile region, which may originally be optically thin, becomes optically thick in a patchy way showing considerable variations in the dust properties during the formation process.  The dust appears in lanes and curls as a result of the interaction with waves, \\ie turbulence, which form larger and larger structures with time. Aiming on a physical understanding of the variability of brown dwarfs, related to structure formation in substellar atmospheres, we work out first necessary criteria for small-scale closure models to be applied in macroscopic simulations of dust forming astrophysical systems.} ",
        "introduction": "\\label{sec:Intro} %---------------------------------------------------------------------------- Substellar objects like brown dwarfs and (extrasolar) planets are largely - but not entirely - convective with considerable overshoot in the upper atmosphere. They also provide excellent conditions for the gas phase transition to solids or liquids (henceforth called {\\it dust}).  The understanding of such object's atmospheres requires therefore the modelling of convection and, because the inertia of the fluid is larger than its friction (Re$\\gg 10^4$), turbulent dust formation must be considered. Hence, substellar atmospheres involve various scale regimes with each being dominated by possibly different physical (\\eg streams, waves, precipitation) and chemical processes (\\eg combustion, dust formation, coagulation).   The large-scale structure of compact, substellar atmospheres is characterised by local and global convective motions (\\eg thunderstorms and monsoon like winds) and -- simultaneously -- by the gravitational settling of the dust. This scale regime has widely been investigated by 1D static, frequency dependent atmosphere calculations applying mixing length theory.  The presence of dust in form of several homogeneous constituents has been modelled by applying local stability and time scale arguments (Rossow 1978, Borrows\\etal 1997, Saeger\\plus Sasselow 2000, Ackermann\\plus Marley 2001, Allard\\etal 2001, Tsuji 2002, Cooper\\etal 2003)\\nocite{ros78,bmhlg97,sesa2000,am2001,ahats2001,tsu2002a,csmlb2002}. Recently, Woitke\\plus Helling (2003a; Paper~II)\\nocite{wh2003a} have proposed a consistent treatment of nucleation, growth, evaporation and gravitational settling of heterogeneous dust particles, which has been applied for the first time to stellar atmosphere models in (Woitke\\plus Helling 2004; Paper~III)\\nocite{wh2003b}.  Modelling macroscopic scales, Ludwig\\etal (2002)\\nocite{lah2002} have presented the first 3D simulations of M-dwarf convective atmospheres, applying the Large Eddy Simulation (LES) technique based on the experience of Nordlund \\& Stein with the solar convection. While the largest scales, which contain most of the energy and energise a small-scale turbulent fluid field, are computationally resolved, smaller scales are modelled by a hyperdiffusion which prevents the energy to accumulate in the smallest scales, thereby smoothing out all small-scale structures. Numerically, it stabilises the flow by filtering out sound and fast mode waves (see \\eg Caunt\\plus Korpi 2001)\\nocite{cako2001}.  Coming from the opposite site of the turbulent energy cascade, Helling\\etal(2001; Paper~I)\\nocite{holks2001b} have investigated the dust formation process in the small, microscopic scale regime ($l_{\\rm ref}\\ll H_{\\rho}$) by direct simulations of acoustic wave interactions.  These investigations of the dense, initially dust-hostile layers in brown dwarf atmospheres have revealed a {\\it feedback loop} which characterises the dust formation process: \\begin{center} \\parbox[h]{8cm}{The interaction of small-scale perturbations of the fluid field can cause a short-term temperature decrease low enough to initiate dust nucleation. The seed particles grow until they reach a size where the dust opacity is large enough to re-enforce radiative cooling which causes the temperature to decrease again below the nucleation threshold.  Dust nucleation is henceforth re-initiated which results in a further intensified radiative cooling. The nucleation rate and consequently also the amount of dust particles increase further.  This {\\sc run-away} process is stopped if either the radiative equilibrium temperature of the gas is reached or all condensible material is consumed.  Meanwhile, the seed particles have grown to macroscopic sizes ($\\mu$m).} \\end{center} Based on this knowledge, we extent our studies to larger and larger spatial scales aiming finally at the simulation of dust formation in the macroscopic scale regime, i.\\,e., the complete atmosphere. The next step is therefore to study the mesoscopic scale regime ($l_{\\rm ref} < H_{\\rho}$) where we model driven turbulence by stochastically superimposed waves in the inertial Kolmogoroff range and study the response of the dust complex. Necessary criteria are derived for a small-scale closure model to be applied in large scale simulations of dust forming systems.  The aim of such a scale-wise investigation is to understand the major physical mechanisms which are responsible for the structure formation in the atmospheres of substellar objects and to provide the necessary informations for building an appropriate sub-grid model needed to solve the closure problem inherent to any macroscopic turbulence simulation (not only) of dust forming media (see also Canuto 1997a, 2000)\\nocite{can97a,can2000}.  The challenge is that only the largest scales of the turbulence cascade are comparable with real astrophysical observations but structure formation is usually seeded on the smallest scales especially if it is correlated with chemical processes.  From the theoretical point of view, turbulence in thin atmospheric layers may be of quasi-2D-nature (Cho\\plus Polvani~1996, Menou\\etal 2003)\\nocite{mcsh2003,copo1996}. The 2D turbulence is characterised by an inverse energy cascade (transfer from small to large scales), contrary to 3D turbulence, which makes a scale-wise investigation even more urgent.  Various model approaches have been carried out to simulate and to study turbulence in different astrophysical scale regimes.  For example, thermonuclear flames in type Ia supernovae have been studied on small scales by R\\\"opke, Niemeyer \\& Hillebrandt~(2003)\\nocite{rnh2003}, investigating the Landau-Darrieus instability in 2D simulations which is responsible for the formation of a cellular structure of the burning front. Reinecke, Hillebrandt \\& Niemeyer~(2002)\\nocite{rehn2002} have performed large-scale calculations in order to model supernovae explosions on scales of the stellar radius. Mac Low~(1999; see also Mac Low\\plus Klessen 2004)\\nocite{ml1999,mk2004}, for instance, have set up isothermal initial velocity perturbations with an initial power spectrum of developed turbulence in the Fourier space and initially constant density. They model decaying turbulence on the small scales of the inertial subrange.  Smith, Mac Low \\& Heitsch~(2000)\\nocite{slh2000} use a similar approach but study the effect of driven turbulence in the same scale regime of star-forming clouds. A stationary but stochastic velocity field was applied by Wallin, Watson \\& Wylf~(1998)\\nocite{www98} in order to perform radiative transfer calculations of Maser spectra of a sub-parsec disk of a massive black whole. A fundamental theoretical investigation of the methods of driven turbulence is provided in (Eswaran \\& Pope 1987, 1988)\\nocite{ep88a,ep88b}. A different approach of turbulence and convective modelling is followed by Canuto (\\eg 1997b)\\nocite{can97b} who treats turbulent convection by a Reynold separation ansatz where decomposed quantities (background field + fluctuations) are introduced into the model equations.  In this paper, we present a model for driven turbulence which allows us to study the onset of dust formation under strongly fluctuating hydro- and thermodynamic conditions in the mesoscopic scale regime. We thereby intend to model a constantly occurring energy input from the convectively active zone.  Section~\\ref{sec:ModelProblem} states the model problem and the characteristic numbers of our astrophysical problem. The turbulence model is outlined in Sect.~\\ref{ssec:turbmodel}. In Sect.~\\ref{sec:Results}, the results are presented for 1D simulations and an illustrative 2D example. Section~\\ref{sec:disc} contains the discussion, and Section~\\ref{sec:concl} the conclusions.     %---------------------------------------------------------------------------- ",
        "conclusions": "\\label{sec:concl} %----------------------------------------------------------------------------  We have studied the onset of the dust formation process in a turbulent fluid field, typical for the dense and initially dust-hostile regions in substellar atmospheres. The main scenario is a convectively ascending fluid element in a brown dwarf atmosphere, which is excited by turbulent motions and just reaches sufficiently low temperatures for condensation, but other applications are also conceivable. The dust formation in a turbulent gas is found to be strongly influenced by the existence of a nucleation threshold temperature $T_S$. The local temperature $T$ must at least temporarily decrease below this threshold in order to provide the necessary supersaturation for nucleation.  Depending on the relation between the local mean temperature $\\overline{T}$ and $T_S$, three different regimes can be distinguished (see l.h.s. of Fig.~\\ref{fig:regimes}): (i) the {\\it deterministic} regime ($\\overline{T}<T_S$) where dust forms anyway, (ii) the {\\it stochastic} regime ($\\overline{T}>T_S$) where $T<T_S$ can only be achieved locally and temporarily by turbulent temperature fluctuations, and (iii) a regime where dust formation is {\\it impossible}. The size of the stochastic regime depends on the available turbulent energy. This picture of {\\it turbulent dust formation} is quite different from the usually applied {\\it thermodynamical picture} (r.h.s. of Fig.~\\ref{fig:regimes}) where dust is simply assumed to be present whenever $T<T_{\\rm sub}$, where $T_{\\rm sub}$ is the sublimation temperature of a considered dust material.  \\begin{figure}[t] \\begin{center} \\epsfig{file=PSFILES/regime_sketch.eps, scale=0.65} \\caption[]{Regimes of turbulent dust formation. $T_S$: nucleation threshold temperature (supersaturation $S\\ll 1$ required), $T_{\\rm sub}$: sublimation temperature.} \\label{fig:regimes} \\end{center} \\end{figure}  The investigations have been performed in the mesoscopic scale regime, where the test volume is excited by a spectrum of waves within a limited $k$-interval at given energy distribution, which are generated at the boundaries (pseudo-spectral method for driven turbulence). Two basic processes are found to be capable to initiate the dust formation even in a host-hostile environment: \\begin{itemize} \\item[ 1)] expansion waves with $\\overline{T}-\\Delta T<T_S$ ({\\it nucleation fronts}). \\item[ 2)] interactions of two or more expansion waves which cannot produce sufficiently low temperatures for themselves, but the superposition of such waves can $\\overline{T}-\\Delta T_1 - \\Delta T_2 -\\ldots<T_S$ ({\\it nucleation events}). \\end{itemize} After initiation, the dust condensation process is completed by a phase of active particle growth until the condensible elements are consumed, thereby preserving the dust particle number density for long times.  However, radiative cooling (as follow-up effect) is found to have an important influence on the subsequent dust formation, if the dust opacity reaches a certain critical value. This cooling leads to a decrease of $\\overline{T}(t)$ which may re-initiate the nucleation. This results in a runaway process (unstable feedback loop) until radiative and phase equilibrium is achieved. Depending on the difference between the initial mean temperature $\\overline{T}(t\\!=\\!0)$ and the radiative equilibrium temperature $T_{\\rm RE}$, a considerable local temperature decrease and density increase occurs. Since the turbulent initiation of the dust formation process is time-dependent and spatially inhomogeneous, considerable spatial variations of all physical quantities (hydro-, thermodynamics, dust) occur during the short time interval of active dust formation (typically a few seconds after initiation), which actually {\\it creates new turbulence}.  Thus, small turbulent perturbations have large effects in dust forming systems. A convectively ascending, initially dust free gas element, which is slightly warmer than its surroundings, can be excited to form dust by waves running through it, even at otherwise dust-hostile temperatures. The newly created dust particles may cause the substellar atmosphere to become almost instantaneously optically thick. Our 2D simulations show that the dust appears in lane-like and curled structures. Small-scale dust structures merge and form larger structures. Vortices appear to be present preferentially in regions without or with only little dust. Non of these structures would occur without turbulent excitation."
    }
}